{
    "0512/astro-ph0512043_arXiv.txt": {
        "abstract": "We revisit the problem of the star formation timescale and the ages of molecular clouds. The apparent overabundance of star-forming molecular clouds over clouds without active star formation has been thought to indicate that molecular clouds are ``short-lived'' and that star formation is ``rapid''. We show that this statistical argument lacks self-consistency and, even within the rapid star-formation scenario, implies cloud lifetimes $\\approx 10$ Myr. We discuss additional observational evidence from external galaxies that indicate lifetimes of molecular clouds and a timescale of star formation of $\\approx 10^7$ yr. These long cloud lifetimes in conjunction with the rapid ($\\approx 1$ Myr) decay of supersonic turbulence present severe difficulties for the  scenario of turbulence-controlled star formation. By contrast, we show that all 31 existing observations of objects for which the linewidth, the size, and the magnetic field strength have been reliably measured are in excellent {\\em quantitative} agreement with the predictions of the ambipolar-diffusion theory. Within the ambipolar-diffusion-controlled star formation theory the linewidths may be attributed to large-scale non-radial cloud oscillations (essentially standing large-amplitude, long-wavelength Alfv\\'{e}n waves), and the predicted relation between the linewidth, the size, and the magnetic field is a natural consequence of magnetic support of self-gravitating clouds. ",
        "introduction": "The ages of molecular clouds and the timescale of star formation are currently at the center of an important debate in the field. However, the debate has a long history. Early on, Giant Molecular Clouds (GMCs) were believed to be very long-lived ($> 10^8$ yr) \\citep{sc79}. That estimate relied on two arguments. First, the distribution of CO emission in galactocentric coordinates lacked a clearly recognizable spiral pattern, indicating that GMCs are situated in both arm and interarm regions \\citep{sss79}. This implied that GMCs must be older than the rotational period of the Galaxy ($\\simeq 10^8$ yr). Second, it was estimated that most of the interstellar hydrogen in the ``molecular ring'' (4-8 kpc) was molecular rather than atomic or ionic. Therefore the gas must spend most of its time in molecular form, which in turn implied GMC lifetimes greater than $10^8$ yr.  Those early arguments for very long lifetimes of GMCs were refuted by \\citet{bls80}. They showed that the molecular-to-atomic hydrogen gas ratio was overestimated, while the random motions of the GMCs were neglected in the estimates of the kinematic distance, thereby leading to an erroneous spatial mapping. Furthermore, they presented a number of arguments that set the upper limit on the ages of GMCs at a few $\\times 10^7$ yr.  It has recently been suggested that GMCs are short-lived ($\\simeq 10^6$ yr), transient objects \\citep{Elmegreen00,HBB01}. The idea of short-lived molecular clouds has been thought to favor a scenario of turbulence-controlled star formation over the ambipolar-diffusion--controlled theory for two reasons. First, if the lifetime of GMCs is smaller than the ambipolar-diffusion timescale, then ambipolar diffusion does not have enough time to operate in molecular clouds and thus cannot be relevant to the star-formation process. Second, short cloud lifetimes help to circumvent the problem of rapid dissipation of supersonic turbulence and ease the energy requirements on the source(s) of turbulence, whatever that(those) might be.  The idea of short-lived molecular clouds is based on observational estimates of the ages of newborn stars in star forming regions and on molecular-cloud core statistics. These observations, however, were shown to be in excellent quantitative agreement with predictions of the ambipolar-diffusion theory for the timescales of the {\\em observed} phases of star formation (Tassis \\& Mouschovias 2004; hereafter TM04). In this paper, we examine additional estimators of the ages of molecular clouds and of the star-formation timescale, and we present further observational evidence in favor of cloud lifetimes $\\simeq 10^7$ yrs that has received little or no attention until now. We discuss the implications that cloud ages $\\simeq 10^7$ yr have for current theories of star formation or for ideas on how the star-formation process is initiated. We also extend the work of Mouschovias \\& Psaltis (1995) and show that all 31 existing observations of objects (clouds, cores and even masers) for which the linewidth, the size, and the magnetic field strength have been reliably measured are in excellent quantitative agreement with the predicted relation between those three quantities, which is a natural and unavoidable consequence of magnetic support of self-gravitating clouds (Mouschovias 1987a). ",
        "conclusions": "In this paper, we have re-examined the argument that an observed overabundance of molecular clouds that are actively forming stars with respect to clouds without active star formation indicates that molecular clouds are ``short-lived'' and that star formation is ``rapid''. According to that argument, ambipolar diffusion does not have enough time to operate in molecular clouds. We have shown that: (a) Even if the observational facts used to support that argument were unbiased and accurate, these statistics imply molecular cloud lifetimes $\\simeq 10$ Myr, even within the ``rapid'' star-formation scenario. (b) Observations of molecular clouds in the solar neighborhood are {\\em not} in fact unbiased or representative --- rather, quiescent clouds are expected to be found mostly close to a galactic shock along spiral arms, as is observed in external, face-on spiral galaxies. (c) The ambipolar-diffusion theory of star formation does {\\em not} require long quiescent periods of molecular clouds or cloud lifetimes $\\gtrsim 10$ Myr, although it can certainly explain and accommodate such a possibility. Furthermore, if molecular clouds were short-lived, ``transient'', ``evanescent'', nongravitating structures \\citep{Elmegreen00,HBB01}, then they should be dispersing fast, in $\\simeq 1$ Myr. This implies that molecular clouds with embedded protostars should exhibit large, (ordered) expansion velocities. Such velocities are {\\em not} observed. \\footnote{Even if it {\\it were} the case that all molecular clouds had embedded stars, the conclusions one would draw within the ambipolar-diffusion theory would be: (1) The cloud ages are longer than the star-formation timescale. (2) The ambipolar-diffusion timescale in every molecular cloud is smaller than the cloud's lifetime. (3) Molecular clouds formed at about the same time (behind a spiral density shock wave). Consequently, statistical arguments cannot be applied to such clouds.}  We have also examined the implications of these results for star-formation theories or ideas. Theories or ideas that depend on supersonic, magnetized or non-magnetized turbulence for fragmentation, core formation, and cloud support are burdened by the requirement of continuously replenishing the turbulent motions during the entire lifetime of a molecular cloud. Proponents of this idea have found that internal driving of the turbulence cannot be reconciled with observations of actual molecular clouds. We have argued that external driving will most likely result in cloud compression and premature collapse. Also, the most promising external driving mechanisms (e.g., supernovae and winds from massive stars) assume the pre-existence of stars, and do not explain the origin of that previous generation of stars. Furthermore, recent observations \\citep{kshs05}, which find molecular clouds to be preferentially elongated along the Galactic plane, are in conflict with driving by supernovae and stellar winds, since these cannot account for the preferred elongation. Without an adequate source of driving, the observed supersonic linewidths cannot be maintained over a cloud's lifetime. Moreover, the linewidth-size relation first established by \\citet{l81} and extended by \\citet{lkm82} and \\citet{m83}, which was a key motivation for invoking turbulence in the first place, has not been explained by current numerical simulations of molecular cloud turbulence. Since superAlfv\\'{e}nic linewidths have never been observed in self-gravitating molecular clouds, turbulent star formation ideas that require superAlfv\\'{e}nic turbulence in order for their simulations to match observed cloud properties (e.g, Padoan \\& Nordlund 1999; Li \\etal\\, 2004) have no relevance to actual molecular clouds.  By contrast, the theory of ambipolar-diffusion--initiated star formation has, over the last thirty years, made numerous {\\em quantitative} predictions that turned out to be in excellent agreement with observations. Cloud envelopes are supported by magnetic fields while ambipolar diffusion allows supercritical cores to form and dynamically collapse in the deep interiors of self-gravitating clouds. This is supported by detailed numerical simulations as well as by polarimetry and Zeeman observations. The linewidths are due to nonradial cloud oscillations, which are essentially standing large-amplitude, long-wavelength Alfv\\'{e}n waves. Linewidth, size, and magnetic field data from 31 clouds, cores, and embedded masers are in excellent quantitative agreement with this theory. The theory does not suffer from a need to replenish turbulent motions in order to support clouds against collapse or to explain the linewidths, although if such replenishment takes place it has no effect on the ambipolar-diffusion theory of fragmentation and star formation: ambipolar diffusion damps the waves precisely over the lengthscales necessary for gravitational formation and contraction of fragments (or cores) (see Mouschovias 1987a, 1991).  We caution against using mass-to-flux ratios from observations of molecular cloud cores to make statements about the mass-to-flux ratios of molecular cloud envelopes. Supercritical cloud cores are a prediction of the ambipolar-diffusion--initiated star formation theory, and do not imply that magnetic support is insignificant in the envelopes. Moreover, the fact that observed cloud cores are critical or slightly supercritical and cloud envelopes are subcritical contradicts the simulations of turbulence-induced star formation, whose results imply that highly supercritical cores are equally likely as slightly supercritical cores.  It is important to clarify the semantic difference between ``rapid'' and ``slow'' star formation. When one refers to a process as being rapid or slow, one must also specify with respect to what. The main theoretical motivation for pursuing the idea of ``short-lived'' molecular clouds was to test whether ambipolar diffusion has enough time to operate and form supercritical cores over the lifetime of a molecular cloud. We have once again pointed out that the timescale of the ambipolar-diffusion--controlled core formation process depends on the mass-to-flux ratio and the degree of ionization of the parent cloud and it can be as short as 1 Myr for mildly subcritical clouds. Hence, from this perspective, clouds with ages of a few Myr are not ``young'' enough to render the ambipolar diffusion theory irrelevant to star formation. Loose terminology should not replace quantitative standards for determining the validity of a theory.  {\\bf Acknowledgments}: This work has been carried out without external support, and this paper would not have been published without the generosity of {\\em The Astrophysical Journal}."
    },
    "0512/astro-ph0512569_arXiv.txt": {
        "abstract": "We summarize the main features of several dark matter candidates in extra-dimensional theories. In particular, we review Kaluza-Klein (KK) gravitons in universal extra dimensions and branons in brane-world models. KK gravitons are superWIMP (superweakly-interacting massive particle) dark matter, and branons are WIMP (weakly-interacting massive particle) dark matter.  Both dark matter candidates are naturally produced in the correct amount to form much or all of dark matter. ",
        "introduction": "In models with universal extra dimensions (UED)~\\cite{Appelquist:2000nn}, gravity and all Standard Model (SM) fields have access to the entire space, which is called bulk space. The first studies of such models concentrated on the Kaluza-Klein (KK) partners of SM particles, but all UED models necessarily also have KK gravitons, and these may be viable DM candidates. Given the general formalism for analyzing the dynamics of gravitons in UED theories~\\cite{Feng:2003nr}, one can find the widths for decays of KK fermions and KK gauge bosons into KK gravitons.  These results are of special relevance when a KK graviton is the lightest KK particle and a superWIMP candidate~\\cite{SWIMP}, as they determine the observable implications of KK graviton DM for Big Bang Nucleosynthesis (BBN) analyses, the cosmic microwave background, the diffuse photon flux~\\cite{Feng:2003nr} and structure formation~\\cite{CFRT}.  The possibility of populating a large number of graviton states at different KK levels implies that the production of gravitons after reheating is extremely efficient and extremely sensitive to the reheat temperature $T_{RH}$. The constraints on $T_{RH}$ are therefore stringent (see Figure~\\ref{fig:reheat}).  \\begin{figure}% \\centering \\resizebox{8cm}{!}{\\includegraphics{reheat.eps}} \\caption{Overclosure and BBN constraints on the reheat temperature $T_{RH}$, as a function of the Kaluza Klein mass $m_{KK}$.  The vertical bands delimit regions of $B^1$ thermal relic abundance $\\Omega_{WIMP}$. See~\\cite{Feng:2003nr} for details. \\label{fig:reheat}} \\end{figure}  This superWIMP scenario can be studied in colliders. Decays to KK gravitons may be observed by trapping KK lepton next lightest KK particles in water tanks placed just outside collider detectors.  By draining these tanks periodically to underground reservoirs, slepton decays may be observed in quiet environments as in the gravitino case~\\cite{Buchmuller:2004rq,Feng:2004gn,Hamaguchi:2004df,Feng:2004yi,Brandenburg:2005he,Feng:2005gj}. Precision studies of KK lepton decays are therefore possible and can provide direct observations of gravitational effects at colliders; measurements of the extra-dimensional size and Newton's constant; precise determinations of the KK graviton's contribution to DM and laboratory studies of Big Bang nucleosynthesis and cosmic microwave background phenomena. ",
        "conclusions": "In this work we have reviewed the main features of KK gravitons in UED and branons in brane-worlds.  The thermal abundance of branons, and the non-thermal abundance of KK gravitons from late decays, are both naturally in the right range to form a significant component of dark matter.  We have considered the main phenomenological signals and constraints of these dark matter scenarios."
    },
    "0512/astro-ph0512619_arXiv.txt": {
        "abstract": "{We have carried out a spectroscopic survey of several {\\B}, {\\Be}, and shell stars in optical and near-infrared regions. Line profiles of the {\\Halpha} line and of selected {\\feii} and {\\oi} lines are presented. ",
        "introduction": "Observations of B, Be, and shell stars in different spectral regions are important for putting constraints on modeling these stars. Echelle spectrographs as well as the high sensitivity of modern detectors in the red part of the spectrum provide a wealth of the information contained in the whole visible and near infrared regions obtained simultaneously. The aim of this paper is a spectroscopic survey of line profiles of {\\Halpha} and selected non-hydrogenic lines of iron and oxygen in the visual and near infrared region for selected bright B and Be stars. We mainly use the echelle observations secured using the {\\HEROS} ({\\em H}eidelberg {\\em E}xtended {\\em R}ange {\\em O}ptical {\\em S}pectrograph) spectrograph attached to the Ond\\v{r}ejov 2m telescope supplemented by several CCD coud\\'e spectra. ",
        "conclusions": ""
    },
    "0512/astro-ph0512105_arXiv.txt": {
        "abstract": "We carry out a systematic search campaign for wide companions of exoplanet host stars to study their multiplicity and its influence on the long-term stability and the orbital parameters of the exoplanets. We have already found 6 wide companions, raising the number of confirmed binaries among the exoplanet host stars to 20 systems. We have also searched for wide companions of Gl\\,86, the first known exoplanet host star with a white dwarf companion. Our Sofi/NTT observations are sensitive to substellar companions with a minimum-mass of 35\\,$\\rm{M_{Jup}}$ and clearly rule out further stellar companions with projected separations between 40 and 670\\,AU. ",
        "introduction": "Some of the exoplanet host stars were found to be components of binary systems and first statistical differences between exoplanets around single stars and exoplanets located in binary systems were already reported by Zucker \\& Mazeh (2002) as well as Eggenberger et al. (2004). In particular, it seems that planets with orbital periods shorter than 40 days exhibit a difference in their mass-period and eccentricity-period distribution.  However, all the derived statistical differences are based only on a small number of known binary systems among the exoplanet host stars, i.e. their significance is sensitive to any changes in the sample size. Furthermore in the statistical analyses it is assumed that most of the exoplanet host stars are single star systems expect these stars known to be a component of a binary system.  In that context it is important to mention that the whole sample of exoplanet host stars was not systematically surveyed so far for neither wide nor close companions, i.e. several more exoplanet host stars, considered today as single stars, might be members of binary systems. Only search campaigns for companions of the exoplanet host stars will clarify their multiplicity status and will finally verify the significance of the reported statistical differences.  Therefore, we have started an imaging search program for wide visual companions of exoplanet host stars, carried out with UFTI/UKIRT, SofI/NTT as well as MAGIC/CA~2.2m. We can find all directly detectable stellar and substellar companions (m$\\ga$40\\,$\\rm{M_{Jup}}$) with projected separations from about 50 up to 1000\\,AU. Thereby companions are identified first with astrometry (common proper motion) and their companionship is confirmed with photometry and spectroscopy later on. So far, 6 wide companions were detected, see Mugrauer et al. (2005a) for further details. ",
        "conclusions": ""
    },
    "0512/astro-ph0512349_arXiv.txt": {
        "abstract": "We determine the distance to \\nclu\\/ clusters of galaxies in the redshift range 0.14$\\leq$z$\\leq$0.89 using X-ray data from \\chandra\\/ and Sunyaev-Zeldovich Effect data from the {\\it Owens Valley Radio Observatory} and the {\\it Berkeley-Illinois-Maryland Association} interferometric arrays.  The cluster plasma and dark matter distributions are analyzed using a hydrostatic equilibrium model that accounts for radial variations in density, temperature and abundance, and the statistical and systematic errors of this method are quantified.  The analysis is performed via a Markov chain Monte Carlo technique that provides simultaneous estimation of all model parameters.  We measure a Hubble constant of \\HO (statistical followed by systematic uncertainty at 68\\% confidence) for an $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}$=0.7 cosmology.  We also analyze the data using an isothermal $\\beta$ model that does not invoke the hydrostatic equilibrium assumption, and find \\HOall; to avoid effects from cool cores in clusters, we repeated this analysis excluding the central 100 kpc from the X-ray data, and find \\HOcut\\/ (statistical followed by systematic uncertainty at 68\\% confidence).  The consistency between the models illustrates the relative insensitivity of SZE/X-ray determinations of $H_0$ to the details of the cluster model.  Our determination of the Hubble parameter in the distant universe agrees with the recent measurement from the {\\it Hubble Space Telescope} key project that probes the nearby universe. ",
        "introduction": "\\label{sec:intro}  Combined analysis of radio and X-ray data provides a method to determine directly the distances to galaxy clusters.  Galaxy clusters are the largest gravitationally collapsed structures in the universe, with a hot diffuse plasma ($T_e \\sim 10^7 - 10^8$ K) that fills the intergalactic space.  Cosmic microwave background (CMB) photons passing through this hot intracluster medium (ICM) have a $\\sim 1$\\% chance of inverse Compton scattering off the energetic electrons, causing a small ($\\sim 1$ mK) distortion of the CMB spectrum, known as the Sunyaev-Zel'dovich Effect (SZE: Sunyaev \\& Zel'dovich 1970, 1972; for reviews see Birkinshaw 1999; Carlstrom, Holder, \\& Reese 2002). The same hot gas emits X-rays primarily through thermal bremsstrahlung.  The SZE is a function of the integrated pressure, $\\Delta T \\propto \\int n_e T_e d\\ell$, where $n_e$ and $T_e$ are the electron number density and temperature of the hot gas, and the integration is along the line-of-sight.  The X-ray emission scales as $S_X \\propto \\int n_e^2 \\Lambda_{ee} d\\ell$, where $\\Lambda_{ee}$ is the X-ray cooling function.  The different dependences on density, along with a model of the cluster gas, enable a direct distance determination to the galaxy cluster.  This method is independent of the extragalactic distance ladder and provides distances to high redshift galaxy clusters.  The $\\sim 1$ mK SZE signal proved challenging for initial searches, but recent improvements in both technology and observational strategies have made observations of the SZE fairly routine.  High signal-to-noise ratio (S/N) detections of the SZE have been made with single dish observations at radio wavelengths (Birkinshaw and Hughes 1994; Herbig et al.\\ 1995; Myers et al.\\ 1997; Hughes and Birkinshaw 1998; Mason et al.\\ 2001), millimeter wavelengths (Holzapfel et al.\\ 1997a,b; Pointecouteau et al.\\ 1999, 2001) and submillimeter wavelengths (Lamarre et al.\\ 1998; Komatsu et al.\\ 1999), while interferometric observations at centimeter wavelengths have produced images of the SZE (Jones et al.\\ 1993; Grainge et al.\\ 1993; Carlstrom et al.\\ 1996, 2000; Grainge et al.\\ 2002; Reese et al.\\ 2000, 2002; Grego et al.\\ 2000, 2001; La~Roque et al.\\ 2003; Udomprasert et al.\\ 2004).  SZE/X-ray distances provide a measure of the Hubble constant that is independent of the extragalactic distance ladder and probe high redshifts, well into the Hubble flow.  The SZE/X-ray determinations of $H_0$ bridge the gap between observations of nearby objects (e.g. the Hubble Space Telescope Key Project, Freedman et al.\\ 2001) and expansion values inferred from CMB anisotropy (Spergel et al.\\ 2003)  and supernova (Riess et al.\\ 2005) measurements. Previous SZE/X-ray determinations of the Hubble parameter have progressed from analysis of individual galaxy clusters, to samples of a few (Myers et al.\\ 1997; Mason et al.\\ 2001; Jones et al.\\ 2005), up to a sample of 18 galaxy clusters using ROSAT X-ray data (Reese et al 2002; for reviews see Reese 2004 and Carlstrom, Holder, \\& Reese 2002).  In most cases, simple isothermal $\\beta$ models were adopted for the cluster gas, since the data did not warrant a more sophisticated treatment.  We present a Markov chain Monte Carlo joint analysis of interferometric SZE observations and \\chandra\\ X-ray imaging spectroscopy observations of a sample of \\nclu\\ galaxy clusters with redshifts $0.14 \\leq z \\leq 0.89$.  The unprecedented spatial resolution of \\chandra\\ combined with its simultaneous spectral resolution allow more realistic modeling of the intracluster plasma than previous studies, thus enabling a more accurate determination of the Hubble constant. ",
        "conclusions": "We analyzed \\nclu\\/ clusters of galaxies with \\chandra\\ X-ray imaging spectroscopy and \\ovro-\\bima\\/ SZE data, the largest sample to date used to measure $H_0$. We applied a hydrostatic equilibrium model that accounts for radial variations in cluster temperature, and for sharp density gradients caused by the cooling of the plasma in the cluster core.  The joint analysis of X-ray and SZE data yields a direct measurement of the cosmic distance scale in the redshift range 0.14$<$z$<$0.89.  We measure a Hubble constant of \\HO\\ for an $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ cosmology (68 \\% confidence interval, statistical followed by systematic uncertainty), which is in agreement with the {\\it Hubble Space Telescope} Key Project results obtained at low redshift.  We also analyze our measurements with a simple isothermal model of the hot plasma without the hydrostatic equilibrium assumption. The results from this simple model are in good agreement with the hydrostatic equilibrium model, indicating that the X-ray/SZE method used to determine the cosmic distance scale is largely insensitive to the details of the hot plasma modeling.  \\vspace{2cm} Acknowledgments:  This paper is dedicated to Leon van Speybroeck and his colleagues on the \\chandra\\/ project, including H.\\ Ebeling, W.\\ Forman, J.P.\\ Hughes, C.\\ Jones, M.\\ Markevitch, H.\\ Tananbaum, A.\\ Vikhlinin and M.\\ Weisskopf; without their effort to construct an exceptional observatory and to obtain deep observations of a large number of clusters, this research would not have been possible.  The support of the BIMA and OVRO staff over many years is also gratefully acknowledged, including J.R.\\ Forster, C.\\ Giovanine, R.\\ Lawrence, S.\\ Padin, R.\\ Plambeck, S.\\ Scott and D.\\ Woody.  We thank C.\\ Alexander, K.\\ Coble, A.\\ Cooray, L.\\ Grego, G.\\ Holder, W.\\ Holzapfel, A.\\ Miller, J.\\ Mohr, D.\\ Nagai, S.\\ Patel and P.\\ Whitehouse for their outstanding contributions to the SZE instrumentation, observations, and analysis. We also thank J.\\ Mohr and D.\\ Nagai for contributions to the development of the hydrostatic equilibrium model.  This work was supported by NASA LTSA grant NAG5-7985 and also in part by NSF grants PHY-0114422 and AST-0096913, the David and Lucile Packard Foundation, the McDonnell Foundation, and a MSFC director's discretionary award.  Research at the Owens Valley Radio Observatory and the Berkeley-Illinois-Maryland Array was supported by National Science Foundation grants AST 99-81546 and 02-28963.  Calculations were performed at the Space Plasma Interactive Data Analysis and Simulation Laboratory at the Center for Space Plasma and Aeronomy Research of the University of Alabama at Huntsville."
    },
    "0512/astro-ph0512380_arXiv.txt": {
        "abstract": "The intergalactic neutral hydrogen which is responsible for the \\lya\\ forest of quasar absorption lines is a tracer of much larger amounts of ionised hydrogen. The ionised component has yet to be detected directly, but is expected to scatter CMB photons via the Sunyaev-Zel'dovich effect. We use hydrodynamic simulations of a LambdaCDM universe to create mock quasar spectra and CMB sky maps from the same volume of space.  We find that the high redshift \\lya\\ forest gas causes temperature fluctuations of the order of 1 ${\\rm \\mu}$K rms in the CMB on arcmin scales.  The kinetic and thermal Sunyaev-Zel'dovich effects have a similar magnitude at redshift three, with the thermal effect becoming relatively weaker as expected at higher redshift. The CMB signal associated with lines of sight having HI column densities $ >10^{18} cm^{-2}$ is only marginally stronger than that for lower column density sightlines. There is a much more significant dependence of rms temperature fluctuation on mean \\lya\\ absorbed flux, however, suggesting that the CMB signal effectively arises in lower density material.  We investigate the extent to which it is possible to cross-correlate information from the \\lya\\ forest and the microwave background to detect the Sunyaev-Zel'dovich effect at redshifts $2-4$.  In so doing, we are enable to set direct limits on the density of diffuse ionised intergalactic baryons.  We carry out a preliminary comparison at a mean redshift $z=3$ of 3488 quasar spectra from SDSS Data Release 3 and the WMAP first year data. Assuming that the baryons are clustered as in a LambdaCDM cosmology, and have the same mean temperature, the cross-correlation yields a weak limit on the cosmic density of ionised baryons $\\Omega_{\\rm b,I}$.  As a fraction of the critical density, we find $\\Omega_{\\rm b,I} < 0.8$ at $95\\%$ confidence. With data from upcoming CMB telescopes, we anticipate that a direct detection of the high redshift ionised IGM will soon be possible, providing an important consistency check on cosmological models. ",
        "introduction": "The \\lya\\ forest seen in quasar spectra is now thought to be produced by neutral hydrogen atoms in an intergalactic medium (IGM) which is in photoionisation equilibrium (see e.g., Rauch 1998 for a review).  At low to moderate redshifts, the fraction of neutral atoms is small (1 in $\\sim 10^{6}$; e.g. Gunn \\& Peterson 1965) but knowledge of the photoionisation rate (from summing the contribution of known sources; e.g. Haardt \\& Madau 1996) makes it possible to extrapolate from this one part in a million measurement to an estimate of the total baryon density $\\Omega_{b}$ (Weinberg \\etal 1997, Rauch \\etal 1997). Loeb (1996) has argued that the ionised component of the IGM should be directly observable by inverse Compton or Doppler scattering of cosmic microwave background (CMB) photons by free electrons (Sunyaev-Zel'dovich [1972,1980]; hereafter, the SZ effect).  Such a detection would provide consistency checks on the currently favoured cosmological model for structure formation, and place constraints on the IGM density, temperature, ionisation state and velocity field.  In this paper, we use a hydrodynamic cosmological simulation to investigate how observations of the CMB temperature and \\lya\\ spectra of quasars drawn from the same region of sky can be used in the search for the hitherto unseen $99.9999\\%$ of the hydrogen associated with the \\lya\\ forest gas at redshifts $z \\approx 2-5$.  In favoured models for structure formation, the optical depth for \\lya\\ forest absorption arises in a continuously fluctuating medium (e.g., Cen \\etal 1994; Zhang, Anninos, \\& Norman 1995; Petitjean, M\\\"ucket, \\& Kates 1995; Hernquist \\etal 1996; Katz \\etal 1996a; Wadsley \\& Bond 1997; Theuns \\etal 1998; Dav\\'e et al. 1999). The density of neutral hydrogen is inversely proportional to the photoionisation rate $\\Gamma$: \\begin{equation} \\Gamma=\\int^{\\infty}_{\\nu_{HI}}d\\nu\\frac{4\\pi J(\\nu)}{h\\nu}\\sigma_{HI}(\\nu) \\, , \\end{equation} and the optical depth to absorption is given by \\begin{equation} \\tau_{Ly\\alpha}=\\frac{\\pi e^{2}}{m_{e}c} f \\lambda H^{-1}(z)n, \\end{equation} where $f=0.416$ is the \\lya\\ oscillator strength, $n$ is the neutral hydrogen number density, and $\\lambda=1216$ \\AA\\ (Gunn and Peterson 1965).  In order for $\\tau$ be of order unity, (as seen, for example at redshifts $z\\sim3$), $n$ must be much smaller than the total density of hydrogen atoms, as noted earlier.  The gas in photoionisation equilibrium lies at densities $\\sim 1-10$ times the mean, and obeys a power law relationship between density and temperature (Gnedin \\& Hui 1998). At $z=3$, according to hydrodynamic simulations (Dav\\'{e} \\etal 2001), more than $90 \\%$ of the baryons are expected to be in this diffuse component. The tiny neutral fraction is readily detectable, but not the dominant ionised portion of the material.  Owing to gravitational evolution, at $z<2$ most of the baryons are expected to reside in a shock heated IGM, with little neutral hydrogen (Cen \\& Ostriker 1999, Dav\\'e \\etal 2001).  The interaction of CMB photons with free electrons in the intergalactic plasma was first considered by Sunyaev and Zel'dovich (1972). Inverse Compton scattering preferentially increases the energy of CMB photons, while conserving photon number, leading to a spectral distortion whose amplitude is proportional to the product of electron temperature and density (the thermal SZ effect). Doppler scattering induces an intensity fluctuation with the same spectral shape as the CMB itself (the kinetic SZ effect). The thermal effect is one of the main sources of secondary CMB anisotropies on small angular scales.  The perturbation in the CMB thermodynamic temperature resulting from scattering of nonrelativistic electrons is \\begin{equation} \\frac{\\Delta T}{T}=y(x\\frac{e^{x}+1}{e^{x}-1}-4) \\end{equation} \\begin{equation} \\simeq -2y \\qquad {\\rm for} \\qquad x << 1, \\end{equation} where $x=h\\nu/kT_{\\rm CMB}\\simeq\\nu/56.85$ GHz is the dimensionless frequency and the second expression is valid in the Rayleigh-Jeans limit, which we assume henceforth. The quantity $y$ is known as the comptonization parameter and is given by \\begin{equation} y\\equiv \\int dl \\frac{n_{e}k(T_{e}-T_{\\rm CMB})}{m_{e}c^{2}}, \\end{equation} where the integral is performed along the photon path.  The kinetic SZ effect arises from the motion of ionised gas with respect to the rest frame of the CMB. The resulting temperature fluctuation is $\\Delta T/T=-b$ where \\begin{equation} b\\equiv\\sigma_{T}\\int dl n_{e} \\frac{v_{r}}{c}, \\end{equation} gives the magnitude of the effect and $v_{r}$ is the component of the gas peculiar velocity along the line of sight (positive sign for receding gas, negative for approaching) to the observer.  The gas which causes \\lya\\ forest absorption has a density near the cosmic mean, with a volume weighted temperature around $2\\times 10^4$ K at $z=3$ (e.g, Schaye \\etal 2000, McDonald \\etal 2001). At this temperature, we expect the contribution from the kSZ to dominate the $rms$ temperature fluctuations of the CMB from the \\lya\\ forest. Using observed \\lya\\ forest line counts, Loeb (1996) estimated that the $rms$ temperature fluctuations from the kSZ would be $\\Delta T/T \\sim 10^{-6}$ for the \\lya\\ forest lying between $z=2$ and $z=5$. This is for angular scales of order 1 arcmin, with a $\\theta^{-1}$ decrease on larger scales. There is much competition from other signals on these small scales (including low redshift thermal SZ from galaxy clusters, the Ostriker-Vishniac (1986) effect, patchy reionisation (see, e.g. McQuinn et al. 2005), and foreground sources). The suggestion to cross-correlate information from SDSS \\lya\\ forest spectra and the CMB was also made by Loeb (1996) in order to try to better extract the signal.  In this paper, we measure the kSZ effect in simulations, as well as the thermal effect. The mass weighted temperature of the IGM at $z\\sim3$ is expected to be closer to $10^6K$ (e.g., Springel, White and Hernquist 2001), and so the thermal effect should be significant. How much of that hot gas is physically associated with \\lya\\ absorption will affect how well the tSZ can be detected by cross-correlation.  We note that both being absorption phenomena, neither the SZ signal nor the \\lya\\ forest opacity are directly affected by the inverse square law. The problem of finding bright QSO background sources aside, they therefore both have an advantage in the hunt for baryons at high redshifts. Indeed, the value of the SZ effect for finding galaxy clusters at the highest redshift has long been recognised (see e.g., Carlstrom \\etal 2002 for a review).  We aim to explore its potential for finding diffuse ionised material at $z>2$.  The plan of this paper is as follows. In Section 2, we give details of the hydrodynamic cosmological simulation and how mock \\lya\\ spectra and CMB temperature sky maps were constructed from the same volume of space. In Section 3, we compute statistical measures of the absorption and examine correlations between the two probes of the intergalactic medium. In Section 4, we compare directly to observational data from the SDSS and WMAP, and in Section 5 we summarise and discuss our results. ",
        "conclusions": "\\subsection{Summary}  Using a cosmological hydrodynamic simulation we have investigated the CMB temperature fluctuations caused by the diffuse IGM at redshifts $z=2-6$, also responsible for the \\lya\\ forest.  Our main findings can be summarised as follows:\\\\ \\\\ \\noindent(1) At redshift $z=2$, the kinetic Sunyaev-Zel'dovich effect from gas within the \\lya\\ to \\lyb\\ region of the forest is predicted to cause CMB fluctuations of amplitude 0.8 $\\mu$K rms averaged in a top-hat filter of angular radius 1 arcmin. This remains roughly constant for higher redshifts.\\\\ \\\\ \\noindent (2) The thermal SZ at redshift $z=2$ yields an rms T of 2.4 $\\mu$K. At higher redshift, the relative fluctuation owing to the thermal SZ compared to kinetic SZ declines as expected because of the decreasing mean temperature. However, there is still a substantial $T_{\\rm rms}$ at higher redshift, 0.4 $\\mu$K at $z=4$.\\\\ \\\\ \\noindent(3) There is a correlation between \\lya\\ forest properties of quasar spectra with CMB temperature fluctuations caused by both the kinetic and thermal SZ effects. In particular, there is a strong relation between the mean effective optical depth \\taueff\\ averaged over the entire \\lya\\ forest region and the CMB temperature.  Pixels (of size $\\sim$ 1 arcmin) in the top 10 percentile of \\taueff\\ are predicted to have a mean CMB temperature 1.5 $\\mu$K lower than the bottom 10 percentile at $z=2$, owing to the thermal SZ effect.  The quantity \\taueff\\ is relatively sensitive to density fluctuations in the diffuse IGM.\\\\ \\\\ \\noindent(4) There is, however, little correlation between the total HI column density in spectra and the CMB temperature, with no relation at all seen for column densities $N_{HI}>10^{16} cm ^{-2}$. This statistic is most sensitive to the neutral hydrogen density in saturated regions, indicating that SZ effects correlate best with relatively diffuse unsaturated gas.\\\\ \\\\ \\noindent(5) We have compared the correlation of \\lya\\ \\taueff\\ with CMB T seen in simulations of the $\\Lambda$CDM cosmology with observational data from the Sloan Digital Sky Survey and WMAP. Because the kinetic SZ signal and pixel noise add in quadrature, no constraints are possible on the IGM velocity field with the present data. However, the sensitivity is within an order of magnitude of that necessary to detect the predicted thermal SZ effect from the \\lya\\ forest. We are able to place a weak limit on the density in ionised baryons at $z=3$ of $\\Omega_{B,i} < 0.80$ at $95 \\%$ confidence. Below we outline future developments which would make a direct detection possible.  \\subsection{Discussion}  \\label{disc}  In the present paper, we have examined the relationship between \\lya\\ forest absorption and the SZ effect in general, without the expectation that a detection of a correlation can be made from current data. In the future, the quality and size of datasets will increase, and it will become worthwhile to plan a strategy for the cross-correlation that makes the best use of the data. Two obvious improvements could be made in the future. In our illustrative comparison of WMAP and SDSS data we did not use the full angular or frequency information, both of which would help with the sensitivity of the test.   As we have seen that the thermal SZ effect offers the best hope for detection of the ionised baryons, making use of the spectral difference between the SZ decrement and the primary CMB signal will enable easier detection of the effect.  One could imagine fitting the \\taueff\\ -T correlation in different CMB frequency bands to do this. For example, in the work of Huffenberger \\etal (2004), a method was derived for combining the WMAP channels in an optimal manner to remove the primary anisotropy, effects of noise and contaminating point sources, enabling the thermal SZ signature to be constrained directly. In this paper we have concerned ourselves with predictions for the Rayleigh Jeans part of the spectrum, which is reasonable given the large observational uncertainties in present data. We note that at higher frequencies the thermal SZ signal becomes positive (for example at 330 GHz the magnitude is the same as the Rayleigh Jeans signal but with the opposite sign). Bolometers tend to be more sensitive than radiometers at present, and they operate at frequencies $>100$ GHz. These differences could be used profitably to increase the signal to noise of any detection.  Also, with the WMAP data, we have limited our analysis to single pixels.  Making use of the full spatial cross-correlation in our detection of the signal will be useful when CMB datasets with higher angular resolution become available. Another possibility could be to use wavelets to isolate the relevant physical scales in a future analysis (see e.g., Vielva \\etal 2005).  One can ask whether emission from the QSOs themselves might create a signal, as the SED of QSOs in the microwave band is not well constrained (Perna \\& Di Matteo 2000, White \\& Majumdar 2004).  Apart from the possible intrinsic microwave emission of QSOs, there will also be an intrinsic thermal SZ effect from galaxy clusters which host QSOs.  Neither of these will necessarily bias any measurement of the \\lya\\ -CMB correlation unless QSOs with more \\lya\\ absorption in their spectra have more intrinsic microwave emission.  This seems unlikely in the former case, although in the case of an SZ signal associated with clusters in which the QSOs lie one could argue that this is part of the signal that we are trying to measure.  Otherwise, it could be removed if we so desire by not using the part of the spectrum closest to the QSO.  The method we have outlined in this paper is one way to potentially detect ionised baryons in the IGM at high redshift. In general, one would like to know if there are any other methods that can be used. At lower redshifts, for example ($z<1$), the ionised baryons are mainly collisionally ionised, with such a low neutral fraction that they cannot be seen in the \\lya\\ forest.  Their higher temperature, however, means that this Warm Hot Intergalactic Medium (WHIM, see e.g., Cen \\etal 1995, 1999, Dav\\'e et al 2001) can potentially be detected by looking for X-ray absorption (e.g., Fang \\etal 2002) or emission (Yoshikawa \\etal 2004, Fang \\etal 2005) or perhaps in \\lya\\ emission (e.g. Furlanetto et al. 2003, 2005).  The thermal SZ signal at low redshifts is also one way to detect the WHIM. At higher redshifts and lower IGM temperatures considered in this paper however, the X-ray signal (see e.g., Croft \\etal 2001) is not likely to be observable.  The overall $y$-distortion in the CMB spectrum from the hot IGM should eventually be detectable. Limits from the COBE FIRAS instrument (e.g., Mather \\etal 1990) are $y < 2.5 \\times 10^{-5}$ at $95 \\%$ confidence, more than an order of magnitude above the mean $y$ caused by the high-$z$ IGM.  However, separating out the contribution from the high $z$ IGM will be difficult without a spatial correlation of the type we have proposed here. One could also use other tracers of the high $z$ density field for a cross-correlation, such as Lyman-break galaxies, or QSOs themselves, although their connection to the IGM would be less easy to interpret.  Another interesting question to ask is how likely it is that we will be able to constrain the high $z$ IGM velocity field, given that the kinetic SZ signal from the photoionised IGM is expected to be relatively important (Loeb 1996). Because of the possible +ve or -ve nature of the effect, depending on the velocity, we cannot use a mean signal, unlike the thermal SZ. By using an $rms$ fluctuation as our statistic, we therefore need a noise level which is extremely low. Since the spectrum of the kSZ is the same as the CMB, the primary fluctuations become even more important sources of noise than with the thermal effect.  Given that the rms fluctuations on scales of 1 arcmin from the kSZ are predicted to be $\\sim 5\\mu$K at $z=2$, one would need the ``noise'' from primary anisotropies and other sources to be less than this in order to detect a signal.  The thermal SZ is more promising. Future experiments which will be useful for improving constraints on $\\Omega_{b,I}$ include the Planck satellite mission\\footnote{http://planck.esa.int}, the Sunyaev Zel'dovich Array (SZA\\footnote{http://astro.uchicago.edu/sza}), the Atacama Cosmology Telescope (ACT, see Kosowsky 2003), and the South Pole Telescope (SPT\\footnote{http://astro.uchicago.edu/spt}). Given the expected noise characteristics and angular resolution/sky coverage of these instruments, we expect substantial improvements over the constraints derived using WMAP in this paper. We have seen in the WMAP case (Section \\ref{wmaps})  that we are limited both by the noise level ($\\sim60 \\mu $K rms) and the resolution of the data. As an example, the ACT angular resolution is projected to be $\\sim 1$ arcmin with an $rms$ detector noise level of $\\sim 2 \\mu$K per pixel in the 145 GHz band. Even without allowing for the higher resolution or using the full angular cross-correlation information, one would therefore expect to extract the same level \\lya\\ forest thermal SZ signal using $(60/2)^{2}$ fewer spectra (assuming the primary anisotropy can be filtered out in frequency space). The ACT signal will be lower by about $50\\%$ at 145 GHz  than the Rayleigh Jeans prediction, but the other frequency bands can be used to increase the significance of the detection. In order to yield a 2-$\\sigma$ detection of $\\Omega_{b}$, using pixel-pdf statistics, as in this paper, we estimate that $\\sim$ 300 quasar \\lya\\ forest spectra would be needed in the ACT survey area, assuming the BBN value of $\\Omega_{b} h^{2}$ (0.019; Burles \\& Tytler 1998) Using angular cross-correlation information would increase the significance of the detection.  The lack of overlap between the ACT fields and the SDSS would necessitate a southern quasar survey, although it should be borne in mind that only extremely low resolution spectra are needed.  We note that in our analysis we have assumed that the ratio of neutral to ionised hydrogen is that expected if the ionising background radiation field does not vary spatially in  intensity. If there are coherent spatial fluctuations in the photoionisation rate ($\\Gamma$ in Equation 1) then the correlation between \\lya\\ optical depth and CMB temperature decrement will be affected by an additional source of scatter. This is likely to be small, however with the spatial cross-correlation reduced by $5\\%$ or less (see e.g., Croft 2004) and could in principle be calibrated out using theoretical modelling in the manner of McDonald \\etal 2005.  Finally, one can ask why it is important to try to detect the ionised component of the baryons at high redshift when the neutral hydrogen tracer has already been well measured (e.g., Rauch \\etal 1997).  One reason is that a detection would act as a consistency check and enable us to verify directly the baryon inventory at these redshifts. Also, at the moment, the best constraints on the baryon density in observed structures (if we set aside for the moment relatively indirect CMB and Big Bang Nucleosynthesis, BBN measurements e.g., Spergel \\etal 2003, Walker \\etal 1991) comes from looking at the mean absorption level in the \\lya\\ forest (e.g., Rauch \\etal 1997, Weinberg \\etal 1997, Hui \\etal 2002, Tytler \\etal 2004).  However, these measurements are degenerate with constraints on the ionising background radiation intensity $J$.  Often, $\\Omega_{b}$ from BBN or CMB is assumed and then the \\lya\\ forest measurement is used to constrain the ionising radiation intensity (e.g., Kirkman \\etal 2005).  If we use actual measurements of $J$ (from e.g., the proximity effect e.g., Scott \\etal 2000) then constraints on $\\Omega_{b}$ directly seen in the \\lya\\ forest become weaker ($J$ is only known to within roughly a factor of two). Measurements of the ionised baryon fraction such as those we propose may therefore become competitive."
    },
    "0512/astro-ph0512455_arXiv.txt": {
        "abstract": "Seed black holes formed in the collapse of population III stars have been invoked to explain the presence of supermassive black holes at high redshift. It has been suggested that a seed black hole can grow up to $10^{5\\sim 6}\\sunm$ through highly super-Eddington accretion for a period of $\\sim 10^{6\\sim 7}$ yr between redshift $z=20\\sim 24$. We studied the feedback of radiation pressure, Compton heating and outflow during the seed black hole growth. It is found that its surrounding medium fueled to the seed hole is greatly heated by Compton heating. For a super-critical accretion onto a $10^3\\sunm$ seed hole, a Compton sphere (with a temperature $\\sim 10^6$K) forms in a timescale of $1.6\\times 10^3$yr so that the hole is only supplied by a rate of $10^{-3}$ Eddington limit from the Compton sphere.  Beyond the Compton sphere, the kinetic feedback of the strong outflow heats the medium at large distance, this leads to a dramatical decrease of  the outer Bondi accretion onto the black hole and avoid the accumulation of the matter. The highly super-critical accretion will be rapidly halted by the strong feedback. The seed black holes hardly grow up at the very early universe unless the strong feedback can be avoided. ",
        "introduction": "Black holes are regarded as an extremely important population in modern cosmological physics. The reionization of the Universe may get started from $z\\sim 17$ deduced from {\\em Wilkinson Microwave Anisotropic Probe} ({\\em WMAP})  which gives Thomson scattering depth $\\tau=0.17\\pm 0.04$ (Spergel et al. 2003). Such an early reionization epoch needs a large population of seed black holes collapsed from population III stars (Madau et a. 2004). Second, the discovery of the currently known highest redshift quasar, SDSS 1148+3251 at $z=6.4$ (roughly 1 Gyr) from {\\em Sloan Digital Sky Survey} (SDSS; Fan et al. 2001) indicates that there are already supermassive black holes with $\\mbh>10^9\\sunm$ (Netzer 2003, Barth et al. 2003, Willott et al. 2003). What is the relation between the seed and supermassive black holes? How to form supermassive black holes?  Rees' diagram shows several possible ways to form supermassive black holes (Rees 1984). A direct collapse of primordial clouds could form supermassive black holes after cosmic background radiation photons remove enough angular momentum through Compton drag (Loeb 1993, Loeb \\& Rassio 1994). This scenario is favored by Ly$\\alpha$ fuzz of the extended emission in quasar where the supermassive black hole has been formed and the galaxy is assembling (Weidinger, Moller \\& Fynbo 2004), especially the recent discovery of an isolated black hole of the quasar HE 0450-2958 without a massive host galaxy (Magain et al. 2005). Second, a rapid growth of a seed black hole with highly super-Eddington accretion rates is used to explain the existence of the black hole $>10^9\\sunm$ at high redshift.  Third a compact cluster of main sequence stars, or neutron stars/black holes will inevitably evolve into a supermassive black hole (Duncan \\& Shapiro 1983, Quinlan \\& Shapiro 1990), and this gets supports from the quasar's metallicity properties (Wang 2001). The different ways to form a supermassive black hole may apply to different redshifts or environments.  A rapid growth of seed black holes is quite a promising model to issue the formation of supermassive black holes at redshift $z\\sim 6$ (Volonteri \\& Rees 2005). The Bondi accretion rate is $\\dot{m}=\\dotmb/\\dotmedd\\sim 40$, where $\\dot{M}_{\\rm Edd}=L_{\\rm Edd}/c^2$ and $L_{\\rm Edd}$ is the Eddington limit,  for a $10^3\\sunm$ black hole surrounded by medium cooled by the hydrogen atomic lines. The seed black hole is able to grow up exponentially $\\mbh=M_0\\exp(\\dot{m} t/t_{\\rm Salp})$, where the Salpeter timescale $t_{\\rm Salp}=0.45$Gyr. However, such a high accretion rate inevitably gives rise to strong interactions of radiation and outflows with the Bondi accretion flow. Consequently, the strong feedback seriously constrains the matter supply to the black hole, even stops the accretion.  In this Letter we discuss how the feedback impacts on the growth of the seed black holes. We find they can not grow up between redshift $z=20\\sim 24$ through accretion. The implications of the present results are discussed. ",
        "conclusions": "We show that the feedback of the radiation pressure, Compton heating and outflows from the super-critical accretion disks will result in strong influence on the accretion itself. We show that the Compton heating almost quenches the super-critical accretion and the outflow from the tiny disk heats up the outer region so that accumulation of matter is avoided.  A rapid growth of seed black holes is greatly suppressed and they are hardly able to grow up unless the strong feedback can be avoided, for example considering the anisotropy of the radiation from the tiny accretion disk."
    },
    "0512/astro-ph0512039_arXiv.txt": {
        "abstract": "Supernovae are essential to understanding the chemical evolution of the Universe.  Type~\\,Ia supernovae also provide the most powerful observational tool currently available for studying the expansion history of the Universe and the nature of dark energy.  Our basic knowledge of supernovae comes from the study of their photometric and spectroscopic properties.  However, the presently available data sets of optical and near-infrared light curves of supernovae are rather small and/or heterogeneous, and employ photometric systems that are poorly characterized.  Similarly, there are relatively few supernovae whose spectral evolution has been well sampled, both in wavelength and phase, with precise spectrophotometric observations.  The low-redshift portion of the Carnegie Supernova Project (CSP) seeks to remedy this situation by providing photometry and spectrophotometry of a large sample of supernovae taken on telescope/filter/detector systems that are well understood and well characterized. During a five-year program which began in September 2004, we expect to obtain high-precision $u'g'r'i'BVYJHK_s$ light curves and optical spectrophotometry for about 250 supernovae of all types. In this paper we provide a detailed description of the CSP survey observing and data reduction methodology.  In addition, we present preliminary photometry and spectra obtained for a few representative supernovae during the first observing campaign. ",
        "introduction": "A universe dominated by normal mass should undergo deceleration as it expands.  Thus, the counter-intuitive discovery of an accelerating universe based on observations of Type~\\,Ia supernovae \\citep{riess98,perlmutter99} is of profound significance for physics.  Evidently ``dark energy,'' in the form of Einstein's cosmological constant or a more general scalar energy field, is the dominant mass/energy constituent of the Universe today.  These important implications depend critically on the quality of the light curves of the Type~\\,Ia supernovae (SNe~Ia, hereafter) and the ability to K-correct, deredden, and normalize these to a standard luminosity.  The evidence for an accelerating universe is based on a differential measurement between local and distant SNe~Ia (at lookback times of 4--10 Gyr). The local samples are very heterogeneous, and as more SNe have been added, the full sample dispersion around the local Hubble flow has increased from 0.12 \\citep{phillips99} to 0.18 mag \\citep{jha02} in units of distance modulus.  Moreover, there are still legitimate concerns about possible systematic errors due to poorly understood photometric systems \\citep{suntzeff00,stritzinger02}, reddening corrections \\citep{phillips99}, and evolutionary effects \\citep{hamuy00,gallagher05}.  A new and larger sample of nearby ($z < 0.07$) SNe, where these sources of observational error have been duly accounted for, is urgently needed.  With that purpose in mind, we have initiated a five-year program, the Carnegie Supernova Project (hereafter CSP), to obtain well-calibrated optical and near-infrared light curves as well as optical spectrophotometry of $\\sim$250 Type~\\,Ia and core-collapse SNe. The CSP is built upon the unique resources of the Las Campanas Observatory (LCO) in Chile. We have guaranteed access to large numbers of nights on the Swope 1~m and the duPont 2.5~m telescopes ($\\sim$300 per year in both telescopes together), which are equipped with high-performance CCD optical imagers, near-infrared (NIR) cameras, and CCD optical spectrographs. In addition to providing densely sampled light curves covering the near-ultraviolet to the NIR ($u'g'r'i'BVYJHK_s$), we have the means to obtain optical spectrophotometry at approximately weekly intervals. The CSP is a follow-up project and relies on SNe discovered in the course of other surveys. A large fraction of our targets come from the Lick Observatory Supernova Search \\citep{li00,filippenko01,filippenko03,filippenko05} conducted with the Katzman Automatic Imaging Telescope (KAIT), and from dedicated SN searches by amateur astronomers (e.g., Tim Puckett, Tom Boles, Berto Monard, Koichi Itagaki), which constitute a growing source of nearby SNe. The targets selected for the follow-up observations by the CSP are SNe discovered before or near maximum light with $z \\lesssim 0.07$ and $\\delta$ $\\lesssim$ +20$^\\circ$.  The primary goal of the CSP is to establish a fundamental data set of optical and NIR light curves in a well-defined and well-understood photometric system for all types of SNe.  A secondary goal is to obtain complementary optical spectrophotometry for these same SNe. The data set for the Type~\\,Ia events will allow us to improve extinction corrections and to investigate systematic effects possibly due to differences in age and metallicity. The data for the Type~\\,II\\,~SNe will be used to establish and refine precise techniques for measuring luminosity distances employing the Expanding Photosphere Method \\citep{eastman96, schmidt94, hamuy01, leonard02, dessart05} or the Standardized Candle Method \\citep{hamuy02}, thereby providing an independent check on the Type~\\,Ia results.  We will be able to explore the use of both SN types for studies of local galaxy flows and independently measure the convergence depth (the distance at which the bulk flows smooth out into the so-called large-scale Hubble flow). Ultimately, the data set will serve as a reference for observations of distant SNe that will be obtained in coming years in the course of the Joint Dark Energy Mission \\citep{aldering05} and those being obtained in the course of the high-$z$ SN surveys such as the CFHT Legacy project \\citep{pritchet05}, ESSENCE \\citep{matheson05}, and our own high-$z$ component of the CSP which seeks to measure rest-band $I$ magnitudes of SNe~Ia at $z \\approx 0.5$ using the Magellan telescopes \\citep{freedman05}.  The low-$z$ CSP data set will also allow us to gain a deeper understanding of the physics of thermonuclear (Type~\\,Ia) events and the different classes of core-collapse SNe (Types~\\,II,~\\,Ib,\\,~Ic). For example, during the first CSP observing campaign, we obtained excellent coverage of SN\\,~2005bf, a peculiar, luminous Type~\\,Ic event which peaked 35 days after explosion and which may represent a transition object between the SNe associated with gamma-ray bursts and ordinary SNe~Ib \\citep{folatelli05}.  The purpose of this paper is to describe the low-$z$ CSP experiment, to explain the general procedures for data acquisition and reduction, to summarize the results obtained during the first (2004--2005) low-$z$ CSP observing campaign, and to present the data for a few representative SNe. In $\\S$ \\ref{inst} we discuss the instrumentation, in $\\S$ \\ref{obs} we describe our observations, in $\\S$ \\ref{red} we explain the data reduction procedures, and in $\\S$ \\ref{res} we show representative light curves and spectra obtained thus far. ",
        "conclusions": "The low-$z$ program of the CSP is underway. During the first 9-month campaign, we were assigned 190 nights with the LCO 1~m telescope, 57 nights with the 2.5~m telescope, and a smaller number of nights with the two Magellan 6.5~m telescopes and the CTIO 1.5~m telescope.  With this allocation, we were able to obtain follow-up photometry and spectroscopy for 38 SNe (17 SNe~Ia, 12 SNe~II, and 9 SNe~Ibc).  In the first year of operation, we developed reduction pipelines and software which allowed us to produce optical+NIR light curves in real time and post them on our public web site. Thanks to the nearly uninterrupted access to the Swope 1~m telescope, our optical light curves have an unprecedented gap-free temporal coverage.  This data set constitutes the first ever for SNe in the SDSS filters. Given that we used only one instrument and filter set, the light curves are very homogeneous in comparison to data sets in the literature such as the ``gold sample'' of \\cite{riess04}. Through careful attention to details, precisions of 0.03 mag in $u'$ and 0.01 mag in $g'r'i'BV$ in single measurements have been achieved. We were able to sample the NIR light curves of SNe~Ia every 5--7 days with longer gaps of 15 days during dark time, and those of SNe~II every 10 nights. The precision of single measurements in the $YJH$ bands was typically 0.02--0.03 mag, consistent with expectations.  Our data processing procedures allowed us to flux and wavelength calibrate hundreds of spectra in a timely manner. During the first campaign, the CSP provided spectroscopic types for 27 SNe. These spectra will serve as a valuable resource for improving K-corrections for SNe~Ia and SNe~II, as well as for measuring expansion velocities and line strengths in order to explore correlations with SN luminosities and thus refine the methods for distance determination."
    },
    "0512/astro-ph0512507_arXiv.txt": {
        "abstract": "We explore whether stellar tidal streams can provide information on the secular, cosmological evolution of the Milky Way's gravitational potential and on the presence of subhalos. We carry out long-term ($\\Delta t\\sim t_{\\rm hubble}$) N-body simulations of disrupting satellite galaxies in a semi-analytic Galaxy potential where the dark matter halo and the subhalos evolve according to a $\\Lambda$CDM cosmogony. All simulations are constrained to end up with the same position and velocity at present. Our simulations account for: (i) the secular evolution of the host halo's mass, size and shape, (ii) the presence of subhalos and (iii) dynamical friction.\\\\ We find that tidal stream particles respond adiabatically to the Galaxy growth so that, at present, the energy and angular momentum distribution is exclusively determined by the present Galaxy potential. In other words, all present-day observables can only constrain the present mass distribution of the Galaxy independent of its past evolution. We also show that, if the full phase-space distribution of a tidal stream is available, we can accurately determine (i) the present Galaxy's shape and (ii) the amount of mass loss from the stream's progenitor, even if this evolution spanned a cosmologically significant epoch. ",
        "introduction": "\\label{sec:int} In the last decade stellar streams in and around the Milky Way, which are possible debris from the disruption of satellite galaxies during the hierarchical assembly of our Galaxy, have become an active topic of investigation for several reasons. Firstly, large scale CCD surveys have provided unprecedented evidence of accretion and tidal disruption of dwarf galaxies around large spirals in the Local Group (Milky Way: see Majewski 2004; M31: Ibata et al. 2002 and beyond: Pohlen et al. 2004). Secondly,  because tidal streams provide strong constraints on the potential of host galaxies, it is possible to estimate the shape of dark matter halos on large scales in contrast to traditional tracers, such as HI or stellar kinematics (see Sackett et al. 1999 for a review), which provide estimates on relatively small scales. The expected shape of dark matter halos depends on the nature of dark matter particles (see, for example, Dubinsky \\& Calberg 1991, Yoshida et al. 2000 and Dav\\'e et al. 2001 for shape estimates for cold, self-interacting and hot dark matter models, respectively), tidal streams represent a useful tool to discriminate between different paradigms. In addition, tidal stream properties also depend on the mass and internal structure of its progenitor (e.g Johnston, Sackett \\& Bullock 2001 and Law, Johnston \\& Majewski 2004, hereinafter LJM), which ultimately results in a complementary way of estimating the progenitor mass and therefore its mass-to-light ratio. Finally, tidal streams can be used to determine the position of the progenitor if this has not previously been detected (e.g. Font et al. 2004, Pe\\~narrubia et al. 2005).   The formation of tidal streams is, conceptually, a simple process: along its orbit, a stellar system crosses regions where the tidal force of its host galaxy supplies kinetic energy to the initially bound particles. If the energy gain is large enough, particles can become unbound and escape from the host system, forming two sub-systems which are kinematically well differentiated: the leading and the trailing tails which, as their names indicate, precede and follow, respectively, the progenitor in its orbit. The orbital evolution of stripped particles is initially similar to that of the progenitor (e.g. Lynden-Bell \\& Lynden-Bell 1995, Johnston 1998) although, as they evolve in the host galaxy potential, orbits diverge from each other monotonically with time. However, even after a large number of orbital periods, tidal stream particles and the progenitor system reside in well-defined regions of the constant of motion space (Helmi \\& White 1999), therefore attesting a common origin.  The complexity of the stream formation and evolution forces the use of N-body calculations in most cases. The existing work can be divided into a) live N-body simulations, where the host galaxy is formed by a given number of particles initially in equilibrium and b) simulations where the host galaxy is represented by a non-responsive potential. Whereas the former takes into account the host galaxy's response to the satellite, the later neglects this in order to save computational resources for extensive orbit surveys. \\\\ However, none of the N-body simulations of tidal stream formation and evolution to date has accounted for the overall build-up of the host galaxy during the 1--10 Gyr that the stream formation may take. % Yet, in the commonly accepted hierarchical scenario, host galaxies experience large changes in mass, size and shape during their history and hence a tidal stream evolving in an unchanging host galaxy can only approximate recent epochs. The main goal of this contribution is to address the effect that the secular evolution of the host galaxy induces on the formation, evolution and interpretation of tidal streams. Specifically, we will explore whether the present-day structure of an extensive tidal stream can constrain the past history of the host's gravitational potential.  In addition to the secular overall mass and size growth of the host's halo, we will also examine the influence of dark matter substructures on tidal streams (Ibata et al. 2002 and Johnston, Spergel \\& Haydn 2002). CDM cosmology predicts a large number of sub-structures in a galaxy-sized halo, many more than the number of observed dwarf galaxies (Klypin et al. 1999, Moore et al. 1999). Recently, it has been proposed that the process of re-ionization in the universe would lead to a significant decrease in the number of ``visible'' sub-structures (Bullock, Kravtsov \\& Weinberg 2000, Benson et al. 2002, Somerville 2002, Tully et al. 2002), while keeping the total number of sub-structures unaltered. This process would establish a minimum mass (corresponding to a minimum circular velocity of $\\sim 30$ km/s) above which gravitationally bound systems would be able to retain baryonic matter and, thus, to form stars. \\\\ The properties of kinematically cold tidal stream are strongly sensitive to the lumpiness of the galaxy potential (Ibata et al. 2002, Johnston, Spergel \\& Haydn 2002) as repeated encounters with dark matter sub-structures alter the energy and angular momentum distribution of tidal stream particles, leading to hotter, broadly dispersed streams. Tidal streams appear to be a unique laboratory to determine the presence of dark matter clumps in galaxy halos. Firstly, because the effects of those clumps on tidal streams are purely gravitational (and so, independent of whether sub-structures enclose baryons or not) and, secondly, because tidal streams can be detected on large scales and can be as old as the host galaxy, and may therefore provide information on the number and spatial distribution of bound sub-structures at different epochs. Yet, the calculations mentioned above do not take into account the evolution of the spatial distribution nor the mass loss of substructures, which may weaken the influence of substructures on the tidal stream evolution.     In this paper we focus on tidal streams in the Milky Way simply because only for our Galaxy can streams be resolved into stars and accurate phase-space information gathered.  The remainder of this paper is arranged as follows. In \\S\\ref{sec:galmod} we describe our models for the Milky Way potential and for the satellite galaxies that are disrupted to form tidal streams; \\S\\ref{sec:code} details our N-body code used for evolving satellite orbits. \\S\\ref{sec:calc} describes the specific set of orbits we explored in this work. \\S\\ref{sec:mass} describes how we address the problem of the unknown mass loss history of the satellite. In \\S\\ref{sec:progorb} we explore the behavior of the satellite galaxy's orbit while in \\S\\ref{sec:tsprop} we explore the properties of the associated tidal streams. In \\S\\ref{sec:shape} we explore whether evolution of the halo shape can influence the properties of tidal streams. \\S\\ref{sec:clumps} examines the influence of dark matter substructures of these tidal streams.  Finally, in \\S\\ref{sec:disc} we present our conclusions. ",
        "conclusions": "\\label{sec:disc} In this paper we have analyzed what information can be extracted from stellar tidal streams on (i) the Milky Way's halo shape, (ii) the halo's secular evolution, (iii) the mass evolution of the stream's progenitor and (iv) the presence of dark matter clumps in our Galaxy. We assumed as a boundary condition of this analysis that the present-day position, velocity and mass of the stream's progenitor can be measured and are identical for all evolutionary scenarios. Under these conditions, we have explored whether the extended tidal debris reflect differences arising from the items (i)--(iv) above. We have carried out this study for a Sagittarius-like dwarf galaxy, although our results are general and can be applied to other systems in the Milky Way.  The main result is that tidal streams do not provide information on the adiabatic evolution of the Milky Way or, in other words, the properties of entire tidal streams only reflect the {\\it present} Galaxy potential. Thus, ground-based observations already available for tidal streams (basically providing distances and radial velocities along the stream) and future satellite data covering the full phase-space (making possible studies in the $E$--$L_z$ plane) can only constrain the present characteristics of the Milky Way potential. As a direct consequence, Galaxy evolution processes can be neglected when modeling tidal streams, which clarifies one of the main caveats in current N-body simulations and confirms that tidal stream models computed under that hypothesis are appropriate  for tracing the distribution of dark matter around our Galaxy.   In contrast, we find that tidal streams are indeed sensitive to the present properties of halo shape. In particular, we confirm that measuring the precession rate is a fairly powerful method to constrain the halo flattening of the gravitational potential. This may not require measurements of proper motions.   We have shown that the study of tidal streams in the $E$--$L_z$ plane provides information on the progenitor's mass loss since the time of accretion.  The energy--angular momentum distribution of stream particles has an average value that only reflects the present position and velocity of the progenitor system. In contrast, the $E$--$L_z$ variance about that mean increases with the initial satellite mass, thus, making it possible to determine the mass loss fraction directly from the present $E$--$L_z$ distribution. Furthermore, since a secular drift in energy and angular momentum is induced by dynamical friction (a drag acceleration that scales in proportion to $M_{\\rm s}$) one could, in principle, reconstruct the mass loss curve $M_s(t)$ from the age of different stream pieces (if the age can be labeled e.g. by metalicity or by theoretical modeling), which ultimately depends on the initial mass profile of the satellite galaxy (see Zhao 2004).   We have analyzed the effects of dark matter clumps on tidal streams in Section~\\ref{sec:clumps}. We simplify the problem by assuming that dark matter clumps do not alter the progenitor's orbit, but only the orbits of stream particles. That approach establishes, therefore, a minimum impact of subhalos on tidal stream properties (we note that a sharp change of the progenitor's $E_p, L_{z,p}$ induced by a collision with a subhalo at $t=t_c$ would be reflected as a discontinuity in the averaged stream's $E, L_z$ at age=$t_f-t_c$). We have confirmed that dark matter subhalos induce only very modest stream ``heating'' by increasing of the angular momentum dispersion in the oldest (age$>$5 Gyr) and coldest (trailing tail) stream parts.\\\\ This raises the question of whether one can constrain the halo lumpiness either with current ground-based techniques or with future astrometric satellite missions (GAIA, SIM).\\\\ At present, the detection of the oldest parts (i.e those stars that became unbound first) of a tidal stream with state-of-the-art ground-based surveys is challenging. Although the trailing tail maintains a coherent structure and should be easier to detect as spatial over-densities (see Sec.~\\ref{sec:xyz}), its surface brightness decreases considerably with time, which reduces the possibility of detection above the Galactic field contamination  in large field-of-view color-magnitude diagrams like those provided by SDSS. Furthermore, it is also expected that tidal streams are composed by old, metal-poor stellar population. Therefore, some  valuable techniques for tracing tidal streams with all-sky surveys (like 2MASS) cannot be applied to detect the oldest stream pieces since these are barely sensitive to metal-poor stars expected in old parts of tidal streams. Also surveys of tidal streams using M-giant stars (Majewski et al. 2004) are limited to the youngest stream pieces (material unbound only 1-2 Gyr ago; LJM), while the oldest wraps still remain hidden in the Galactic halo. Current theoretical models of the two largest, brightest streams in the Milky Way (Sgr: LJM;  Monoceros: Pe\\~narrubia et al. 2005) indicate that there is no detection of tidal debris that became unbound more than 2-3 Gyr ago.  Finally, stream tails are not located on the progenitor's orbital plane (see Sec.~\\ref{sec:orbplane}), which further complicates tagging of debris as part of a known stellar system and needs of accurate stream models in order to determine a possible common origin and to estimate the stream's age. \\\\ In the future, the most powerful method to search for ancient debris in the halo should come with the next generation of astrometric satellites, which will permit analysis of tidal streams in the constant-of-motion space, thus, providing the most straightforward way to identify tidal debris independently of the stream age. However, it is important to remark that observational errors may introduce strong limitations, as indicated by Brown et al. (2005). Possibilities of identifying satellite remnants in the Milky Way halo are considerably reduced after taking into account the observational errors expected for the GAIA catalogue and the large number of background stars."
    },
    "0512/astro-ph0512488_arXiv.txt": {
        "abstract": "% Combination of high-precision photometry and spectroscopy allows the detailed study of the upper main sequence in open clusters. We are carrying out a comprehensive study of a number of clusters containing Be stars in order to evaluate the likelihood that a significant number of Be stars form through mass exchange in a binary. Our first results show that most young open clusters contain blue stragglers. In spite of the small number of clusters so far analysed, some trends are beginning to emerge. In younger open clusters, such as NGC~869 and NGC~663, there are many blue stragglers, most of which are not Be stars. In older clusters, such as IC~4725, the fraction of Be stars among blue stragglers is very high. Two Be blue stragglers are moderately strong X-ray sources, one of them being a confirmed X-ray binary. Such objects must have formed through binary evolution. We discuss the contribution of mass transfer in a close binary to the formation of both blue stragglers and Be stars ",
        "introduction": "% Blue stragglers (BSs) are stars lying above the main sequence (MS) turnoff region in colour-magnitude diagrams, a region where, if the BSs had been normal single stars, they should already have evolved away from the main sequence \\citep{stry}. Several mechanisms have been proposed to explain the formation of BSs in different environments. \\begin{itemize} \\item BSs may be stars formed later than the rest of the cluster or association. Though this may well happen in some regions with sequential star formation, in many clusters there are no stellar sequences connecting the BSs with the turn-off, arguing against a second epoch of star formation. \\item BSs might be evolved stars back in the blue region of the HR diagram after a red supergiant phase. However, abundances of Nitrogen in blue stragglers are much lower than those predicted by models for stars in blue loops \\citep{smartt}. \\item BSs may be stars that, for some reason, have evolved bluewards. In particular {\\it homogeneous evolution} has been proposed as a mechanism that can result in blueward evolution for very high (near-critical) initial rotational velocity \\citep{maeder}. However there is no strong {\\it a priori} reason to believe that many such extreme rotators will be formed. \\item BSs may be formed by coalescence of two stars. This mechanism requires a high stellar density environment. Though it is probably the major channel for the formation of BSs in globular cluster, it is unlikely to be able to explain BSs in OB associations. \\item BSs result from mass exchange in close binaries. Examples of this process abound amongst massive binaries (see review by Negueruela in these proceedings). Therefore this channel must contribute some BSs in young clusters and associations. \\end{itemize}  While the properties of BSs in globular clusters have been the object of many studies, the mechanisms for forming BSs in young open clusters have deserved little attention. The most relevant works are those of \\citet{mermi} and \\citet{mathys} ",
        "conclusions": "Though the sample of open clusters for which a detailed analysis has been made is still quite small, we observe the following interesting trends:  \\begin{itemize} \\item Blue stragglers are found in all open clusters surveyed \\item In clusters $\\sim15\\:$Myr old, most evolved stars appear too massive for these ages. In particular, many of these clusters  contain Ia and Iab supergiants, which must be descended from stars with $\\ga20M_{\\sun}$ according to all theoretical evolutionary paths. Such massive stars have lifetimes $<10\\:$Myr \\citep{mm03}, rendering all these supergiants BSs. \\item Among clusters younger than $\\sim25\\:$Myr, very few BSs are Be stars. \\item From the data on IC~4725 and NGC~6649, incomplete samples in NGC~2516 and IC~2488 and data in the literature \\citep{mermi}, it would seem that most BSs in clusters in the $50-150$~Myr range are Be stars. \\item Our data suggest that a non-negligible fraction of BSs in moderately young open clusters are formed by mass transfer in a close binary (X-ray sources in NGC~663 and NGC~6649, eclipsing binary in NGC~3766), but not all BSs are binaries, making it unlikely that this might be the only channel. \\item Likely different mechanisms are dominant at different ages. \\end{itemize}"
    },
    "0512/astro-ph0512441_arXiv.txt": {
        "abstract": "We report on {\\em XMM-Newton} observations of the young open cluster NGC~2547 which allow us to characterise coronal activity in solar-type stars, and stars of lower mass, at an age of 30\\,Myr.  X-ray emission is seen from stars at all spectral types, peaking among G-stars at luminosities (0.3--3\\,keV) of $L_{\\rm x}\\simeq 10^{30.5}$~erg\\,s$^{-1}$ and declining to $L_{\\rm x} \\leq 10^{29.0}$~eg\\,s$^{-1}$ among M-stars with masses $\\geq 0.2\\,M_{\\odot}$. Coronal spectra show evidence for multi-temperature differential emission measures and low coronal metal abundances of $Z\\simeq 0.3$. The G- and K-type stars of NGC~2547 follow the same relationship between X-ray activity and Rossby number established in older clusters and field stars, although most of the solar-type stars in NGC~2547 exhibit saturated or even super-saturated X-ray activity levels. The median levels of $L_{\\rm x}$ and $L_{\\rm x}/L_{\\rm bol}$ in the solar-type stars of NGC~2547 are very similar to those in T-Tauri stars of the Orion Nebula cluster (ONC), but an order of magnitude higher than in the older Pleiades. The spread in X-ray activity levels among solar-type stars in NGC~2547 is much smaller than in older or younger clusters.  Coronal temperatures increase with $L_{\\rm x}$, $L_{\\rm x}/L_{\\rm bol}$ and surface X-ray flux. The most active solar-type stars in NGC~2547 have coronal temperatures intermediate between those in the ONC and the most active older ZAMS stars. We show that simple scaling arguments predict higher coronal temperature in coronally saturated stars with lower gravities.  A number of candidate flares were identified among the low-mass members and a flaring rate (for total flare energies [0.3--3\\,keV] $>10^{34}$~erg) of 1 every $350^{+350}_{-120}$\\,ks was found for solar-type stars, which is similar to rates found in the ONC and Pleiades. Comparison with {\\it ROSAT} HRI data taken 7 years previously reveals that only 10--15 percent of solar-type stars or stars with $L_{\\rm x}>3\\times10^{29}$~erg\\,s$^{-1}$ exhibit X-ray variability by more than a factor of two. This is comparable with clusters of similar age but less than in both older and younger clusters. The similar median levels of X-ray activity and rate of occurrence for large flares in NGC~2547 and the ONC demonstrate that the X-ray radiation environment around young solar-type stars remains relatively constant over their first 30\\,Myr. ",
        "introduction": "X-ray emission from the hot coronae of cool stars is now a well established phenomenon (e.g. see the review by G\\\"udel 2004). The emission arises from magnetically confined and heated structures with temperatures in excess of $10^{6}$\\,K. In stars that have reached the zero age main sequence (ZAMS) or older, the driving mechanism for this magnetic activity is thought to be a stellar dynamo: stars with convective envelopes and rapid rotation are relatively luminous X-ray sources compared with slower rotating stars of similar spectral type. There is now a well-founded age-rotation-activity paradigm (ARAP -- see Jeffries 1999; Randich 2000), established via observations of many open clusters with ages from 50\\,Myr to several Gyr (e.g. Stauffer et al. 1994; Stern, Schmitt \\& Kahabka 1995; Jeffries, Thurston \\& Pye 1997), whereby younger stars tend to be more rapidly rotating and hence exhibit strong X-ray emission up to a saturated level, where the ratio of X-ray to bolometric flux, $L_{\\rm x}/L_{\\rm bol} \\simeq 10^{-3}$. As stars get older, they lose angular momentum and eventually spin down to rotation rates where $L_{\\rm x}/L_{\\rm bol}<10^{-3}$ and decreases further thereafter.  For very young stars in star forming regions with ages $<10$\\,Myr, a direct connection between rotation and X-ray activity is much less clear and the presence of an (accretion) disc may play a role in either stimulating or inhibiting the observed levels of X-ray activity (see Feigelson et al. 2003; Flaccomio et al. 2003a; Flaccomio, Micela \\& Sciortino 2003b; Stassun et al. 2004; Preibisch et al. 2005). Feigelson et al. (2003) find no correlation between rotation and X-ray activity and a ``saturation level'' of only $L_{\\rm x}/L_{\\rm bol}\\simeq 10^{-3.8}$ for pre-main-sequence (PMS) stars, both with and without discs, in the Orion Nebula Cluster (ONC). They suggest that a less efficient, turbulent ``distributed'' dynamo may act throughout the convective zones of these stars. On the other hand Flaccomio et al. (2003a,b) and Stassun et al. (2004) suggest that ONC stars with accretion discs bias the average $L_{\\rm x}/L_{\\rm bol}$ downwards, perhaps as a result of intrinsic absorption or changes in the magnetic field geometry.  Preibisch et al. (2005) show that active accretion, rather than the mere presence of a disc is possibly responsible for the wider spread and lower median level of X-ray activity among the very young ONC stars.  NGC 2547 is an interesting open cluster in the context of studying the transition between the early behaviour of stellar coronae in star forming regions and the development of the well tested ARAP at older ages.  It has a precisely determined age of either $30\\pm5$\\,Myr determined from fitting isochrones to its 0.3--1.2$M_{\\odot}$ stars as they descend their PMS tracks (Naylor et al. 2002), or $35\\pm 3$\\,Myr determined from the re-appearance of lithium in the atmospheres of even lower mass stars (Jeffries \\& Oliveira 2005). It is old enough that inner circumstellar discs have dispersed -- no accretion-related H$\\alpha$ emission or $L$-band near infrared excesses are seen from its solar-type members (e.g. Jeffries, Totten \\& James 2000; Young et al. 2004).  However, cluster members with $M<1.4\\,M_{\\odot}$ are still in the PMS phase, stars with $M<0.4\\,M_{\\odot}$ are fully convective (Siess, Dufour \\& Forestini 2000; D'Antona \\& Mazzitelli 1997) and it is significantly younger than other well-studied open clusters like IC\\,2391 ($50\\pm 5$\\,Myr) and the Alpha Per cluster ($90\\pm10$\\,Myr).  NGC 2547 was observed at X-ray wavelengths by the {\\it ROSAT} High Resolution Imager (HRI). Jeffries \\& Tolley (1998) found a rich population of low mass, X-ray active cluster candidates with $10^{29}<L_{\\rm x}<10^{31}$\\,erg\\,s$^{-1}$. Puzzlingly, the solar-type stars of NGC 2547 seemed slightly {\\it less} X-ray active than their counterparts in older clusters. Their activity peaked at $L_{\\rm x}/L_{\\rm bol}=10^{-3.3}$, whereas the most active stars at lower masses had $L_{\\rm x}/L_{\\rm bol}=10^{-3}$ as expected.  Jeffries, Totten \\& James (2000) ruled out anomalously slow rotation rates in the cluster as an explanation; they found both very fast and slow rotators among the X-ray selected members.  A further possibility is that the energy distribution of the X-ray emitting plasma is rather different for the solar-type stars of NGC 2547 than for those in older clusters or cooler stars -- the HRI observations had no spectral resolution and so a uniform conversion factor was used to estimate X-ray fluxes from count rates.  In this paper we present the results of an {\\it XMM-Newton} observation of NGC 2547 using the European Photon Imaging Camera (EPIC), which seeks to characterise the coronal emission of solar-type (and lower mass) stars at $\\simeq 30$\\,Myr.  The sensitivity of these X-ray images is better than the {\\it ROSAT} HRI data, enabling us to identify cluster members with lower activity and measure X-ray emission from cluster members with lower mass.  There is also some spectral resolution available with the EPIC data that allows us to test whether the X-ray spectra of the solar-type members of NGC 2547 are significantly different to active stars in other open clusters. Finally, we are able to look for possible variability in the level of X-ray emission of these young stars on a timescale of 7 years, which is comparable with the solar magnetic activity cycle.  Section~2 describes the observations, data analysis and identification of X-ray sources with members of NGC~2547. Section~3 deals with spectral analysis of the X-ray data, whilst section~4 uses the spectral information to calculate intrinsic luminosities and search for evidence of dynamo-related activity. Section~5 looks at the X-ray variability of NGC~2547 members, both within the observation (flares, rotational modulation) and on the longer 7 year timescale. Section 6 places NGC~2547 in context with younger and older clusters and discusses the evolution of X-ray activity, coronal temperatures and coronal variability. Our conclusions appear in section~7. ",
        "conclusions": "With a reasonably long and sensitive {\\it XMM-Newton} observation of NGC~2547 we have been able to go some way towards characterising the X-ray emission from analogues of the Sun (and at lower masses) at ages of $\\simeq 30$\\,Myr. The main findings of this study can be summarised as follows: \\begin{enumerate} \\item Candidate cluster counterparts have been found for 108 X-ray sources. X-ray emission is seen from stars of all spectral types. X-ray luminosities peak among the G-stars at $10^{30.5}$~erg\\,s$^{-1}$ declining to $\\leq 10^{29.0}$~erg\\,s$^{-1}$ among M stars with masses as low as $0.2\\,M_{\\odot}$. The ratio of X-ray to bolometric luminosity increases rapidly with decreasing $T_{\\rm eff}$ among F-type stars, reaching a saturated peak of $L_{\\rm x}/L_{\\rm bol} \\simeq 10^{-3}$ for G stars and cooler.  \\item Optically thin thermal plasma models have been used to fit the X-ray spectra of the most active stars in NGC~2547. Where there were sufficient detected X-ray photons, a multi-temperature model fares better than a single temperature model. Coronal metal abundances are clearly sub-solar. There are reasonable correlations between increasing coronal temperature, indicated by a hardness ratio, and X-ray luminosity, surface flux or $L_{\\rm x}/L_{\\rm bol}$. However, none of these relations is entirely deterministic, with a significant scatter in hardness ratio at any activity level.  \\item The relationship between X-ray activity and Rossby number among the G- and K-type stars of NGC~2547 follows that established in a number of older clusters and field stars. The majority of the solar-type stars in NGC~2547 have saturated or even super-saturated levels of coronal activity that are commensurate with their rotation rates. Based on their saturated activity levels, we predict that none of the K stars in NGC~2547 have spun down below equatorial velocities of 10\\,km\\,s$^{-1}$.  \\item The X-ray luminosity functions (XLFs) of NGC~2547 have been compared with younger and older stars in several mass ranges. For G- and K-type stars the median activity levels in NGC~2547, indicated by $L_{\\rm x}/L_{\\rm bol}$, are very similar to stars of equivalent mass in the Orion Nebula cluster (ONC), but nearly an order of magnitude higher than in the Pleiades at $\\simeq 120$\\,Myr. The X-ray luminosities of G-type stars in NGC~2547 are only marginally lower than their counterparts in the ONC. The spreads in the XLFs of NGC~2547 are much smaller than in older clusters and the ONC. These phenomena are readily interpreted in terms of a magnetic dynamo and the age-rotation-activity paradigm, providing account is taken of the evolution of convection zone parameters and angular momentum loss using Rossby numbers. The reported low activity levels of some very young accreting PMS stars in the ONC do not straightforwardly fit into this scenario.  \\item Coronal temperatures in the most active solar-type stars in NGC~2547 are intermediate between those of older clusters and the very young ONC stars. We show that correlations proposed between coronal temperature and X-ray activity for field stars do not predict the temperatures in NGC~2547 or the much higher temperatures in the ONC.  Simple scaling law arguments predict that coronal temperatures should increase as surface gravity decreases for stars with saturated levels of magnetic activity. These arguments require that the scale heights of coronae in the ONC stars are an order of magnitude larger than in ZAMS stars. Alternative explanations such as enhanced flaring or different modes of dynamo action in younger, fully convective PMS stars are possible, but there is no clear evidence for very hot coronae in the fully convective stars of NGC~2547 and the flaring rate (at least for strong flares) is comparable in the solar-type stars of the ONC, NGC~2547 and the Pleiades (see below). However, the hottest coronal temperatures in our sample do arise from clear flaring behaviour.   \\item A number of candidate flares were detected among the low-mass stars of NGC~2547, with peak luminosities (0.3--3.0\\,keV) exceeding $10^{30}$~erg\\,s$^{-1}$ and integrated flare energies of $10^{34}$--$10^{35}$~erg. The occurrence rate for flares with energy $>10^{34}$~erg is approximately 1 every $350^{+350}_{-120}$\\,ks and comparable to flare frequencies in the ONC and Pleiades.  \\item By comparison with earlier {\\it ROSAT} HRI observations we find that only 10--15 per cent of G- and K-type stars, or stars of any spectral types with $L_{\\rm x}>3\\times10^{29}$~erg\\,s$^{-1}$, show variations of a factor of two or more on timescales of 7 years. The level of long term variability is incompatible with solar-like magnetic activity cycles with periods of 20 years or less. This is comparable with variability seen in similarly aged clusters (IC~2391) but less than observed in older clusters and field stars {\\it and} in the younger PMS stars of the ONC.  \\end{enumerate}"
    },
    "0512/astro-ph0512394_arXiv.txt": {
        "abstract": "The fluctuating accretion model of \\citet{Lyub} and its extension by \\citet{Kotov}, seeks to explain the spectral-timing properties of the X-ray variability of accreting black holes in terms of inward-propagating mass accretion fluctuations produced at a broad range of radii. The fluctuations modulate the X-ray emitting region as they move inwards and can produce temporal-frequency-dependent lags between energy bands, and energy-dependent power spectral densities (PSDs) as a result of the different emissivity profiles, which may be expected at different X-ray energies.  Here we use a simple numerical implementation to investigate in detail the X-ray spectral-timing properties of the model and their relation to several physically interesting parameters, namely the emissivity profile in different energy bands, the geometrical thickness and viscosity parameter of the accretion flow, the strength of damping on the fluctuations and the temporal coherence (measured by the `quality-factor', $Q$) of the fluctuations introduced at each radius.  We find that a geometrically thick flow with large viscosity parameter is favoured, and confirm that the predicted lags are quite robust to changes in the emissivity profile, and physical parameters of the accretion flow, which may help to explain the similarity of the lag spectra in the low/hard and high/soft states of Cyg~X-1.  We also demonstrate the model regime where the light curves in different energy bands are highly spectrally coherent.  We compare model predictions directly to X-ray data from the Narrow Line Seyfert~1 galaxy \\ngc\\ and the BHXRB Cyg X-1 in its high/soft state and show that this general scheme can reproduce simultaneously the time lags and energy-dependence of the PSD. ",
        "introduction": "A common characteristic in Active Galactic Nuclei (AGN) and X-ray binary systems (XRBs) is their strongly variable X-ray emission. Most of the variability power is in the form of aperiodic fluctuations that span several orders of magnitude in temporal frequency \\citep[e.g.][]{vanderKlis,McHardy4051}.  Rapid, large amplitude X-ray variations, on time-scales down to milliseconds in XRBs and minutes in AGN, are commonly observed \\citep[e.g.][]{RevHifreq,GierlinFlares,McHardy4051}, supporting the expectation from accretion theory that the bulk of the emission must originate close to the central black hole.  However, in many sources (in both AGN and XRBs), large-amplitude X-ray variations, of tens of per cent fractional rms, are observed over several decades of time-scales \\citep[e.g.][]{ReigCygX1,McHardy4051,Uttley3227}, including values orders of magnitude longer than the viscous time-scale of the inner disc, suggesting that these fluctuations must originate at large radii. This discrepancy motivated the appearance of propagating-fluctuation models, as proposed by \\citet{Lyub}, who noted that while the \\emph{emission} might be produced only in the innermost regions, the \\emph{variability} can originate throughout the accretion flow. In this model, accretion rate fluctuations arise over a wide range of radii and, correspondingly, a wide range of time-scales, and propagate inwards to modulate the central X-ray emission.  As noted by \\citet{Uttley} the model of \\citet{Lyub} is consistent with the observed linear relation between rms-variability amplitude and mean count rate in XRB and AGN X-ray light curves, which implies a coupling between fluctuations on different time-scales, ruling out models where the variability arises from strictly independent flares or active regions \\citep{UttleyNonlin}.  Fluctuating accretion models, on the other hand, are consistent with this relation: the fluctuations couple together as they propagate down to the centre, so low frequency fluctuations produced at large radii modulate higher frequency fluctuations produced further in.  A linear rms-flux relation is also observed in neutron star XRB systems, where the emission mechanism is most likely different to the one operating in black-hole systems \\citep{Uttley_sax,Uttley}. This result provides additional evidence for fluctuating-accretion scenarios because it suggests that the variability originates in the accretion flow and not just locally via the mechanism producing the emission \\citep{Uttley_sax}.  Besides explaining the broad range of variability time-scales and the rms-flux relation, the fluctuating accretion model can also explain the spectral-timing properties of the variability, as noted by \\citet{Kotov}.  For example, it has long been known \\citep[e.g.][]{Miyamoto,NowakGX339} that X-ray variations in black hole X-ray binaries (BHXRBs) often show hard lags, i.e. a delay in hard X-ray variations with respect to soft X-rays, which is larger for variations on longer time-scales, and at higher energies. The magnitudes of the lags are typically of order one per cent of the variability time-scale.  Similar time-scale-dependent, hard lags have recently been discovered in AGN, albeit on much longer time-scales, commensurate with their higher black hole masses \\citep[e.g.][]{Papadakis,VaughanMCG,McHardy4051}. In their analytical extension to Lyubarskii's model, \\citet{Kotov} explain these lags by invoking a radially extended X-ray emitting region with an energy-dependent profile.  In this scenario, the response of the emission to a fluctuation in the accretion flow is a function of the inward propagation time-scale, and hence radius of origin, of the fluctuation, combined with the emissivity profile. Hence, if the emissivity profile is more centrally concentrated at higher energies, hard-band lags are produced such that the lags are larger for longer time-scale variations.  \\citet{Kotov} show that the same basic picture can also explain the energy dependence of the PSD of BHXRBs and AGN, where there is relatively more high-frequency power observed at higher energies than at lower energies \\citep[e.g.][]{Nowak_lags,McHardy4051}. This is because the emitting region acts as a low-pass filter and, as the emitting region of the soft band is more extended, the higher frequency variations are filtered more strongly in the soft band.  Due to its success in explaining many aspects of the X-ray variability data, the fluctuating accretion model of \\citet{Lyub} and \\citet{Kotov} warrants a deeper investigation, which is the aim of this paper.  As shown by the brief analytical treatment of \\citet{Kotov}, in principle, the model produces time-scale-dependent lags of approximately the same amplitude as observed in the data, and energy-dependent PSDs.  However it is not clear exactly how the lag spectrum (lag as a function of Fourier frequency) and energy dependence of the PSD varies as a function of emissivity profile, or the parameters of the accretion flow (assuming the standard disc model of \\citealt{Shakura} e.g. the viscosity parameter $\\alpha$ or the scale-height of the flow), or due to the effects of radial damping of variations in the accretion flow.  Furthermore, since \\citet{Kotov} analytically determined spectral-timing properties by making the simplifying assumption that the perturbation introduced into the accretion flow at each radius is a  delta-function in time and radius, it is important to determine the effects on spectral-timing properties of a more realistic model, where  stochastic variations over a broader range of time-scales are introduced at each radius. Finally, the fluctuating accretion model should, in principle, produce light curves which are highly correlated in different bands (i.e. spectrally coherent, see \\cite{Vaughan_coh} and Appendix~B of this paper), but this aspect of the model has not yet been studied. To study these various effects in more detail, in Section~\\ref{model} we introduce a computational toy model for a fluctuating accretion flow, based on the work of \\citet{Lyub} and \\citet{Kotov}, and explore the dependence of the PSD, time lags and coherence on the model parameters in  Section~\\ref{properties}. We show how the model predictions fit X-ray data from AGN and the BHXRB Cyg~X-1 in Sections~\\ref{agn} and \\ref{xrb}.  Specifically, we will concentrate on explaining the spectral-timing properties of the BHXRB Cyg~X-1 in its high/soft state, and AGN which show similar variability properties, since these states show rather simple $1/f$-type PSD shapes, without complex quasi-periodic oscillations (QPOs). We discuss the implications of our results in Section~\\ref{discussion} and summarise our results in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} We have used a computational fluctuating-accretion model to reproduce the spectral-timing properties of X-ray light curves from AGN and BHXRBs. In the model, the variability is produced throughout the accretion flow as perturbations of the accretion rate on the local viscous time-scale, produced by many signals which are geometrically spaced in temporal frequency.  The fluctuations propagate inward at the radial drift velocity and modulate the fluctuations produced further in. This general scheme, introduced by \\citet{Lyub}, produces light curves which possess $1/f$ type PSDs over a broad range of time-scales and follow a linear rms-flux relation. Following \\citet{Kotov}, X-ray emission from a radially extended region is assumed to track the local accretion rate fluctuations and to show a power law radial emissivity profile, which steepens for higher energy bands. The different emissivity profiles and the inward propagation produce differences between the PSDs of each energy band and introduce time lags. Our implementation of the model is successful in reproducing the observed variability properties of AGN and BHXRBs, as follows:  \\begin{enumerate} \\item We reproduce the expected $\\sim1/f$ frequency-dependence of time lags between energy bands, and lag amplitudes of a few per cent of the variability time-scales, predicted by \\citep{Kotov} and observed in BHXRB and AGN variability data.  \\item We demonstrate that the amplitudes and slopes of the lag spectra are only weakly dependent on changes in the emissivity indices, the radial dependence of the product of disc parameters $(H/R)^2 \\alpha$, and the strength of damping of accretion fluctuations. The robust nature of the lag spectra may help to explain the similarity of lag spectra in the low/hard and high/soft states in Cyg~X-1 \\citep{Pottschmidt}, despite the very different PSD and energy-spectral shapes in these states.  \\item The extended emitting region suppresses short time-scale fluctuations, producing a gradual bend in the PSD at high frequencies. The `filtered' PSD shapes resemble those of AGN and BHXRBs in the high/soft state (e.g. \\citealt{McHardy4051}) and can be approximately fitted with a bending power-law model.  The bend frequencies obtained match the observed frequencies if the fluctuations are produced on the viscous time-scales of a geometrically thick accretion flow. Values of $(H/R)^2\\alpha \\sim 0.3$ gave adequate bend frequencies compared to data from Seyfert galaxy \\ngc\\ and the BHXRB Cyg~X-1.  \\item The fact that all energy bands are produced by the same emitting region can maintain a high coherence ($>0.95$) between light curves of different energy bands, as observed by e.g. \\citet{Nowak_lags}, \\citet{VaughanMCG}. However, the coherence can be substantially reduced at high frequencies if input signals at each radius have fairly broad PSDs and the difference between emissivity profiles in different bands is large. Coherence is also reduced if damping of the inward-propagating fluctuations is significant.  These results may provide an explanation for the drops in coherence observed at high frequencies in various BHXRBs and AGN \\citep{Nowak_lags, VaughanMCG,McHardy4051}.  \\item The model can simultaneously reproduce both the observed energy-dependence of PSD shape and the lag spectrum for the AGN NGC~4051 and the BHXRB Cyg~X-1 in the high/soft state, although the inferred propagation times in Cyg~X-1 are slightly slower than the drift velocity expected for a standard disc.  In the low/hard state of Cyg~X-1, our simple model can reproduce the general slope of the lag spectrum but cannot account for the stepped shape of the lag spectrum. Modelling the detailed features of this spectrum would require additional assumptions about the mechanism producing the variability.  \\item Direct comparison with spectral-timing data from AGN and BHXRBs shows good general agreement with the model. However, we note that if the model is correct, the empirical bending or broken power-law models used to fit the PSDs of AGN and BHXRBs in the high/soft state must break down for high signal-to-noise data, because they only approximate the gradual bending shape predicted by our model. Also, we note that the model implies that the PSD bend time-scale measured for AGN and BHXRBs is a function of both the underlying PSD of accretion fluctuations (which can contain intrinsic breaks) {\\it and} the emissivity profile of the observed energy band (which introduces a high-frequency bend).  Therefore, the PSD bend time-scale is not necessarily associated with the outer radius of the emitting region, and this region can, in principle, be extremely large, provided the emissivity profile is relatively steep.  \\end{enumerate}"
    },
    "0512/astro-ph0512111_arXiv.txt": {
        "abstract": "{ The fraction of AGN with photoelectric absorption in the X-rays ranging from N$_H$ of 10$^{22}$ up to about 10$^{24}$  cm$^{-2}$ (Compton-thin) appears observationally to be anticorrelated to their luminosity $L_x$. This recently found evidence is used to investigate the location of the absorbing gas. The molecular torus invoked in the unified picture of AGN, while it can be regarded as confirmed on several grounds to explain the Compton-thick objects, do not conform to this new constraint, at least in its physical models as developed so far. In the frame of observationally based evidence that in Compton-thin sources the absorbing gas might be located far away from the X-ray source, it is shown that the gravitational effects of the black hole (BH) on the molecular gas in a disk, within 25-450 pc (depending on the BH mass, from 10$^6$ to 10$^9$ M$_{\\odot}$), leads naturally to the observed anticorrelation, under the assumption of a statistical correlation between the BH mass and $L_x$. Its normalization is also reproduced provided that the surface density, $\\Sigma$, of this gas is larger than about 150-200 $M_{\\odot}$ pc$^{-2}$, and assuming that the bolometric luminosity is one tenth of the Eddington limit. Interestingly, the required values are consistent with the value of the 300 pc molecular disk in our own galaxy, namely 500 $M_{\\odot}$ pc$^{-2}$. In a sample of nearby galaxies from the BIMA SONG survey, it is found that half of the objects have central $\\Sigma$ larger than 150 $M_{\\odot}$ pc$^{-2}$. Given the simplicity of the proposed model, this finding is very encouraging, waiting for future higher resolution surveys in CO on more distant galaxies. ",
        "introduction": "Similar to dust extinction in the optical, photoelectric absorption in X-rays provides clues on the environment of Active Galactic Nuclei (AGN). In general the latter provides information which is more straightforward to interpret than the former, because it depends a) on atomic properties and consequently only on the chemical composition of the medium, b) on the adoption of a shape for the X-ray continuum, which is empirically known to follow a simple Power Law (PL). The values of the spectral photon index, $\\Gamma$, display a rather modest dispersion (e.g. Perola et al. 2002; Piconcelli et al. 2005); furthermore, in any single object under consideration, and provided that the absorption is Compton-thin (i.e. with N$_H$ not exceeding $\\sigma_{Th}^{-1}$=1.5$\\times$10$^{-24}$ cm$^{-2}$), the exact value of $\\Gamma$ can be directly measured, given an adequately wide observational band and sufficient statistics for the spectral counts. Here, though, it is more important to stress the limitations, which are mainly of three sorts:\\begin{enumerate} \\item The source of the PL photons is, for most practical purposes,  point-like, and therefore it remains questionable how to relate the total absorbing matter density (the absorbing column N$_{H}$) to bi-dimensional information (either direct or indirect) on the circumnuclear matter from observations in other wave-bands. In other words, usually it cannot be excluded that the observed column might be pertaining to that unique line of sight, rather than to the overall configuration of the circumnuclear matter. \\item While it is spectroscopically possible (as amply demonstrated even before the use of very high resolution instruments) to disentangle the imprints of a multi-phase (in ionization terms) absorber, it is by no means immediately possible to recover the space distribution of the phases along the line of sight. When the ionization is attributed to the photoelectric effects of the PL source as a reasonable hypothesis (better still when it can be observationally proven not to be of collisional origin), it is generally regarded as most likely that the high ionization gas is very close to the centre. This leaves however quite open the possibility that the low (in practice, indistinguishable from neutral) ionization gas might lie either much further away or equally close but dense enough to remain neutral. \\item When N$_{H}$ is larger than about $10^{24}$ cm$^{-2}$, the column is Compton-thick, that is a diffusive process adds to the absorption process to deplete the specific intensity along the line of sight. A truly thick absorber ($\\tau_{Th}\\gg$ 1) reduces the intensity by several orders of magnitude all the way throughout the hard X-ray band, and renders in fact its quantitative evaluation impossible, except for a ``lower limit\" (e.g. Matt et al. 1999). \\end {enumerate}  It must be emphasized that, in general, AGN spectra contain another component in addition to the PL, which becomes evident beyond 8-10 keV. It is due to reflection from Compton-thick, neutral gas, and, as expected on physical grounds, it is accompanied by a strong fluorescent iron line. From a pure observational point of view, this component is often the only visible one below $\\sim$ 10 keV when the PL is photoelectrically absorbed by N$_{H}$ of about $10^{24}$ cm$^{-2}$ (corresponding to a turnover in energy at about 10 keV), or greater. Its intensity is typically a few percent of the primary one in the 2-10 keV band. In such systems, it is regarded as very reasonable to assume that the absorbing and the reflecting gas belong to one and the same circumnuclear structure. Historically, it was the spectropolarimetric measurements in the optical band of NGC 1068 which led Antonucci \\& Miller (1985) to propose the so-called ``unified model\". This model (Antonucci 1993) envisages the existence of a thick and dusty ``torus\" of molecular gas around the nucleus, extending out to a few tens of pc at most (Risaliti et al. 1999) and with a substantial covering factor to account for the large ratio of type 2 to type 1 Seyferts, at least in the local Universe (4:1, if also Seyfert 1.8 and 1.9 are included: Maiolino \\& Rieke 1995; Ho et al. 1997; it is worth recalling that the local Universe is dominated by relatively low luminosity objects as compared to the Quasi Stellar Objects, QSO). In its basic form, the model has been widely confirmed, most notably by measurements in the hard X-ray band, which showed that most type 2 Seyferts are X-ray absorbed. The situation, however, is likely to be more complex than the zero-th order model would predict. On one hand, there is increasing evidence of a substantial population of optically inactive, obscured AGN, both in the local (Maiolino et al. 2003) and more distant (Comastri et al. 2002; Brandt \\& Hasinger 2005) Universe. This implies that the ratio bewteen X-ray obscured to unobscured AGN may be larger than that between type 2 and type 1 Seyferts. On the other hand, other absorbing regions,  besides the torus, may exist in the complex environment of AGN: dust lanes (Malkan et al. 1998), starburst regions (Weaver 2002), the galactic disk (Maiolino \\& Rieke 1995). Indeed, it has been suggested (see e.g. Matt 2000, 2004 and references therein) that Compton--thick absorption (which, in the local Universe, is observed in about half of optically selected type 2 Seyferts, Risaliti et al. 1999, Guainazzi et al. 2005, the percentage rising when ``elusive'' AGN are also taken into account, Maiolino et al. 2003) is related to the torus, while Compton--thin absorption is due to one (or more) of the other possible absorbing regions. (If a dust/gas ratio typical of the cold interstellar medium of our Galaxy is adopted, Compton--thin absorption, i.e. N$_{H}$ in the 10$^{22}$--10$^{24}$ cm$^{-2}$ range, is sufficient to explain the extinction of the BLR lines).  This paper will address the question of the location of the absorber in the light of new results from X-ray surveys in the 2-10 keV band. The sensitivity reached in this band with the XMM-Newton and Chandra satellites made it, for the first time, possible to study AGN of both type 1 and 2, with sufficient statistics to investigate the  N$_{H}$ distribution as a function of the intrinsic (PL) luminosity and of the cosmic epoch out to z about 4.  Surveys limited to 10 keV on the high energy side (and with instrument sensitivity dropping rapidly beyond 5-7 keV) are doomed to miss almost completely the Compton-thick objects. Thus the results must be read as pertinent, among the AGN classified through optical spectroscopy as type 2, only to those with  N$_{H}$ less than $10^{24}$ cm$^{-2}$. Basically, the question addressed in this paper is whether the new results can be understood in the context of the unified model based on just one element, the torus mentioned above, or, if not, if one can conceive of just one realistic (astrophysically) additional element able to provide a positive answer, with the least possible number of assumptions. In particular, the new contribution of this paper is the use of one further constraint, namely the decrease of the fraction of Compton-thin obscured AGN with increasing X-ray luminosity, an anticorrelation which have recently emerged from the statistical analysis of the abovementioned types of survey (Ueda et al. 2003; La Franca et al. 2005).  In Sect. 2 this new constraint is described. In Sect. 3 the current physical models of the torus are discussed and shown to be unable to account, even qualitatively, for the anticorrelation. In Sect. 4 it is shown that, if the obscuration is due to molecular gas in galactic disks, then the gravitational effects of the black hole on its spatial distribution within 25--450 pc (depending on the mass) leads naturally to an anticorrelation. An encouragingly reasonable quantitative match is achieved under two conditions: that the molecular mass surface density exceeds $\\sim$150 $M_{\\odot}$ pc$^{-2}$ within 25-450 pc, and that in a statistical sense the bolometric luminosity is proportional to the black hole mass, and equal to about 10\\% of the Eddington luminosity. In Sect. 5 these two conditions are discussed in the light of the evidence available, and a positive conclusion is drawn. ",
        "conclusions": "The main result of this paper is that the anticorrelation between  $\\xi$(L$_x$) (namely the fraction of AGN with N$_H$, as measured by their X-ray spectra, greater than 10$^{22}$ cm$^{-2}$) and L$_x$ in Compton-thin objects is at least qualitatively reproduced provided that: a) there is a sufficiently large number of molecular clouds within the radius where the BH gravitational influence is dominant; b) there is, in a statistical sense, a correlation between the BH mass and the AGN luminosity. These two conditions are now briefly discussed.  As shown in Fig.~\\ref{anticorr}, the slope and the normalization of the anticorrelation are consistent with the observational results for values of $\\Sigma \\geq$ 150-200 $M_{\\odot}$ pc$^{-2}$, if a bolometric luminosity 0.1 times the Eddington luminosity is assumed (as it seems to be the case in the local universe, e.g. Peterson et al. 2004). Thus, the first question arising is what are the typical values of this quantity in the innermost region of spiral galaxies (the hosts of Seyfert active nuclei). The relevant radius of this region, 2$R_{infl}$, lies between about 25 and 450 pc (for $M_{BH}$ from one and one thousand million solar masses), which corresponds to 0.5-9 arcsec at a distance of 10 Mpc, and requires the best angular resolution currently achievable, with interferometry at the 3 mm CO J=1-0 line, on relatively nearby galaxies. In this respect, the BIMA Survey of Nearby Galaxies (BIMA SONG) at Hat Creek, CA, with the 10-element Berkeley-Illinois-Maryland Association (BIMA) millimeter interferometer (Helfer et al. 2003) is the best available for our purpose, because the sample was deliberatly selected without reference to CO or infrared brightness. It consists of 44 spiral galaxies of all morphological types (the Sa type are somewhat undersampled relative to the later types). The angular resolution of 6 arcsec corresponds to 360 pc (that is 180 pc radius for the central region) at the average distance of the galaxies, 12 Mpc. From the CO surface brightness distribution (see Table 5 in Helfer et al. 2003) it turns out that the peak brightness is attained within the central 6 arcsec in 20 out of 44 galaxies. In these 20 objects the corresponding surface density $\\Sigma_{centr}$ lies between 15 and 1854 $M_{\\odot}$ pc$^{-2}$, with most of them (18) above 90, and the other two around 15 $M_{\\odot}$ pc$^{-2}$. Notably, among the other 24 galaxies, 10 have $\\Sigma_{centr}$ $\\geq$ 50 $M_{\\odot}$ pc$^{-2}$, with values up to 1151. The galaxies with $\\Sigma_{centr}$ $\\geq$ 150 $M_{\\odot}$ pc$^{-2}$ are 22/44, namely 50\\%. The straight mean value of  this 22 galaxies is 511 $M_{\\odot}$ pc$^{-2}$, notably close to the value of $\\sim$500 $M_{\\odot}$ pc$^{-2}$ of the 300 pc molecular disk in our own galaxy (G\\\"usten 1989).  The sample contains 8 objects classified as Seyfert galaxies, 6 of them with $\\Sigma_{centr}$ greater than 90 and up to 466 $M_{\\odot}$ pc$^{-2}$, 1 (NGC 4725) with $\\Sigma_{centr}$ = 21 $M_{\\odot}$ pc$^{-2}$, 1 (NGC3031) with $\\Sigma_{centr}$ below the detection limit: the proportion of galaxies with an active nucleus, which possess a high value of $\\Sigma_{centr}$, appears to be, within the small statistics, similar to that in the general population.    The fraction of objects which meet our requirements is encouraging, especially because it does not require a special correlation with the nuclear activity, and represents a relatively large scale (compared to the close circumnuclear environment) property fairly common in spiral galaxies. In other words, the condition required is already in place in about half of the galaxies before a high rate nuclear activity lights up, which in turn is unlikely to significantly influence the preexisting condition. Much higher resolution and sensitivity surveys are needed to explore more distant objects, and therefore draw more solid conclusions on the validity of our hypothesis: to pursue this goal an array of millimeter telescopes like ALMA is necessary.   Concerning the increase of the fraction of absorbed sources with redshift (La Franca et al. 2005), an interesting possibility is that it might be due to an average increase of $\\dot{M}$ with the redshift, if one were allowed to assume that the molecular content and distribution in spiral galaxies is independent of $z$. Indeed, McLure \\& Dunlop (2004) found that the $L_{bol}/L_{Edd}$ ratio raises from about 0.15 at $z\\sim$0.2 to 0.5 at  $z\\sim$2. We intend to investigate this issue further, in particular within the frame of astrophysically self consistent models, which associate the growth of the supermassive BH in galaxies with the evolution in cosmic time of the AGN population (e.g. Menci et al. 2004)."
    },
    "0512/astro-ph0512327_arXiv.txt": {
        "abstract": "The physical nature of the currently observed dark energy in the universe is completely unclear, and many different theoretical models co-exist. Nevertheless, if dark energy is produced by vacuum fluctuations then there is a chance to probe some of its properties by simple laboratory tests based on Josephson junctions. These electronic devices can be used to perform `vacuum fluctuation spectroscopy', by directly measuring a noise spectrum induced by vacuum fluctuations. One would expect to see a cutoff near 1.7 THz in the measured power spectrum, provided the new physics underlying dark energy couples to electric charge. The effect exploited by the Josephson junction is a subtile nonlinear mixing effect and has nothing to do with the Casimir effect or other effects based on van der Waals forces. A Josephson experiment of the suggested type will now be built, and we should know the result within the next 3 years. ",
        "introduction": "It would be nice to start this paper with a clear definition of what dark energy is and what it is good for. Unfortunately, the answer to this question is completely unclear at the moment. What is clear is that various astronomical observations \\cite{bennett03, spergel03} (supernovae, CMB fluctuations, large-scale structure) provide rather convincing evidence that around 73\\% of the energy contents of the universe is a rather homogeneous form of energy, so-called `dark energy'. It behaves similar to a cosmological constant and currently causes the universe to accelerate its expansion. Dark energy may just be vacuum energy (with an equation of state $w=p/\\rho=-1$, where $p$ denotes the pressure and $\\rho$ the energy density). In that case its energy density $\\rho$ is constant and does not change with the expansion of the universe. Or, $w$ may be just close to -1, in which case the dark energy density evolves dynamically and changes with the expansion of the universe. The remaining contents of the current universe is about 23\\% dark matter and 4\\% ordinary matter. With 96\\% of the universe being unknown stuff, there is enough room (and, indeed, the need) for new theories. It seems that in order to construct a convincing theory of dark energy that explains why it is there and what role it plays in the universe one has to be open-minded to new physics.  A large number of different theoretical models exist for dark energy, but an entirely convincing theoretical breakthrough has not yet been achieved. Popular models are based on quintessence fields, phantom fields, quintom fields, Born-Infeld quantum condensates, the Chaplygin gas, fields with nonstandard kinetic terms, to name just a few (see e.g.\\ \\cite{review, weinberg, caroll, pad} for reviews). All of these approaches contain `new physics' in one way or another, though at different levels. However, it is clear that the number of possible dark energy models that are based on new physics is infinite, and in that sense many other models can be considered as well. Only experiments will ultimately be able to confirm or refute the various theoretical constructs. ",
        "conclusions": ""
    },
    "0512/astro-ph0512057_arXiv.txt": {
        "abstract": "We present near ultraviolet imaging with the Hubble Space Telescope Advanced Camera for Surveys, targeting young radio galaxies (Gigahertz Peaked Spectrum and Compact Steep Spectrum sources), in search of star formation regions in their hosts. We find near UV light which could be the product of recent star formation in eight of the nine observed sources. However, observations at other wavelengths and colors are needed to definitively establish the nature of the observed UV light. In the CSS sources 1443+77 and 1814--637 the near UV light is aligned with and is co-spatial with the radio source, and we suggest that in these sources the UV light is produced by star formation triggered and/or enhanced by the radio source. ",
        "introduction": "The relationship between black hole mass and galaxy mass implies that the growth and evolution of black holes (therefore AGN) and their host galaxies must somehow be related \\citep[e.g.,][]{Gebhardt00}. Mergers and strong interactions can trigger AGN activity in a galaxy \\citep[e.g.,][]{Heckman86, Baum92, Israel98}. These events can also produce instabilities in the ISM and trigger star formation (e.g., Ho 2005). Numerical simulations \\citep[e.g.,][]{Mellema02,Rees89} suggest that the advancement of the jets through the host galaxy environment can also trigger star formation. UV studies of large 3CR sources find traces of episodes of star formation around the time when the radio source was triggered \\citep[i.e. $\\lesssim 10^7 - 10^8$~yr,][]{Koekemoer99, Allen02, O'Dea01, O'Dea03, Martel02} suggesting a possible link between both. %  Gigahertz Peaked Spectrum (GPS) sources and Compact Steep Spectrum (CSS) are young, smaller \\citep[GPS $\\lesssim 1$ kpc, CSS $\\lesssim$15 kpc, for a review see][]{O'Dea98} versions of the large powerful radio sources, so they are expected to exhibit signs of more recent star formation. Their size makes them excellent probes of the interactions between the expanding lobes and the host. They have not yet completely broken through the host ISM, so these interactions are expected to be even more important than in the larger sources.  The near UV observations are very sensitive to the presence of hot young stars and therefore will trace recent star formation events. We have obtained high resolution HST/ACS UV images of these young compact sources to study their morphology and the extent of recent star formation. %  Our sample is chosen to be representative of GPS and CSS sources with z$\\lesssim$ 0.5, nearby enough to eliminate strong effects due to evolution with cosmic time. The objects are drawn primarily from the well-defined samples of \\citet{Fanti90, Fanti01} and \\citet{Stanghellini97}. \\\\  % ",
        "conclusions": "We have obtained HST/ACS near-UV high resolution images of young radio sources: GPS and CSS galaxies. We detect near UV emission (point sources and/or  clumps) in eight of the sources, consistent with the presence of recent star formation. In two CSS sources, 1443+77 and 1814--637 the near UV emission is aligned with and co-spatial with the the radio emission and we suggest that star formation has been triggered/enhanced by  expansion of the radio source through the host. Observations at other wavelengths and measurement of the colors are needed to further asses the nature of the observed UV properties.            \\newcommand\\aj{AJ} \\newcommand\\apj{ApJ} \\newcommand\\apjl{ApJ} \\newcommand\\apjs{ApJS} \\newcommand\\mnras{MNRAS} \\newcommand\\aaps{A\\&AS} \\newcommand\\aap{A\\&A} \\newcommand\\nat{Nature} \\newcommand\\aapr{A\\&A~Rev.} \\newcommand\\pasp{PASP}"
    },
    "0512/astro-ph0512261_arXiv.txt": {
        "abstract": "Curves in a family derived from powers of the polar coordinate formula for ellipses are found to provide good fits to bound orbits in a range of power-law potentials. This range includes the well-known $1/r$ (Keplerian) and logarithmic potentials. These approximate orbits, called p-ellipses, retain some of the basic geometric properties of ellipses. They satisfy and generalize Newton's apsidal precession formula, which is one of the reasons for their surprising accuracy. Because of their simplicity the p-ellipses make very useful tools for studying trends among power-law potentials, and especially the occurence of closed orbits. The occurence of closed or nearly closed orbits in different potentials highlights the possibility of period resonances between precessing, eccentric orbits and circular orbits, or between the precession period of multi-lobed closed orbits and satellite periods. These orbits and their resonances promise to help illuminate a number of problems in galaxy and accretion disk dynamics. ",
        "introduction": "In recent decades, the recognition that galaxies are surrounded by extended halos of dark matter has spurred a greatly increased interest in the structure of orbits in gravitational potentials that are shallower than the Keplerian $1/r$ potential (e.g., \\citealt{bin87}, \\citealt{ber00}). Similarly, the mass distribution in galactic bulges and elliptical galaxies is also extended, and the potential shallower than the Keplerian one. The logarithmic potential, $\\Phi \\propto \\ln(r)$, is among the simplest and most studied examples of such a shallow potential, and will be used as a primary example in this work. The celestial mechanics of such potentials is not nearly as well developed as the centuries old literature of the classical Kepler-Newton potential, but there has been much recent progress (see \\citealt{boc96} and \\citealt{con02}).  The orbits in most physically relevant potentials can be numerically integrated very quickly and accurately, except near singularities and resonances. Direct numerical integration has provided much information about orbit families in both general symmetric, and non-axisymmetric potentials relevant to barred galaxies (including non-axisymmetric logarithmic potentials, see e.g., \\citealt{mir89}).  There has also been much progress in the areas of analytic approximation and global properties of orbits in non-Keplerian potentials, and power-law potentials in particular. Concerning global properties, \\citet{sto03} have studied the topology of energy surfaces of the logarithmic potential and derived a complete qualitative description of orbits in the axisymmetric case. They also proved that orbits are centrophobic for the non-axisymmetric case as well. \\citet{val05} have presented extensions of Newton's apsidal precession theorem \\citep{new87} for power-law potentials, and in the process found no evidence for cases of zero precession at high eccentricities. The work below provides further support for that conjecture.  With regard to analytic approximation, we note first of all that classical epicyclic approximations are the most well developed (e.g., see \\citealt{bin87}, \\citealt{deh99}, and \\citealt{con02}). Beyond this, \\citet{con90} have used the multiple Fourier series method of \\citet{pre82} to find orbit families and to derive accurate orbit approximations for the symmetric and non-axisymmetric logarithmic potentials. \\citet{ada05} have examined spirograph or epicycloid orbit approximations and found that they accurately fit a wide range of orbits in the Hernquist and other halo potentials (with only two ``spirograph'' circles). The spirograph or epicycloid approximation is formally distinct from the epicyclic approximation, though similar in that they both can be extended to arbitrary accuracy.  More recently, \\citet{tou97} developed a relatively simple symplectic mapping for non-axisymmetric potentials that captures the behavior of orbits in these potentials. The map is based on the approximation that orbital evolution is driven by two processes: apsidal precession in the axisymmetric part of the potential, and torques by the non-axisymmetric potential, concentrated at apoapse. Because this symplectic map is based on apsidal precession, the Valluri et al. paper and the results below on precession complement it.  With this recent work the knowledge base on orbits in galaxy potentials (and similarly shallow potentials in galaxy groups and clusters) has become quite extensive. However, the basis of much of classical celestial mechanics (and many specific approximation schemes) is the simplicity of Kepler's elliptical orbit solution. The conceptual and analytical foundation of galactic dynamics and other fields that use power-law potentials is inevitably weaker without such simple orbital solutions. Unfortunately, mechanics textbooks teach us that analytic orbital solutions only exist in a few special cases among the power-law forces (\\citealt{gol80}, \\citealt{dan88}). The simple epicyclic approximation provides a versatile computational tool, but with limitations as a conceptual tool if multiple or large epicycles are needed for accurate calculations (e.g., \\citealt{con02}, section 3.1). The references above show that other accurate approximations have been developed, but they are complex.  The purpose of this paper is to present a family analytic orbit approximations for the logarithmic and other spherically symmetric, or axisymmetric potentials. These orbit functions are simply powers of precessing ellipses, and so, very simple indeed. Nonetheless, for small-to-moderate eccentricities, they are also surprisingly accurate. For the logarithmic potential, the approximate orbits librate around accurate numerical solutions, but do not diverge from them over time. Equally surprising, these approximate solutions extend continuously across all members of a family of potentials that range from the Keplerian potential to potentials with slightly rising rotation curves. Although they do not provide complete analytic solutions, their simplicity and accuracy allow them to fill a large part of the role of simple analytic solutions in providing a readily understandable conceptual picture of the nature of orbits in these potentials and the relations between them. ",
        "conclusions": "In the sections above, we have shown that curves derived as powers of ellipse functions, called p-ellipses, provide simple, yet surprisingly accurate approximations to orbits in a range of power-law potentials, including the well-known logarithmic potential, and softened power-law potentials. For a given power-law potential the p-ellipse function is nearly as simple as an ellipse, but the family of p-ellipses extends continuously across a physically interesting range of potentials.  There are at least two reasons why p-ellipses are good orbital approximates across this range of potentials, and are likely to be the best simple, analytic functions to do so. First, p-ellipses match the tendency for orbits of a given energy and angular momentum to become more circular as the potential changes from Keplerian to solid-body (and $\\delta$ decreases). In this sense the p-ellipses adjust well to the appropriate orbital shape in a given potential. Secondly, the precession rate obtained by demanding that the p-ellipse satisfy the equation of motion in a given potential to first order in eccentricity is the identical to that given by Newton's theorem. Thus, the p-ellipses precess correctly. Moreover, a second order approximation yields an eccentricity dependence of the precession rate that is in qualitative agreement with the \\citet{val05} semi-analytic results.  A number of the results obtained above for p-ellipses, like the apsidal precession rates in power-law potentials, are not new, and good approximations to individual orbits can be obtained with epicycles or numerical integration. However, a family of simple curves like the p-ellipses allow us to readily see trends across a range of potentials, so they provide a simple, conceptual picture for orbital variations, as discussed in the introduction.  Moreover, as shown in the previous sections the p-ellipses provide a very powerful tool for studying characteristics like the occurrence of (non-circular) closed orbits as a function of eccentricity in different potentials. The key feature of the p-ellipses in this regard is a relatively simple approximate formulae for precession rates as a function of $\\delta$ and $e$. Newton's theorem is valid in the limit of small eccentricity, and Valluri et al.'s extensions involve integrals that must be evaluated numerically. Equations (19) and (20), though approximate, offer convenience. (Also see eq. (25).)  The example of the softened power-law potentials holds out the hope that the p-ellipses can provide useful orbit approximations in other non-power-law potentials. Non-axisymmetric potentials have not been considered in this paper, but it seems reasonable to hope that p-ellipses could provide good approximations to loop orbits in such potentials. This issue deserves more study.  In Section 5.3 I described several examples of how the systematics of p-ellipses might shed light on important astrophysical problems. This conclusion will be even more general if p-ellipses prove to be good orbital approximations in more types of potential.  Orbit theory is more general than celestial mechanics, and the p-ellipse approximations should also be relevant to any field than involves orbits in general potentials. Electron orbits in general, steady (gradient) electric and magnetic fields are obvious examples.  In sum, there seems to be room for a great deal more development of the theory and application of these simple curves. Their simplicity may allow us to address a number of complex issues that would be hard to study directly through the accumulation of numerical examples."
    },
    "0512/astro-ph0512526_arXiv.txt": {
        "abstract": "We present photometric data of the type II-P supernova (SN) 2004dj in NGC 2403. The multicolor light curves cover the SN from $\\sim$ 60 to 200 days after explosion, and are measured with a set of intermediate-band filters that have the advantage of tracing the strength variations of some spectral features. The light curves show a flat evolution in the middle of the plateau phase, then decline exponentially at the late times, with a rate of 0.10$\\pm$0.03 mag (10 days)$^{-1}$ in most of the filters. In the nebular phase, the spectral energy distribution (SED) of SN 2004dj shows a steady increase in the flux near 6600~\\AA\\, and 8500~\\AA, which may correspond to the emission lines of  H$\\alpha$ and Ca~II near-IR triplet, respectively. The photometric behavior suggests that SN 2004dj is a normal SN II-P. Compared with the light curves of another typical SN II-P 1999em, we estimate the explosion date to be June 10$\\pm$21 UT, 2004 (JD 2453167$\\pm$21) for SN 2004dj. We also estimate the ejected nickel mass during the explosion to be $M(^{56}\\rm{Ni})$ = 0.023 $\\pm$ 0.005 $M_{\\odot}$ from two different methods, which is typical for a SN II-P. We derive the explosion energy $E \\approx 0.75^{+0.56}_{-0.38}\\times10^{51}$ erg, the ejecta mass $M \\approx 10.0^{+7.4}_{-5.2}$ $M_{\\odot}$, and the initial radius $R \\approx 282^{+253}_{-122}$ $R_{\\odot}$ for the presupernova star of SN 2004dj, which are consistent with other typical SNe II-P. ",
        "introduction": "Supernova (SN) 2004dj was discovered by K. Itagaki \\citep{nak04} on 2004 July 31.76 (UT dates are used throughout this paper) in the nearby SBcd galaxy NGC 2403 at a distance of 3.3$\\pm$0.1 Mpc \\citep{kar04}. With a peak magnitude of 11.2 mag in the $V$ band, SN 2004dj was the brightest supernova in the past decade. Spectroscopic observations of SN 2004dj show that it is a typical Type II Plateau SN (SN II-P), with prominent P-Cygni profiles in hydrogen Balmer lines \\citep{pat04}. The main photometric characteristic of a SN II-P is that its light curve, unlike other SNe, does not decay rapidly after maximum, but shows a plateau phase for 60-100 days before decaying exponentially. The plateau phase originates from a balance between the receding photosphere in the expanding ejecta when the supernova is powered by the recombination of hydrogen previously ionized by the supernova shock. SNe II-P have long been thought to be produced by core-collapse of massive ($>$8 $M_{\\odot}$) red supergiants that do not experience significant mass loss and retain  most of their hydrogen-rich envelopes.  The optical position of SN 2004dj is measured to be R. A. = 07$^{\\rm{h}}$37$^{\\rm{m}}$17$^{\\rm{s} }$.04, Dec = +65$^{\\circ}$35$^{\\prime}$57$^{\\prime\\prime}$.84 (J2000.0; \\citet{nak04}). This is in good agreement with the {\\it Chandra} X-ray position \\citep{poo04} and the  MERLIN radio position \\citep{bes05}. Using the astrometry of SN 2004dj, and also from geometrical registration between images of NGC 2403 before and after SN 2004dj occurred, both \\citet{mai04} and \\citet{wxf05} have convincingly shown that SN 2004dj occurred at a position coincident with Sandage Star 96 (hereafter S96) in the list of luminous stars and clusters in NGC 2403 published by \\citet{san84}. \\citet{mai04} suggest that S96 is a young compact cluster with an age of 13.6 Myr and a total stellar mass of $\\sim$ 24,000 $M_\\odot$, and estimate that the progenitor of SN 2004dj had a main-sequence mass of about 15 $M_\\odot$. \\citet{wxf05}, on the other hand, suggest that S96 is an older ($\\sim$ 20.0 Myr) and more massive ($\\sim$ 96,000 $M_\\odot$) cluster, and that the progenitor of SN 2004dj may have a mass of $\\sim$ 12 $M_\\odot$.  Studies of SN 2004dj itself may shed light  on the nature of its progenitor, thus providing additional constraints on the properties of S96. In this paper, we present the results from our campaign of photometric followup of SN 2004dj from Aug 2004 to Jan 2005. The observations and data reduction are described in $\\S$ 2, and the multicolor light curves and spectral energy distribution (SED) are presented in $\\S$ 3. We estimate the explosion date, the reddening, the synthesized $^{56}$Ni mass, and some explosion parameters for SN 2004dj in $\\S$ 4. The conclusions are summarized in $\\S$ 5. ",
        "conclusions": "This paper presents the photometry of SN 2004dj taken with 12 intermediate-band filters, and obtained from $\\sim$ 60 to 200 days after the explosion. Our observations show that SN 2004dj was discovered in the middle of the plateau phase. The multicolor light curves show a plateau phase of about 45 days, a transition period of about 30 days during which the SN declines dramatically, and a nebular phase during which the SN declines at a rate of $\\sim$ 0.10$\\pm$0.03 mag (10 days)$^{-1}$ (consists with the decay rate of $^{56}$Co).  The SEDs for SN 2004dj are constructed from the measured flux in the 12 passbands. These SEDs show an evolution that is similar to the spectral evolution of a normal SN II-P. A flux peak near 6600~\\AA\\, is observed, which is consistent with the strong H$\\alpha$ emission seen in these objects. The flux around 8500~\\AA\\ evolves from a flux deficit to a flux peak, consisting with the evolution of the strong Ca~II near-IR triplet line.  By comparing the light curve of SN 2004dj to the well observed SN II-P 1999em, we estimate the explosion date for SN 2004dj to be JD 2453168$\\pm$21 (June 11, 2004), approximately 50 day before the actual discovery of SN 2004dj.  Finally, we estimate the $^{56}$Ni mass synthesized during the SN 2004dj explosion using two methods: the bolometric tail luminosity and the ``steepness\" parameter. These methods yield a value of $M_{\\rm{Ni}}$ = 0.023$\\pm0.005$ $M_\\odot$ for SN 2004dj, which is typical for a normal SN II-P. Comparing our observed parameters of SN 2004dj with the analytical models of SN II-P light curves \\citep{lina85}, We derive for SN 2004dj an explosion energy of $0.75^{+0.56}_{-0.38}\\times10^{51}$ erg, an ejecta mass of $10.0^{+7.4}_{-5.2} M_{\\odot}$, and an initial radius of $282^{+253}_{-122} R_{\\odot}$. These parameters suggest that the progenitor of SN 2004dj is likely to be a K0--M0 type supergiant."
    },
    "0512/astro-ph0512460_arXiv.txt": {
        "abstract": "We searched  for steady PeV gamma-ray emission from the Monogem ring region with the Tibet air shower array from 1997 February to 2004 October . No evidence for statistically significant gamma-ray signals was found in a region 111$\\degr$~$\\leq$ R.A. $<$~114$\\degr$,~12$\\fdg$5~$\\leq$ decl. $<$ 15$\\fdg$5  in the Monogem ring where the MAKET-ANI experiment recently claimed a positive detection of PeV high-energy cosmic radiation, although our flux sensitivity is approximately 10 times better than MAKET-ANI's. We set the most stringent integral flux upper limit at a 99$\\%$ confidence level of 4.0 $\\times$ 10$^{-12}$~cm$^{-2}$~s$^{-1}$~sr$^{-1}$ above 1 PeV on diffuse gamma rays extended in the 3$^{\\circ}$ $\\times$ 3$^{\\circ}$ region. ",
        "introduction": "In a recent observation at PeV energies, the MAKET-ANI air shower experiment at Mount Aragats (E44$\\degr$10$\\arcmin$, N40$\\degr$30$\\arcmin$N; 3200 m above sea level) claimed a detection of significant excess (6$\\sigma$) of cosmic-ray events within a 3$\\degr$~$\\times$~3$\\degr$ search window~(111$\\degr$~$\\leq$ $\\alpha$ $<$ 114$\\degr$, 12$\\fdg$5 $\\leq$ $\\delta$ $<$ $15\\fdg5$) in the Monogem ring region using the air shower data recorded from 1997 to 2003\\\\ \\cite{Chilingarian2003} Naturally, the significant excess may be attributed to PeV gamma rays, because the Larmor~radius $R_{\\rm L}$ $\\sim$ 0.4 $/$ $Z$~pc at 10$^{15}$~eV in the galactic magnetic field of 3 $\\mu$G is too small to reach the Earth without deflection compared with the distance of 300~pc between the Earth and the Monogem~ring, and the mean decay length of a neutron $\\lambda$ $\\sim$ 10~pc at 10$^{15}$~eV is also too short.  The Monogem~ring is a diffuse (extended with a diameter of 25$\\degr$ in the sky) supernova remnant (SNR) associated with the radio pulsar PSR~B0656+14 at a distance of approximately 300~pc \\cite{Thorsett2003}.  This SNR is a bright source in the soft X-ray region. According to a detailed observation of the Monogem~ring made by the $ROSAT$ X-ray survey \\cite{Plucinsky1996}, the average temperature of thermal emission is 6.15 in log($T$/1 K).  If the Monogem~ring is modeled as an SNR in the adiabatic stage of evolution, the initial ambient density is estimated to be 5.2~$\\times$~10$^{-3}$~cm$^{-3}$, the initial explosion energy is estimated to be 1.9~$\\times$~10$^{50}$~erg, the radius to the shock front is estimated to be 66.5~pc, and the age is estimated to be 86 thousand years. The Monogem ring has been considered to be a possible candidate for the particle acceleration site of cosmic electrons \\cite{Kobayashi2004} and nuclei \\cite{Thorsett2003} in our Galaxy.  The cosmic-ray flux we observe at the Earth roughly shows a power-law spectrum at energies from 10$^{9}$ to 10$^{20}$~eV.  There is a general consensus that the stochastic shock acceleration at SNRs in our Galaxy could explain the cosmic-ray energy spectrum up to approximately $10^{15}$~eV.  The maximum energy is estimated to be only $\\sim Z \\times 10^{14}$ eV by parallel shock acceleration where the normal axis of shock front is parallel to the direction of the interstellar magnetic field. On the other hand, the oblique shocks accelerate particles more efficiently than the parallel shocks \\cite{Jokipii1987}. The maximum energy based on the oblique shock acceleration increases by several orders of magnitude, and could explain the cosmic-ray energy spectrum up to approximately $10^{17}$~eV \\cite{Kobayakawa2002}.  In this case, the cosmic rays beyond $10^{17}$~eV are assumed to be of extragalactic sources, such as active galactic nuclei and gamma-ray burst .  Among recent models, the ``single source (SS) model'' \\citep{Berezhko1996,Erlykin} is interesting in that a single nearby SNR, such as the Monogem~ring, mainly contributes to the cosmic-ray intensity observed at the Earth.  To explain the shape and intensity of the cosmic-ray energy spectrum using the SS model, the most likely parameters of the single SNR should be 300--350 pc distant and 90--100 thousand years old \\cite{Erlykin2003}.  If such a strong cosmic-ray accelerator lies near the Earth, we may observe an anisotropy of cosmic rays from its direction, or we may detect high-energy photons that are emitted from these high-energy charged particles by the nonthermal processes.  In this paper, we report on the search for diffuse/pointlike PeV gamma-ray emission based on the data recorded from 1997 to 2004 around the Monogem~ring region by a large air shower array with a total area of 36,900~m${^2}$ constructed in Tibet. ",
        "conclusions": "No evidence for statistically significant gamma-ray signals was found in a region 111$\\degr~\\leq \\alpha<~114\\degr$,~12$\\fdg5~\\leq \\delta < 15\\fdg$5  in the Monogem Ring where the MAKET-ANI experiment recently claimed a positive detection of PeV high energy cosmic radiation, although our flux sensitivity is approximately 10 times better than MAKET-ANI's. We set the most stringent integral flux upper limit at a 99$\\%$ confidence level of 1.1 $\\times$ 10$^{-14}$ cm$^{-2}$s$^{-1}$/(2.66 $\\times$ 10$^{-3}$sr) =~4.0 $\\times$ 10$^{-12}$~cm$^{-2}$~s$^{-1}$~sr$^{-1}$ on steady diffuse gamma rays $>1$~PeV extended within the rectangular region in the Monogem Ring assuming a differential spectral index $-2.0$. One of the potential possibilities for explaining the discrepancy between the two experiments could be transient emission that occurred at occasions when we stopped data acquisition for annual maintenance, calibration, upgrading jobs, etc. Another could be the strong transient emission of PeV gamma rays that incidentally or periodically occurred during the 3 hr when only the MAKET-ANI experiment could observe because the two experimental sites are separated by about 45$\\degr$ in longitude. Third, the excess events that they detected might be due to statistical fluctuation \\citep{ChilingarianICRC}.      No significant signal was found in the whole Monogem Ring region based on a 0$\\fdg$5 $\\times$ 0$\\fdg$5 window search for a pointlike source at energies $>1$ PeV.  We also set 99$\\%$ confidence-level flux upper limits of 2.6 and 5.4 $\\times$ 10$^{-15}$ cm$^{-2}$ s$^{-1}$ on the steady gamma rays $>1$~PeV from PSR~B0656+14 and Geminga, respectively, assuming pointlike sources with a differential energy spectral index $-$2.0. The KASCADE group also reported that no significant signal was seen in the sub-PeV region at the suggested location by the MAKET-ANI experiment and PSR~B0656+14 by a pointlike source analysis \\cite{Antoni2004}. Furthermore, we reported on the result of wide sky survey for steady TeV gamma-ray pointlike sources elsewhere \\cite{all}, although no significant point source was found in the Monogem Ring region above a few TeV and above 10 TeV."
    },
    "0512/astro-ph0512183_arXiv.txt": {
        "abstract": "{New NTT/EMMI spectrophotometry of single WN2--5 stars in the Small and Large Magellanic Clouds are presented, from which He\\,{\\sc ii} $\\lambda$4686 line luminosities have been derived, and compared with observations of other Magellanic Cloud Wolf-Rayet stars from the literature. SMC WN3--4  stars possess line luminosities which are a factor of 4 times lower than LMC counterparts, incorporating several binary SMC WN3--4 stars from the literature. Similar results are found for WN5--6 stars, despite reduced statistics, incorporating observations of single LMC WN5--9 stars from the literature.  C\\,{\\sc iv} $\\lambda$5808 line luminosities of carbon sequence WR stars in the SMC and IC\\,1613 (both WO subtypes) from the recent literature are  a factor of 3 lower than LMC WC stars from Mt Stromlo/DBS spectrophotometry, although similar results are also obtained for the sole LMC WO star. We demonstrate how reduced  line luminosities at low metallicity follow  naturally if WR winds are metallicity-dependent, as recent empirical and theoretical results suggest.  We apply mass loss-metallicity scalings to atmospheric non-LTE models of Milky Way  and LMC WR stars to predict the wind signatures of WR stars in the metal-poor star forming WR galaxy I\\,Zw~18. WN He\\,{\\sc ii} $\\lambda$4686 line luminosities are 7--20 times lower than in metal-rich counterparts of identical bolometric luminosity, whilst WC C\\,{\\sc iv} $\\lambda5808$ line luminosities are 3--6 times lower. Significant He$^{+}$ Lyman continuum fluxes are predicted for metal-poor early-type WR stars. Consequently, our results suggest the need for larger population of WR stars  in I\\,Zw~18 than is presently assumed, particularly for WN stars,  potentially posing a severe challenge to evolutionary models at very low metallicity. Finally, reduced wind strengths from WR stars at low metallicities impacts upon the immediate circumstellar  environment of long duration GRB afterglows, particularly since the host galaxies of high-redshift GRBs tend to be metal-poor. ",
        "introduction": "Wolf-Rayet stars -- subdivided into nitrogen (WN) and carbon (WC) sequences -- are the evolved descendants of the O stars, and possess strong, broad, emission lines due to their dense, stellar winds. WR galaxies represent a subset of emission-line galaxies with active massive star formation via the direct signature of WR stars.  To date, over a hundred WR galaxies are known in the nearby universe (Schaerer et al. 1999) via a broad C\\,{\\sc iii} $\\lambda$4650/He\\,{\\sc ii} $\\lambda$4686 (blue bump) and/or C\\,{\\sc iv} $\\lambda$5808     (yellow bump) emission features seen in the integrated optical spectrum of individual sources. Indeed, broad He\\,{\\sc ii} $\\lambda$1640 emission, attributed to WR stars, can be easily seen in the average rest-frame spectrum of $z\\sim 3$ Lyman Break Galaxies (LBGs, Shapley et al. 2003).  The O star content of WR galaxies is typically derived indirectly using nebular hydrogen emission line fluxes to determine the total number of ionizing photons, from which the equivalent number of O7V stars is commonly calculated. The actual O star content depends primarily upon the age and mass function, and relates to the equivalent number of O7V stars via a correction factor (Vacca 1994). As such, this quantity is fairly metallicity {\\it independent}, although O7V stars at lower metallicity will likely have higher temperatures -- and so higher Lyman continuum fluxes -- than their high metallicity counterparts (Massey et al. 2005; Mokiem et  al. 2006).  In contrast, the WR content is routinely obtained merely by dividing the observed emission bump fluxes, corrected for reddening and distance, by average line luminosities of Milky Way and Large Magellanic Cloud WR stars (Schaerer \\& Vacca 1998). As such, line luminosities of WR stars are {\\it assumed} to be metallicity independent.  Following this technique, the WR content of metal-rich (e.g. Mrk 309: Schaerer et al. 2000)  and metal-poor (e.g. I\\,Zw~18: Izotov et al. 1997) galaxies have been derived and compared, generally successfully, with evolutionary synthesis models.  Recent observational and theoretical evidence suggests WR winds depend upon the heavy metal content of the parent galaxy (Crowther et al. 2002; Gr\\\"{a}fener \\& Hamann 2005; Vink \\& de Koter 2005). Since WR stars are believed to be the immediate precursors of some long-soft Gamma Ray Bursts (GRBs) the circumstellar environment of low metallicity GRBs is expected to differ substantially from those in metal-rich regions (Eldridge et al. 2005). If WR winds depend upon metallicity, do their line luminosities?  In the present study, we present new optical spectrophotometry of single, early-type WN stars in the Large and Small Magellanic Clouds in Sect~\\ref{obs}. The Magellanic Clouds were selected on the basis of their known distances, resolved stellar content and low interstellar reddenings. Indeed, reduced He\\,{\\sc ii} $\\lambda$4686 line luminosities for SMC WN stars ($\\log$~O/H$+12\\sim$8.1) relative to the LMC ($\\log$~O/H$+12\\sim$8.4) are obtained in Sect.~\\ref{line}, incorporating observations of late-type WN stars from the recent literature.  We also demonstrate that carbon sequence WR stars at low metallicity -- dominated by the rare WO subclass -- also possess lower line luminosities than typical LMC counterparts based upon spectrophotometry of single and binary WC stars, published in part by Crowther et al. (2002). In Sect~\\ref{models}, we consider whether reduced WR line luminosities are expected at lower metallicities for metallicity dependent WR winds. The results of our study are discussed in Sect.~\\ref{discussion}, with particular application to I\\,Zw~18, and brief conclusions are drawn in Sect~\\ref{summary}. ",
        "conclusions": "\\label{discussion}  We have demonstrated that line luminosities of early--mid WN stars in the LMC exceed those in the SMC by $\\sim$4--5. If this is supported for other metal-poor galaxies, one would need to apply such a corrective factor when determining WR populations at low metallicity. Unfortunately, within the Local Group only NGC~6822 and IC~10 --  with log (O/H)+12 $\\sim$ 8.25 (Pagel et  al. 1980; Garnett 1990) -- possess metallicities below that of the LMC. To date, only 4 and 11 WN stars have been confirmed in these respective galaxies (Massey \\& Johnson 1998; Crowther et al. 2003) whilst spectrophotometry is not yet available, preventing robust line flux comparisons since reliable reddening corrections rely upon accurate colours of isolated WR stars.   We propose that a correction factor may need to be applied when estimating WR populations from observations of unresolved clusters/galaxies at sub-LMC metallicities, where {\\it adopted} values from metal-rich WR calibrations may greatly exceed those of individual stars within those galaxies. In the extreme case of I\\,Zw~18, our test calculations suggest factors of $\\sim$5--20 may be appropriate.     \\subsection{WR population of I\\,Zw~18 from optical studies}  Izotov et al. (1997) and Legrand et al. (1997) presented observations of I\\,Zw~18-NW in which broad blue (C\\,{\\sc iii} $\\lambda$4650/C\\,{\\sc iv} $\\lambda$4658/He\\,{\\sc ii} $\\lambda$4686, FWHM$\\sim$70\\AA) and yellow (C\\,{\\sc iv} $\\lambda$5808, FWHM$\\sim$50\\AA) emission features were observed, together with nebular He\\,{\\sc ii} $\\lambda$4686 emission.  Izotov et al. derived blue and yellow line luminosities of 4.8$\\times 10^{37}$ erg/s and 2.1$\\times 10^{37}$ erg/s, respectively, for a revised distance of $\\sim$14.1\\,Mpc (Izotov \\& Thuan 2004). Applying  our own LMC WC calibration or that from Smith et al. (1990a) to the yellow feature would require $\\sim$7 equivalent WC4 stars ~(De  Mello et al. 1998). In contrast, with a typical low metallicity WC star  contributing a factor of $\\sim$5 less C\\,{\\sc iv} $\\lambda$5808  flux, we suggest a  much larger WC population of $\\geq$30 may be necessary to explain the observed line flux in I\\,Zw~18.  Depending upon individual temperatures, these stars would display either a weak-lined early WC, or a WO spectrum.   Legrand et al. (1997) noted that the observed line width of the C\\,{\\sc iv} feature more closely matches that of LMC WC4 stars than WO stars (recall Fig.~\\ref{wc}). However, WO stars are known to display decreasing wind velocities at lower metallicity, as demonstrated in Fig.~\\ref{wo}. Consequently, one might expect low metallicity WO stars to have unusually narrow lines with respect to their metal-rich counterparts.  Izotov et al. noted that the C\\,{\\sc iii-iv} $\\lambda$4650 flux far exceeded that expected from the number of WC stars inferred from C\\,{\\sc iv} $\\lambda$5808, assuming they were typical of LMC WC4 stars, which they exclusively attributed to  He\\,{\\sc ii} $\\lambda$4686 in late-type WN stars.  If WC stars in I\\,Zw~18 mimic those of the LMC one would expect C\\,{\\sc iii-iv} $\\lambda$4650/C\\,{\\sc iv} $\\lambda$5808 =1.5--1.6 (Table~\\ref{average_wc}; Smith et al. 1990a) so these would provide $\\sim$2/3 of the C\\,{\\sc iii-iv} $\\lambda$4650 flux observed by Izotov et al. (1997). The remainder could be attributed to $\\sim$10 WN5--6 stars, or  $\\sim$20 WN7--9 (or WN2--4) LMC-like stars (Table~\\ref{average_wn}; see  also De Mello et al. 1998).  We have shown that low metallicity WC stars are likely to possess rather different ratios of C\\,{\\sc iii} $\\lambda$4650 to C\\,{\\sc iv} $\\lambda$5808 line fluxes from near-Solar counterparts, so we {\\bf advise} caution when indirectly inferring WR populations at extremely low metallicities in this way.  Observationally, it is challenging to establish the presence of WN stars in the Izotov et al.  (1997) dataset since the $\\lambda$4650 feature is much broader than in late-type WN stars, with FWHM $\\sim$70\\AA. Late-type LMC metallicity WN stars possess FWHM$\\sim$20\\AA, at comparable spectral resolution (Fig.~\\ref{wn}).  The difficulty in  spectroscopically identifying WN stars in I\\,Zw~18 is almost certainly due to the extremely low He\\,{\\sc ii} $\\lambda$4686 line luminosity of individual WN stars, owing to the steeper metallicity dependence of their winds relative to WC stars (Vink \\& de Koter 2005). Since late-type WN stars positively shy away from low metallicity environments (recall Fig.~\\ref{wnl_model}), line luminosities of individual mid-type or early-type WN stars in I\\,Zw~18 may be a factor of $\\sim$10 times smaller than for  LMC late-type WN stars according to Schaerer \\& Vacca (1998). Indeed, if we assume $\\sim$1/3 of the C\\,{\\sc iii-iv} $\\lambda$4650 flux observed by Izotov et al. (1997) is due  to SMC early-type WN stars, we would require 100--300 WN stars, depending upon whether we adopt average values from Table~\\ref{average_wn} including or excluding the WN+O binaries. Consequently, large numbers of WN stars would be required for their  spectroscopic detection  via broad He\\,{\\sc ii} $\\lambda$1640 or $\\lambda$4686 emission. Alternatively, their presence  in smaller numbers may be seen  indirectly via strong nebular  He\\,{\\sc ii} $\\lambda$4686 emission,  which {\\it is} indeed observed in I\\,Zw~18.  \\begin{figure}[h] \\centerline{\\psfig{figure=4298f12.ps,width=8.8cm,angle=0.}} \\caption{Comparison between wind velocities of all known Local Group WO stars, from Kingsburgh et al. (1995), Kingsburgh \\& Barlow (1995) and Drew et al. (2004). For Milky Way WO stars we have adopted $\\log$(O/H)+12=8.6 for WR30a which lies beyond the Solar circle, 8.7 for Sand~5 which lies at the Solar circle and 8.8 for Sand~4 and WR93b which lie within the Solar circle.} \\label{wo} \\end{figure}  \\subsection{WR population of I\\,Zw~18 from UV studies}  Brown et al. (2002) scanned the NW part of I\\,Zw~18 with HST/STIS, revealing two  clusters, exhibiting strong C\\,{\\sc iv} $\\lambda$1550 and He\\,{\\sc ii} $\\lambda$1640 emission. Adjusting their fluxes to a distance of 14.1\\,Mpc (Izotov \\& Thuan 2004) indicates He\\,{\\sc ii} $\\lambda$1640 luminosities of L$_{\\rm 1640}$ = 3.0 $\\times 10^{37}$ erg/s and 4.0 $\\times 10^{37}$ erg/s, respectively. Brown et al. applied the Schaerer \\& Vacca (1998) He\\,{\\sc ii} $\\lambda$1640 calibration of a representative {\\it Milky Way} WC5 star,  implying 6 and 8 WC stars, respectively, with   N(WC)/N(O)$\\sim$0.2 in the former cluster.  A  similar exercise for the (distance adjusted)  C\\,{\\sc iv} $\\lambda$1550 luminosities of L$_{\\rm 1550}$ = 6.9 $\\times 10^{37}$ erg/s and 2.8 $\\times 10^{37}$ erg/s would lead to 3--4 and 1--2 stars, respectively, according to Sect.~\\ref{1548} assuming representative LMC-type WC4 stars. The reduced numbers with respect to Brown et al. (2002) is due to the higher line luminosities of LMC WC4 stars relative to Milky Way WC5 stars, plus the low $\\lambda$1550/$\\lambda$1640 ratio of $\\sim$0.7 for the second cluster, suggesting a primary contribution by WN stars rather than WC stars.  If we instead assume a representative C\\,{\\sc iv} $\\lambda$1550 line luminosity of 4$\\times 10^{36}$ erg/s for a single WC star in I\\,Zw~18, i.e. 5 times lower than typical LMC WC4 stars, we would require $\\sim$17 and $\\sim$7 weak-lined WC stars for the two clusters. Brown et al. remarked upon their unusually high N(WR)/N(O) populations, which would be exacerbated if larger WC populations are inferred at reduced metallicities.  \\subsection{Challenges for evolutionary models of single massive stars at low metallicity}   The potential presence of much greater numbers of WR stars in I\\,Zw~18 than is currently appreciated naturally causes problems for evolutionary models of  single massive stars at low metallicity.  Non-rotating, high mass-loss evolutionary models were calculated by de Mello et al. (1998), revealing progression through to the WN and WC phases for stars of initial mass $\\sim 90M_{\\odot}$ and 120$M_{\\odot}$, respectively. For an instantaneous burst with a Salpeter IMF and an upper mass limit of $\\sim150M_{\\odot}$ the maximum N(WR)/N(O) ratio is $\\sim$0.02, with WN stars dominating the WR population, i.e. N(WC)/N(O)$\\leq$0.003. In the case of a WR population with $\\sim$30 WC or WO stars, and potentially $\\sim$200 WN stars, one would obtain  N(WC)/N(O)$\\sim$0.02   and N(WN)/N(O)$\\sim$0.1  in I\\,Zw~18, based  upon the $\\sim$2000 O star content from Izotov et al. (1997, again  adjusted to a distance of  14.1\\,Mpc) greatly exceeding evolutionary  predictions for single stars. Further comparison awaits the calculation of evolutionary models for single stars at very low metallicity including rotation and contemporary mass-loss rates.  \\subsection{Circumstellar environment of long duration GRBs}  Within the past decade, several long duration ($\\geq$ 2s) GRBs have been positively  identified with Type Ic core-collapse SN (Galama et al. 1998; Stanek et al. 2003), supporting the collapsar model of MacFadyen \\& Woosley (1999) involving Wolf-Rayet stars. The ejecta strongly interacts with the circumstellar material, probing the immediate vicinity of the GRB itself, thus  providing information on the progenitor (Li \\& Chevalier 2003, van Marle et al. 2005).  If WR stars possess metallicity-dependent winds, one would potentially expect rather different environments for the afterglows of long-duration GRBs, that were dependent  upon the  metallicity of the host galaxy. WN stars in a  galaxy of 1/100$Z_{\\odot}$ may possess wind densities a factor of $\\sim$25 times lower than those in the Milky Way. In general, the metallicity dependence of wind velocities for WR stars is unclear, although amongst carbon sequence WR star, lower velocity winds are seen in WO stars from metal-poor environments (Fig.~\\ref{wo}).  Overall, the immediate environment of GRBs that involve Wolf-Rayet precursors may differ substantially from those of Solar metallicity WR stars (Eldridge et al. 2005). Indeed, the host galaxies of high-redshift GRBs tend to be rather metal-poor, from medium to high resolution spectroscopy obtained immediately  after the burst.  For example, Vreewwijk et al. (2004)  suggest 1/20$Z_{\\odot}$ for the host galaxy of GRB 030323 at $z$=3.37 and Chen et al. (2005) conclude 1/100$Z_{\\odot}$ for the host galaxy of GRB 050730 at $z$=3.97."
    },
    "0512/astro-ph0512306_arXiv.txt": {
        "abstract": "We observed the HH\\,211 jet in the submillimeter continuum and the CO(3--2) and SiO(8--7) transitions with the Submillimeter Array.  The continuum source detected at the center of the outflow shows an elongated morphology, perpendicular to the direction of the outflow axis. The high-velocity emission of both molecules shows a knotty and highly collimated structure.  The SiO(8--7) emission at the base of the outflow, close to the driving source, spans a wide range of velocities, from $-20$ up to 40 \\kms. This suggests that a wide-angle wind may be the driving mechanism of the HH\\,211 outflow.  For distances $\\ge 5''$ ($\\sim 1500$ AU) from the driving source, emission from both transitions  follows a Hubble-law behavior, with SiO(8--7) reaching higher velocities than CO(3--2), and being located upstream of the CO(3--2) knots. This indicates that the SiO(8--7) emission is likely tracing entrained gas very close to the primary jet, while the CO(3--2) is tracing less dense entrained gas.  From the SiO(5--4) data of Hirano et al.\\ we find that the SiO(8--7)/SiO(5--4) brightness temperature ratio along the jet decreases for knots far from the driving  source. This is consistent with the density decreasing along the jet, from  (3--10)$\\times 10^6$ \\cmt\\ at  500 AU to (0.8--4)$\\times 10^6$ \\cmt\\ at 5000 AU from the driving source. ",
        "introduction": "HH\\,211 is a warm and energetic molecular outflow located in the IC\\,348 complex at 315 pc, which was discovered by McCaughrean \\et \\ (1994) from observations of H$_2$ (at 2.12 $\\mu$m). The inclination from the plane of the sky is supposed to be small, around 10$\\degr$ (Hirano \\et\\ 2005).  The outflow in the CO(2--1) transition shows a well collimated structure at high velocities, and traces the outflow cavity walls at low velocities (Gueth \\& Guilloteau 1999, hereafter GG99). On the other hand, the SiO emission has been detected  with single-dish telescopes up to the 11--10 transition, indicative of gas densities $> 10^6$ \\cmt\\ along the jet (Nisini \\et \\ 2002; Chandler \\& Richer 2001). Recently, Hirano \\et \\ (2005) observed HH\\,211 in the SiO(5--4) transition and found a highly collimated structure consisting of a chain of knots. The innermost knots likely trace the primary jet launched at the close vicinity of the protostar. In order to follow up the study of the excitation conditions along the outflow, we have carried out observations of the high-J transitions CO(3--2) at 345.796 GHz and SiO(8--7) at 347.331 with high-angular resolution. ",
        "conclusions": "\\subsection{SiO(8--7) versus SiO(5--4) \\label{sdsiosio}}  Observations toward HH\\,211 in the SiO(5--4) transition at 217 GHz were carried out by Hirano \\et \\  (2005), with similar angular resolution. In order to compare the maps from both transitions, we convolved the moment-zero maps (integrated over all velocities) with a Gaussian to achieve a final beam of $1\\farcs95$, the largest major axis of the SiO(8--7) and SiO(5--4) beams.    We computed the SiO(8--7)/SiO(5--4) ratio map after correcting for the offset found  in the position of the continuum source of both images ($\\sim 0.2''$ in declination), and the  result is shown in Fig.\\ \\ref{sio8754ratiomap}. The uncertainty in the ratio is $\\sim 20$\\%. While the value for the ratio at the position of the innermost knots, B1 and R1, is $\\sim 1$, far away from the protostar the  ratio decreases down to $\\sim 0.5$. Comparing the ratio with the results of LVG modeling of Nisini \\et\\  (2002),  we set ranges for the density. Note that the  SiO(8--7)/SiO(5--4) ratio is not very sensitive to temperature variations for  $T \\gg 100$ K, since  the SiO(8--7) upper level energy is $\\sim 75$ K. At the position of B1 and R1 ($\\sim 500$ AU) we estimate that, for temperatures in the range 100--1000 K, the density must be (3--10)$\\times 10^6$ \\cmt.  On the other hand, $\\sim 15''$ (or 5000 AU) away from the center of the jet the ratio is 0.5, and this yields a density of (0.8--4)$\\times 10^6$ \\cmt. This is consistent with the density derived by Hirano \\et \\ (2005) from the SiO(5--4) jet and by Nisini \\et \\  (2002) from the single-dish data. Therefore, the density of the innermost knots seems to be about one order of magnitude higher than that of the knots further out along the jet.    \\subsection{SiO(8--7) versus CO(3--2)  \\label{sdsioco}}  Position-velocity (p-v) plots were computed for both CO(3--2) and  SiO(8--7) emission from 0.7 \\kms \\ wide channel maps along the axis of the jet,  PA$=116 \\arcdeg$. Since the CO emission is somewhat extended in the  direction perpendicular to the jet axis, we have smoothed the image with a  gaussian twice the beam of the observations, and with a position angle  perpendicular to the axis of the jet, in order to enhance the S/N of the CO emission in the p-v plot.  An overlay of SiO on the CO p-v plot is presented in  Fig.\\ \\ref{fpv}.  First of all, there is a distinct gap in CO emission from 7 to 8  \\kms, affecting all positions  along the jet. CO(3--2) and (1--0) observations of low angular resolution ($\\sim 15''$; Hirano \\et, in preparation) show an absorption feature at the same velocity. The gap is probably  due to an  intervening cold cloud along the line of sight.  In the CO(3--2) p-v plot, one can see a low-velocity component, extending along all positions and  tracing the shell structure seen in the low-velocity map from Fig.\\ \\ref{fm0}a,  and a second component tracing the high-velocity material, with velocities  increasing with distance from the protostar (Hubble-law), up to velocities of  $\\sim -14$ \\kms \\   (blue side) and $\\sim 40$ \\kms \\  (red side). Note that for CO no high-velocity emission comes from the  positions close to the protostar.  As for SiO(8--7), the p-v plot shows several features.  Contrary to the CO case, only very weak SiO emission is coming from the low-velocity material. The SiO emission resembles the CO emission at distances greater than $\\sim 5''$, with velocities increasing  with distance. However, the most remarkable characteristic of the SiO p-v plot is that the emission close to the protostar has the widest  range of velocities, including the highest, up to $-20$ \\kms \\ in the blue side and up to 40 \\kms \\ in the red side. This is a striking feature of the SiO jet,  that is,  that very high velocities are found very close ($\\sim 500$ AU) to the  protostar.  A possible explanation for the wide range of velocities found for SiO at the spatial scales of the disk would be that SiO close to the protostar is tracing a protostellar wind, with a large opening angle, and thus yielding a  maximum spread of velocities. This interpretation favors a wide-angle wind as the mechanism for driving the outflow (\\eg \\ Shu \\et  \\ 1991), since a pure jet model, in which velocity vectors point only in the polar direction, cannot produce this feature (\\eg \\ Masson \\& Chernin 1993; Smith \\et \\  1997). However, the overall structure of the SiO emission is highly collimated, and is very reminiscent of a pure jet (see Fig.\\ \\ref{fm0}b,c). In the wind-driven model, such a collimated structure would be the densest portion of a wide-angle wind. In this model, at distances $\\sim 500$ AU (the position of the  innermost SiO knots), the density decreases steeply with distance perpendicular to the axis of the jet, while the velocity vectors still span a wide angle around the jet axis (see \\eg \\ Shang \\et \\ 1998).  Note that the SiO emission close to the driving source is not due to mixing with entrained material.  The highly collimated morphology, together with the observed high velocity (up to 40 \\kms), and the derived high density ($> 10^6$ \\cmt) and high temperature ($>300$ K, Hirano \\et \\ 2005) for the SiO gas very close to the center are strongly suggestive of material from the primary jet, and not of ambient material being entrained.  Finally, we also see from the p-v plot that for distances larger than $\\sim 5''$ (1500 AU), the velocity of both CO and SiO emission increases with distance. This is consistent with high-velocity CO and SiO tracing entrained material, dragged by the primary jet.  However, the velocities reached by SiO are $\\sim 5$ \\kms \\ higher than those of CO, especially on the red side. This feature is also seen when superposing the SiO(5--4) p-v plot on the CO(2--1) (Hirano \\et \\ 2005). Presumably, SiO at distances larger than $\\sim 5''$ from the protostar comes from entrained material with higher density than CO. This is consistent with the critical density of SiO(8--7) being higher than that of CO(3--2). Such higher-density material would be likely closer to the primary jet, resulting in entrained SiO showing higher velocities than entrained CO.  In order to  compare the peak positions of the high-velocity SiO emission with those of the CO emission,  we  cross correlated the map of high-velocity SiO with that of high-velocity CO.   The cross correlation function on the blue side of the jet has a single maximum  at $3\\farcs8$, meaning that the SiO knots are on average $\\sim 4''$ closer to the  protostar than the CO knots on the blue side. For the red side, the cross  correlation function had two maxima, one at $\\sim 1''$, and the other one at  $\\sim 6''$, with the SiO knots closer to the protostar. Thus, for both sides of the jet, the SiO knots are found closer to the driving source than the CO knots. This suggests that chemical differentiation can be important in the jet.  In particular, we measured (in the high-velocity maps) the distance to the center for the SiO knots B1 and R1, and found that the peak is at $\\sim 1\\farcs5$ (470 AU) from the center. Since B1 and R1 are only slightly resolved by the beam (even if we include the low-velocities), no significant  emission of SiO is closer to the protostar than the peak of the innermost knots.   In summary, from the comparison of SiO(8--7) with SiO(5--4) and CO(3--2), it seems that the SiO(8--7) emission close to the protostar has contributions from the  primary jet, which could be driven by a wide-angle wind. At projected distances $\\ge 1500$ AU from the protostar, the SiO(8--7) shows velocities increasing with distance, likely tracing entrained gas. For the same distances, the CO(3--2) also shows velocities increasing with distance, but reaching systematically lower velocities than the SiO(8--7). We interpret this feature as SiO(8--7) tracing entrained gas which is denser, and therefore closer to the primary jet, than the entrained CO(3--2) gas."
    },
    "0512/astro-ph0512130_arXiv.txt": {
        "abstract": "We present the first high-resolution ($R \\approx 31\\,000$) spectra of the cool sdL \\twom, and what was originally identified as an early-type L subdwarf (sdL) \\lsr. Our work, in combination with contemporaneous work by Cushing and Vacca, makes it clear that the latter object is more probably a mid-M dwarf with an unusual composition that gives it some sub-dwarf spectral features.  We use the data to derive precise radial velocities for both objects and to estimate space motion; both are consistent with halo kinematics. We measure the projected rotational velocities, revealing very slow rotation for the old sd?M6 object \\lsr.  \\twom\\ exhibits rapid rotation of $v\\,\\sin{i} = 65 \\pm 15$\\,km\\,s$^{-1}$, consistent with the behavior of L dwarfs. This means that the braking time for L dwarfs is extremely long, or that perhaps they never slow down.  A detailed comparison of the atomic Rb and Cs lines to spectra of field L dwarfs shows the spectral type \\twom\\ is consistent with being mid- to late-L. The Rb~I and K~I lines of \\lsr\\ are like an early-L dwarf, but the Cs~I line is like a mid-M dwarf.  The appearance of the Ca~II triplet in absorption in this object is very hard to understand if it is not as least as warm as M6. We explain these effects in a consistent way using a mildly metal-poor mid-M model. M subdwarfs have weak metal-oxides and enhanced metal-hydrides relative to normal M dwarfs. \\lsr\\ exhibits metal-hydrides like an M dwarf but metal-oxides like a subdwarf. The same explanation that resolves the atomic line discrepancy explains this as well.  Our spectra cover the spectral region around a previously unidentified absorption feature at 9600\\,\\AA, and the region around 9400\\,\\AA\\ where detection of TiH has been claimed. We identify the absorption around 9600\\,\\AA\\ as due to atomic lines of Ti and a small contribution of FeH, but we cannot confirm a detection of TiH in the spectra of cool L subdwarfs. In \\twom, both metal-oxides \\emph{and} metal-hydrides are extremely strong relative to normal L dwarfs.  It may be possible to explain the strong oxide features in \\twom\\ by invoking effects due to inhibited dust formation. High resolution spectroscopy has aided in beginning to understand the complex molecular chemistry and spectral formation in metal-deficient and ultracool atmospheres, and the properties of early ultralow-mass objects. ",
        "introduction": "The most immediate relics of the early Galaxy are the cool subdwarfs of spectral type sdK and later -- lifetimes of such cool stars are well in excess of the age of the Galaxy. Their metal-deficient atmospheres make them appear hotter than solar-metallicity main-sequence stars of the same mass, which in turn renders them ``subluminous'' \\citep{Kuiper39}.  The spectral classification of solar abundance K- and M-dwarfs is well understood, and is entirely due to changes in effective temperature. Classification of M-subdwarfs is based on the metallicity sensitive ratio of absorption bands of metal-oxides and metal-hydrides \\citep{Gizis97, Lepine03a}; metal-hydrides are generally stronger in metal-deficient atmospheres, while metal-oxides are weaker \\citep[e.g.,][]{Mould76}. In the ultracool atmospheres of late M- and L-dwarfs, however, refractory grains become an important ingredient since they are competing with molecules for available metals. Furthermore, their opacity influences the optical depth, and hence which part of the atmosphere is visible. Since the formation and distribution of dust grains is not well understood even in solar abundance stars, this is an even more severe problem in the classification of metal-deficient L-subdwarfs.  The last two years have seen the first discoveries of L-type subdwarfs.  These ultracool dwarfs exhibit colors too red for an M-type object, and do not fit on the expanded subdwarf classification scheme of \\cite{Lepine03a}. Another argument for surface temperatures lower than M-dwarf temperatures is the extremely strong absorption lines of alkali atoms as K~I, Na~I, and Rb~I. The first L-type subdwarf, 2MASS~J05325346+8246465 (hereafter \\twom), was discovered by \\citet[][hereafter B03]{Burgasser03}. Its spectrum is very similar to the L7 dwarf \\denis, exhibiting strong alkali lines \\citep[we use the optical classification scheme of][]{Kirkpatrick99}.  \\twom\\ shows enhanced metal hydride bands, but in contrast to M-subdwarfs it also has strongly enhanced bands of TiO. The first proposed early-type L-subdwarf, \\lsr, was discovered by \\citet[][hereafter L03]{Lepine03b}.  It also exhibits strong Rb~I absorption and L03 report enhanced CaH and TiO bands.  Metal hydrides FeH and CrH, however, are relatively weak. We conclude in this paper that \\lsr\\ is not really an L subdwarf, but in the M temperature range. The same conclusion is reached in a paper based on lower resolution spectra that was posted during our refereeing process by \\citet[][hereafter CV05]{CV05} . The spectral peculiarities of these unusual objects provide a preview of the complex chemical processes occurring in cool atmospheres of metal-deficient L-dwarfs. The list of subdwarfs of spectral type L and late M is constantly growing; \\citet{Burgasser04a} discovered the L-subdwarf LSR~$1626 + 3925$, and two late M-type subdwarfs at the end of the subdwarf classification scheme with features very similar to the L-type subdwarfs have been reported by \\citet[][SSSPM~J1013\\,--\\,1356]{Scholz04a} and \\citet[][SSSPM~J1444\\,--\\,2019]{Scholz04b}.  The presumably large age of M- and L-subdwarfs also makes them ideal tracers of the rotational evolution of late-type objects. Cool stars of spectral type K and M are known to suffer rotational braking on a relatively short timescale (1 Gyr), but no projected rotation velocity below $v\\,\\sin{i} = 10$\\,km\\,s$^{-1}$ has been reported in an L-dwarf \\citep{Mohanty03, Bailer04}. Whether this is due to extremely long braking times or the lack of rotational braking at all in L-dwarfs may be answered by observations of the rotation of the oldest objects of very late spectral type.  So far, no high-resolution spectrum of an L-type subdwarf has been available to investigate the details of line strengths, take a close look for unidentified features detected in L-subdwarf spectra, \\citep[among them the tentative detection of TiH,][]{Burgasser04b}, and to investigate rotation velocities. In this paper, we present the first high resolution spectra of the late-type sdL \\twom\\ and the peculiar sd?M \\lsr.  We describe our observations in \\S~2.  Radial velocities are calculated in \\S~3, where we also derive new constraints on the space motion on the basis of the accurate radial velocities. Surface rotation is investigated in \\S~4, and individual spectral features are discussed in \\S~5. ",
        "conclusions": "We have presented the first high resolution spectra of the peculiar possibly metal-poor sd?M6 \\lsr\\ and the late-type L-subdwarf \\twom. We have derived accurate radial and rotation velocities for both objects, and we compared their spectral features to spectra of known field L-dwarfs. The spectra of \\lsr\\ and \\twom\\ are very different in many respects. These are discussed below.  The radial velocity of \\lsr\\ is consistent with the value reported by L03 if we assume that their determination suffered similar systematic effects as mentioned by B03 for the case of \\twom. For the latter, we derive a radial velocity consistent with the result of B03. We computed the space motion in order to check membership of the halo population. Our result for \\lsr\\ significantly differs from the result reported by L03; the radial velocity of \\lsr\\ does not fall outside the $2\\sigma$-region reported for thick disk stars by \\cite{Chiba00}. However, the $V$-velocity of \\lsr\\ ($-111$\\,km\\,s$^{-1}$) is comparably high, and the space motion vector does not contradict its halo membership. For \\twom, our estimate of the $V$-velocity is $V = -285$\\,km\\,s$^{-1}$, which is in strong support of halo membership.  The rotation velocity of \\lsr\\ is below the detection limit of about 5\\,km\\,s$^{-1}$. This is consistent with the bulk of its properties more closely resembling the M5.5V dwarf Gl\\,406 than an L1 object. We have analyzed the reasons why it was initially classified as an early-L subdwarf, and explain them with a mild metal-deficiency. For the true L subdwarf, \\twom, we measure a rotation velocity of $v\\,\\sin{i} = 65$\\,km\\,s$^{-1}$. If this object is indeed a relic of the young galaxy, that means that the braking time for a mid L-dwarf is much longer than 10\\,Gyr. Either the braking time exceeds the age of the galaxy by this spectral type, or there is no rotational braking in mid to late L-type objects at all.  In M-subdwarfs metal-oxides are weaker and metal-hydrides are stronger than in the metal rich M-dwarfs, but both species are stronger in \\twom\\ than in its metal-rich counterpart \\denis. The strong oxide-bands in \\twom\\ can be explained by invoking a reduced level of dust formation in the metal-deficient atmosphere, which would leave behind more TiO than in the metal rich L-dwarfs (B03). The extremely strong metal-hydride bands in \\twom\\ are generally consistent with observations of enhanced metal-hydride bands in M-subdwarfs.  Chemistry and line formation in the ultracool atmospheres of L-dwarfs is much more complex than it is in hotter stars because of dust formation and depletion of refractory metals from the atmosphere. The chemical anomalies of subdwarfs heighten these complexities, and currently yield a confusing situation.  High-resolution spectra provide crucial information for our eventual understanding of the ultracool metal-deficient atmospheres which were common in the early galaxy."
    },
    "0512/astro-ph0512289_arXiv.txt": {
        "abstract": "{We present the results of a photometric multisite campaign on the $\\delta$ Scuti Pre-Main-Sequence star IP Per.  Nine telescopes have been involved in the observations, with a total of about 190 hours of observations over 38 nights. Present data confirms the multiperiodic nature of this star and leads to the identification of at least nine pulsational frequencies. Comparison with the predictions of linear non-adiabatic radial pulsation models allowed us to identify only five of the nine observed frequencies, and to constrain the position of IP Per in the HR diagram. The latter is in good agreement with the empirical determination of the stellar parameters obtained by \\cite{miro}.  An initial interpretation of the observed frequencies using the Aarhus non-radial pulsation code suggests that three frequencies could be associated with non-radial ($l$=2) modes. \\par Finally, we present new evolutionary and pulsation models at lower metallicity (Z=0.008) to take into account the possibility that IP Per is metal deficient, as indicated by \\cite{miro}. ",
        "introduction": "Intermediate mass (M $\\geq$ 1.5 M$_{\\odot}$) Pre-Main-Sequence (PMS) stars are known as Herbig Ae/Be stars (\\cite{herbig}). This class of stars is characterized by spectral type A or B with emission lines, an infrared excess due to hot or cool circumstellar dust or both, and luminosity class III to V (\\cite{waters}).  Herbig Ae/Be are also well known for their photometric and spectroscopic variability on time scales of minutes to years mainly due to photospheric activity and interaction with the circumstellar environment (see e.g. \\cite{catala}).  However, the fact that these young stars during their contraction towards the Main-Sequence (MS) move across the pulsation instability region of more evolved stars has prompted the suggestion that at least part of the activity could be due to stellar pulsation (see \\cite{baade}; \\cite{kurtz}).   \\begin{table} \\caption{Journal of the observations. Note that the data coming from Sierra Nevada and San Pedro Martir (SPM) have been transformed in the Johnson system as discussed in section 2.2.} \\label{jou} \\begin{tabular}{cccc} \\hline \\noalign{\\smallskip} HJD-2450000 & HJD-2450000  & Duration &  Filter \\\\ start (days)         &   end (days)        &  (hours) &         \\\\ \\noalign{\\smallskip} \\hline & \\bf{Loiano CCD  (Italy)}  & & \\\\ \\noalign{\\smallskip} 2545.585 &\t2545.659&  1.8 & B \\\\ 2547.599 &\t2547.681&  2.0 & B\\\\ 2548.547 &\t2548.611&  1.5 & B\\\\ 2549.503 &\t2549.623&  2.9 & B\\\\ 2628.347 &\t2628.627&  6.7 & B\\\\ 2629.286 &\t2629.453&  4.0 & B\\\\ 2630.454 &\t2630.513&  1.4 & B\\\\ 2654.252 &\t2654.332&  1.9 & B\\\\ \\noalign{\\smallskip} &\\bf{Fairborn-APT (USA)}\t&      &\\\\ \\noalign{\\smallskip} 2976.875 &\t2976.892&   0.4&  BV\\\\ 2977.851 &\t2977.936&   2.0&  BV\\\\ 2978.721 &\t2978.933&   5.1&  BV\\\\ 2983.573 &\t2983.920&   8.3&  B \\\\ 2984.711 &\t2984.917&   4.9&  BV\\\\ 2987.765 &\t2987.902&   3.3&  BV\\\\ \\noalign{\\smallskip} &\\bf{SOAO (Korea)}    &      &\\\\ \\noalign{\\smallskip} 2977.170 &\t2977.333&   3.9&  BV\\\\ 2980.910 &\t2981.063&   3.7&  V\\\\ 2983.007 &\t2983.191&   4.4&  BV\\\\ \\noalign{\\smallskip} & \\bf{OSN (Spain)}\t&      &\\\\ \\noalign{\\smallskip} 2988.267 &\t2988.500&   5.6&  uvby\\\\ 3047.364 &\t3047.472&   2.6&  uvby\\\\ 3048.299 &\t3048.459&   3.8&  uvby\\\\ \\noalign{\\smallskip} &\\bf{Loiano TTCP (Italy)}\t&      &\\\\ \\noalign{\\smallskip} 2946.330 &\t2946.703&   9.0&  BV\\\\ 2947.287 &\t2947.441&   3.7&  BV\\\\ 2973.393 &\t2973.636&   5.8&  BV\\\\ \\noalign{\\smallskip} &\\bf{SARA (USA)} &      &\\\\ \\noalign{\\smallskip} 2970.631 &\t2970.756&   3.0&  B\\\\ \\noalign{\\smallskip} &\\bf{SPM (Mexico)}    &      &\\\\ \\noalign{\\smallskip} 2970.494 &\t2970.818&  7.8 & uvby\\\\ 2972.495 &\t2972.841&  8.3 & uvby\\\\ 2973.493 &\t2973.775&  6.8 & uvby\\\\ 2974.487 &\t2974.585&  2.4 & uvby\\\\ 2977.485 &\t2977.847&  8.7 & uvby\\\\ \\noalign{\\smallskip} & \\bf{BAO (China)}\t&      &\\\\ \\noalign{\\smallskip} 2972.974 &\t2973.383&  9.8 & V\\\\ 2974.032 &\t2974.146&  2.7 & V\\\\ 2974.939 &\t2975.382& 10.6 & V\\\\ 2976.940 &\t2977.078&  3.3 & V\\\\ 2984.196 &\t2984.292&  2.3 & V\\\\ \\noalign{\\smallskip} &\\bf{Teide-OGS (Spain)}&     &\\\\ \\noalign{\\smallskip} 2977.506 &\t2977.603&  2.3 & V\\\\ \\noalign{\\smallskip} & \\bf{Serra la Nave (Italy)}\t  \t&      &\\\\ \\noalign{\\smallskip} 2966.462 &\t 2966.62&  3.9&  BV\\\\ 2967.337 &\t 2967.61&  6.7&  BV\\\\ 2973.348 &\t 2973.61&  6.4&  BV\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table}   The possible presence of pulsators among Herbig Ae/Be stars is particularly attractive since the precise observables which can be measured, i.e.  the pulsation frequencies can, in principle, allow us to test evolutionary models by constraining the internal structure using asteroseismological techniques.  The existence of pulsating Herbig stars was originally suggested by \\cite{breger1} who discovered two candidates in the young open cluster NGC 2264. This initial finding was confirmed by subsequent observations of $\\delta$ Scuti-like pulsations in the Herbig Ae stars HR5999 (\\cite{kurtz}) and HD104237 (\\cite{donati}).\\par  This empirical evidence stimulated the first theoretical investigation of the PMS instability strip based on non-linear convective hydrodynamical models (Marconi \\& Palla 1998).  As a result, the topology of the PMS instability strip for the first three radial modes was identified.  \\cite{marconi} also pointed out that the interior structure of PMS stars entering the instability strip differs significantly from that of more evolved Main Sequence stars (with the same mass and temperature), even though the envelopes structures are similar. This property was subsequently confirmed by \\cite{suran} who made a comparative study of the seismology of a 1.8 $M_{\\odot}$ PMS and post-MS star. \\cite{suran} found that the unstable frequency range is roughly the same for PMS and post-MS stars, but that some non-radial modes are very sensitive to the deep internal structure of the star. In particular, it is possible to discriminate between the PMS and post-MS stage using differences in the oscillation frequency distribution in the low frequency range ($g$ modes, see also \\cite{templeton}). \\par  Up to now new observational programs have been carried out by various groups. The current number of known or suspected candidates amounts to about 30 stars (see the updated list at http://ams.astro.univie.ac.at/pms\\_corot.php, and the reviews by \\cite{zwintz}, \\cite{marconi2004} and \\cite{marconi2004a}).  However, only a few stars have been studied in detail, so that the overall properties of this class of variables are still poorly determined.  In this context, our group has started a systematic monitoring program (see \\cite{marconi2001}, \\cite{h254}, \\cite{v346}, \\cite{v351}, \\cite{bernabei}) of Herbig Ae stars with spectral types from A to F2-3 with the following aims: 1) to identify the largest number of pulsating objects in order to observationally determine the boundaries of the instability strip for PMS $\\delta$ Scuti pulsation; 2) to study in detail through multisite campaigns selected objects showing multiperiodicity (see \\cite{marconi2001}, \\cite{h254},\\cite{v346}, \\cite{v351}, \\cite{bernabei}).  The multiperiodic pulsators are potential candidates for future asteroseismological analysis.   \\begin{table}  \\caption[]{List of telescopes and instruments involved in the multisite campaign. SCP means Single Channel Photometer; TCP means Three Channel Photometer. To give an idea of the relative photometric precision of each site, the last column shows the white noise in the Fourier space (amplitude) for each site. All the calculations are in the V filter except for Loiano (2002) and SARA sites for which we have only B filter data.} \\label{tab1} \\begin{tabular}{lccc} \\hline \\noalign{\\smallskip} Observatory  & Telescope & Instrument & noise \\\\ & & & (mmag)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Loiano (Italy)    & 1.5m  &  BFOSC- 2002  & 0.2 \\\\ Loiano (Italy)    & 1.5m  &  TTCP - 2003  & 0.2 \\\\ BAO (China)       & 0.85m &  TCP & 0.4 \\\\ SPM (Mexico) & 1.5m  & $uvby$ Phot. & 0.6\\\\ SARA (USA)  & 0.9m  & CCD  &  1.5 \\\\ Teide (Spain)     & 1.0m OGS &  CCD   & 0.6 \\\\ Fairborn (USA)    & 0.75m T6 & SCP  & 0.8 \\\\ OSN (Spain) & 0.9m &  $uvby$ Phot.  & 0.6\\\\ Serra la Nave (Italy) & 0.9m & SCP & 1.3 \\\\ SOAO (Korea)         & 0.6m & CCD  & 0.7 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular}  \\end{table}   During this project our attention turned to the star IP Per, already listed as Herbig Ae star by \\cite{the} with spectral type A3e, and studied in detail by \\cite{miro} both photometrically and spectroscopically.  The main properties of this interesting object are:  \\begin{itemize}  \\item IP Per is a typical UX Ori type star showing photometric variations with an amplitude of $\\sim$0.3 mag and with duration of the minima of 10-50 days.  \\item the fundamental stellar parameters are: spectral type A7, $T_{\\rm eff} \\sim 8000\\pm100$ K, $\\log g \\sim 4.4 \\pm 0.1$, $\\xi_t=2.0\\pm1.0$ km s$^{-1}$ and [M/H] $\\simeq -0.41\\pm0.1$.  Thus, IP Per is a dwarf metal-poor star.  \\item its radial velocity and proper motion suggest that it most likely belongs to the Per OB2 association ($D \\simeq$ 300 pc, \\cite{zeeuw}). At $D \\simeq$ 300 pc and $\\log L/L_{\\odot}$=1.0, IP Per falls onto the ZAMS at a mass of 1.8 M$_{\\odot}$.  \\end{itemize}  In the following, we present the results of a photometric multisite campaign to study in detail the pulsational properties of IP Per.  In Section 2. we describe the observations and data reduction techniques; in Section 3. we discuss the frequency analysis; Section 4 presents a comparison with theoretical models. A brief discussion concludes the paper. Preliminary results of our observations have been presented in \\cite{ripmarc} and \\cite{ripepiiau}. ",
        "conclusions": "A total of about 190 hours of observations obtained during 38 nights at 9 different telescopes on the PMS $\\delta $ Scuti star IP Per have been presented. The Fourier analysis of this data set confirms the multiperiodic nature of this pulsator: we have identified nine frequencies of pulsation which are significant on the basis of the \\cite{breger93} and \\cite{kus97} criteria.  Comparison with the predictions of linear non-adiabatic radial pulsation models allows us to recover only five out of the nine observed frequencies for a 1.77 $M_{\\odot}$ model.  Non-radial pulsation is also present in this star. A preliminary interpretation of the observed frequencies through the Aarhus non-radial code, applied to the evolutionary structure of the 1.77 $M_{\\odot}$ model reproducing $f_1$, $f_2$, $f_4$, $f_5$ and $f_9$ with radial modes, seems to indicate that $f_3$, $f_6$ and $f_8$ are associated with non-radial modes with $l=2$. Specific post-MS evolutionary and pulsational models were computed in order to investigate the dependence of radial and non-radial output frequencies on the assumed evolutionary status. The resulting post-MS solution has similar stellar parameters and $p$ mode frequencies.\\par  \\medskip Finally the possible effect of the metal poor nature of IP Per detected by \\cite{miro} on both pulsation and evolutionary properties is discussed.  We find that if the metallicity of IP Per is as low as Z=0.008 the best fit radial model has a significantly lower mass than the case at solar chemical composition but the pulsation characteristics are similar.Also the estimated position of IP Per in the HR diagram appears to be in good agreement with the independent determination by \\cite{miro}. Whether the low metallicity is a property only of the surface layers or represents a systematic deficit throughout the interior, as we have assumed in our modeling, should be clarified before final conclusions on the stellar parameters of IP Per can be reached."
    },
    "0512/astro-ph0512240_arXiv.txt": {
        "abstract": "We report on new X-ray observations of the large-scale jets recently discovered in X-rays from the black hole candidate 4U 1755-33.  Our observations in 2004 show that the jets found in 2001 are still present in X-rays.  However, sensitive radio observations in 2004 failed to detect the jets.  We suggest that synchrotron radiation is a viable emission mechanism for the jets and that thermal bremsstrahlung and inverse-Compton emission are unlikely on energetic grounds.  In the synchrotron interpretation, the production of X-rays requires acceleration of electrons up to $\\sim 60$ TeV, the jet power is $\\sim 4 \\times 10^{35} \\rm \\, erg \\, s^{-1}$, and the radio non-detection requires a spectral index $\\alpha > -0.65$ ($S_{\\nu} \\propto \\nu^{\\alpha}$) which is similar to the indexes found in lobes surrounding some other compact objects.  We find an upper limit on the flux of 4U 1755-33 in quiescence of $5 \\times 10^{-16} \\rm \\, erg \\, cm^{-2} \\, s^{-1}$ in the 0.3-8~keV band. ",
        "introduction": "Relativistic jets are often produced by accreting compact objects, including both stellar-mass X-ray binaries and active galactic nuclei. Jets appear to be important in the dynamics of the overall accretion flow in such systems and a substantial fraction of the accretion energy of X-ray binaries may be dissipated in jets.  Understanding the properties and production of jets is important for our understanding of the energetics and dynamics of the accretion process.  Further, the jets from X-ray binaries in the Milky Way are a potentially significant source of energy input to the Galactic ISM and, if hadronic, a likely source of a cosmic rays \\citep{heinz02} and a possible source of Galactic light element nucleosynthesis \\citep{butt03}.   Recent XMM-Newton observations have led to the discovery of a large scale X-ray jet from the long-term X-ray transient and black hole candidate 4U 1755--33 \\citep{angelini03}.  The X-ray source 4U 1755--33 was discovered with Uhuru \\citep{jones77} and was, therefore, active in 1970 and may have been active earlier.  It was later found to have an unusually soft spectrum \\citep{whitem84,white84} and a hard X-ray tail \\citep{pan95} suggesting a black hole candidate.  The source shows X-rays dips which indicate that the system has a high inclination and an orbital period of 4.4 days \\citep{white84,mason85}.  The source was still active in 1993 \\citep{church97}, and then found in quiescence in 1996 \\citep{roberts96}.  The source was active for at least 23 years. The distance to the source is poorly constrained.  It is likely greater than 4~kpc because the optical counterpart (identified during outburst; McClintock, Canizares, \\& Hiltner 1978) was not detected in quiescence \\citep{wachter98}, but less than 9~kpc because of the low level of visual extinction \\citep{mason85}.  \\citet{angelini03} found a linear X-ray structure which is roughly symmetric about the position of 4U 1755--33 extending about $3\\arcmin$ to the north-west and $3\\arcmin$ to the south-east of the black hole candidate.  There appear to be multiple knots in the jets.  For estimated distances of 4--9~kpc, the angular size corresponds to jet lengths of 3--8~pc. Therefore, the jet must have taken at least 10-30 years to form.  The source was active for at least 23 years which appears sufficient to have formed the jet.  Chandra observations do not show point sources along the jet, but do give a detection of emission over the area of the jet \\citep{park05}.  This indicates that the jet emission seen with XMM-Newton is truly diffuse and not an alignment of point sources.  The primary questions concerning the jet are: what is the X-ray emission mechanism, how is the jet powered, and what is the total energy in the jet?  We obtained new observations of 4U 1755--33 using XMM-Newton and the Australia Telescope Compact Array (ATCA) to attempt to measure a multiwavelength spectrum and observe the evolution of the jet.  We describe the X-ray observations and analysis in \\S 2, the radio observations and analysis in \\S 3, and draw conclusions regarding the properties of the jet in \\S 4.   \\begin{deluxetable}{rlllrrl} \\tabletypesize{\\scriptsize} \\tablecaption{XMM-Newton X-ray sources near 4U 1755-33 \\label{xmmsources}} \\tablewidth{0pt} \\tablehead{ & \\colhead{RA} & \\colhead{DEC} & \\colhead{Error} & \\colhead{Flux A}  & \\colhead{Flux B}  & \\colhead{Counterparts} } \\startdata 1 & 17 59 00.86 & -33 45 48.1 & 0.2 & 243.9 $\\pm$10.0 & 171.9 $\\pm$ 4.0 & AW2003 1, 2mass \\\\ 2 & 17 57 58.88 & -33 46 24.9 & 0.4 & 218.4 $\\pm$12.1 &  45.1 $\\pm$ 2.7 & 2mass \\\\ 3 & 17 59 21.83 & -33 53 15.6 & 0.7 &  77.1 $\\pm$ 7.1 &  14.9 $\\pm$ 1.9 & AW2003 3 \\\\ 4 & 17 58 26.89 & -33 59 53.8 & 0.8 &  76.2 $\\pm$ 9.7 &    -            & \\\\ 5 & 17 58 42.02 & -33 41 48.8 & 0.3 &  29.3 $\\pm$ 3.8 &  71.6 $\\pm$ 2.9 & 1WGA J1758.6-3341 \\\\ 6 & 17 59 22.79 & -33 50 25.2 & 0.7 &  63.8 $\\pm$ 7.0 &  24.4 $\\pm$ 2.3 & 1WGA J1759.3-3350, 2mass \\\\ 7 & 17 58 20.25 & -33 42 51.8 & 0.6 &  63.4 $\\pm$ 7.7 &  33.7 $\\pm$ 2.2 & \\\\ 8 & 17 58 49.48 & -33 41 40.8 & 0.5 &  56.6 $\\pm$ 5.7 &  42.7 $\\pm$ 2.7 & 2mass \\\\ 9 & 17 58 36.67 & -33 40 27.3 & 0.6 &    -            &  55.0 $\\pm$ 5.7 & \\\\ 10 & 17 58 38.88 & -33 57 00.3 & 0.4 &    -            &  49.0 $\\pm$ 2.5 & \\\\ 11 & 17 58 21.03 & -33 46 54.7 & 0.3 &  34.1 $\\pm$ 4.3 &  41.8 $\\pm$ 1.9 & 2mass \\\\ 12 & 17 59 27.97 & -33 49 10.3 & 1.5 &  37.0 $\\pm$ 4.5 &    -            & \\\\ 13 & 17 58 44.04 & -33 46 13.0 & 0.5 &  29.2 $\\pm$ 3.4 &  11.7 $\\pm$ 1.2 & \\\\ 14 & 17 57 40.65 & -33 45 53.9 & 0.6 &    -            &  25.4 $\\pm$ 2.3 & \\\\ 15 & 17 58 59.19 & -33 50 16.3 & 0.5 &  25.3 $\\pm$ 3.2 &  20.4 $\\pm$ 1.5 & \\\\ 16 & 17 58 44.16 & -33 45 09.5 & 0.6 &  24.5 $\\pm$ 3.2 &  12.2 $\\pm$ 1.4 & \\\\ 17 & 17 59 03.24 & -33 59 18.1 & 0.9 &    -            &  17.8 $\\pm$ 2.2 & \\\\ 18 & 17 58 51.23 & -33 49 13.8 & 0.7 &    -            &  17.1 $\\pm$ 1.4 & \\\\ 19 & 17 58 25.56 & -33 57 46.8 & 2.8 &    -            &  15.8 $\\pm$ 1.7 & \\\\ 20 & 17 58 24.23 & -33 52 58.8 & 0.5 &    -            &  14.8 $\\pm$ 1.3 & \\\\ 21 & 17 58 49.58 & -33 57 14.0 & 0.6 &    -            &  14.1 $\\pm$ 1.6 & 2mass \\\\ 22 & 17 58 12.44 & -33 43 56.6 & 0.9 &    -            &  14.0 $\\pm$ 1.7 & \\\\ 23 & 17 58 14.22 & -33 49 08.7 & 0.8 &    -            &  11.7 $\\pm$ 1.2 & \\\\ 24 & 17 58 43.37 & -33 58 17.8 & 1.0 &    -            &  11.6 $\\pm$ 1.5 & \\\\ 25 & 17 59 08.17 & -33 46 47.7 & 0.7 &    -            &   9.5 $\\pm$ 1.3 & \\\\ 26 & 17 58 46.92 & -33 48 44.7 & 1.0 &    -            &   8.0 $\\pm$ 1.1 & \\\\ 27 & 17 58 22.29 & -33 47 43.7 & 1.0 &    -            &   7.9 $\\pm$ 1.0 & \\\\ \\enddata  \\vspace{-12pt}\\tablecomments{Table~\\ref{xmmsources} includes for each source: the source number; RA and DEC -- the position of the source in J2000 coordinates; Error - the statistical error on the source position in arcseconds, note that this does not include errors in the overall astrometry; Flux A -- the source flux for observation A in units of $10^{-14} \\, \\rm erg \\, cm^{-2} \\, s^{-1}$ in the 0.3--10~keV band calculated assuming a power law spectrum with photon index of 1.5 and corrected for the Galactic absorption column density of $3.1 \\times 10^{21} \\rm \\, cm^{-2}$; Flux B -- the same for observation B; Counterparts - indicates counterparts in \\citet{angelini03} (AW2003), the 2Mass catalog (2mass), or the WGA catalog (WGA).} \\end{deluxetable}  \\begin{figure*}[tb] \\centerline{\\psfig{figure=f1.eps,width=5.75in,angle=0} } \\caption{XMM-Newton images of 4U 1755--33.  The image on the left is observation A (8 Mar 2001) on the right is observation B (18 Sep 2004). The green cross indicates the position of 4U 1755--33.  The green rectangle encloses the jet and has a size of $44\\arcsec$ by $430\\arcsec$.  On the left image, the green X's mark the positions of XMM-Newton point sources reported in Table~\\ref{xmmsources}.  On the right image, the green X's mark the positions of Chandra point sources reported in Table~\\ref{chandrasources}.  The green arrow indicates North and is $2\\arcmin$ long.  Radio contours are superimposed on the image on the right in magenta.  The radio contours represent the flux density at 13~cm at levels of 0.3, 0.6, 1.2, and 2.4~mJy.  Also, on the right image, the green circle indicates the extraction region used to find the spectrum of the jet bright spot described in the text.} \\label{bothobs} \\end{figure*} ",
        "conclusions": "There is now known to exist a broad range of jets from stellar-mass X-ray binaries.  Persistent compact jets with lengths of tens of A.U.\\ are produced in the low/hard X-ray spectral state.  Impulsive jets produced at state transitions have been detected on lengths scales from hundreds of A.U.\\ out to parsecs.  Stationary lobes have been found in the radio at separations of several parsecs for sources such as 1E1740.7-2942 and GRS 1758-258 and in the X-ray with separations up to 70~pc from SS 433 \\citep{watson83}.  The jet size of 3--8~pc in 4U 1755--33 is larger than the transient large-scale moving jets of XTE J1550-564 \\citep{corbel02} and H~1743$-$322 \\citep{corbel05}, but is similar to the total size of the stationary radio lobes of 1E1740.7-2942 and GRS 1758-258 \\citep{mirabel99}.  The latter two sources are persistent X-ray emitters, similar in properties to 4U 1755--33 while it was X-ray bright.  This may suggest that 4U 1755--33 represents the formation of a large-scale, nearly stationary jet.  The central source in 4U 1755--33 has turned off and we now see the decay of the jet over a time scale of 10--40 years or perhaps longer for the Northern part of the jet.  While we can place only a lower bound of 23 years on the time over which the central source was active, and the jet was being energized, our detection of the decay of the jet suggests that the time scale for the energization of the jet is similar to or shorter than the decay time.  We note that we detect flux decay only for the southern jet.  The change in morphology is also stronger for the inner regions of the southern jet.  This may suggest that the southern jet is the approaching and the northern jet is receding. Continued monitoring of 4U 1755--33 may reveal motion and decay of the northern jet.  A key question is the nature of the jet emission.  We consider three possible emission mechanisms: thermal bremsstrahlung, inverse-Compton, and synchrotron.  If the X-ray emission is thermal bremsstrahlung, then the total mass and energy of the jet can be estimated from its observed luminosity, temperature, and volume.  For the volume, we assume that the jet occupies a roughly cylindrical volume with a diameter perpendicular to the direction of motion of $0.2\\arcmin$ and a linear dimension of $6\\arcmin$ multiplied by a filling factor of 0.2.  The volume is then $2 \\times 10^{54} \\rm \\, cm^{3}$ for an assumed distance of 4~kpc.  We took the flux from the X-ray spectral fits, and set $kT = 5 \\rm \\, keV$ which is the minimum temperature consistent with the fits.  We find that the density of the jet material is $4 \\rm \\, cm^{-3}$, the total energy of the jet is $2 \\times 10^{47} \\rm \\, erg$ and the total mass in the jet is $1 \\times 10^{31} \\rm \\, g$.  The cooling time of the gas would then be $6 \\times 10^{6} \\rm \\, yr$ which is much longer than the observed decay time.  Also, if the jet were fed by the outflow from a mass accretion rate of $10^{19} \\rm \\, g/s$ corresponding the Eddington rate for a $10 M_{\\odot}$ compact object, at least $10^{4}$ years would be required to accumulate the needed mass.  This is much longer than the observed decay time scale and we conclude that the jet emission is unlikely to be thermal bremsstrahlung.  The best candidate source of seed photons for inverse-Compton scattering is the interstellar radiation field (ISRF).  The ISRF varies strongly with Galactocentric radius and height above the Galactic plane.  On the sky, 4U 1755-33 is rather close to the Galactic center and it could be physically close to the Galactic center, although the low optical extinction suggests that the source is actually nearer than the Galactic bulge.  We assume that 4U 1755-33 is at a distance of 8.5~kpc, close to the Galactic center, which maximizes the ISRF energy density.  We adopt an energy density of 10~eV~cm$^{-3}$, equal to the maximum found anywhere in the Milky Way \\citep{strong00} and assume, for simplicity, that all of the radiation is in the dominant component of ISRF near 1~$\\mu$m \\citep{mathis83}.  To produce X-rays in the 0.3-10~keV, electrons with energies from $\\sim$5 to $\\sim$100~MeV are required.  For an assumed X-ray spectral index $\\alpha = -0.5$, defined as $S(\\nu) \\propto \\nu^{\\alpha}$, the total energy in relativistic electrons required is $\\sim 10^{50} \\rm \\, erg$.  A $10 M_{\\sun}$ black hole producing energy at the Eddington rate with all of the energy going to perfectly efficient acceleration of electrons in the desired energy range would require $\\sim 2000$ years to power the jet.  This is much longer than the observed decay time scale.  Therefore, it appears unlikely that the jet emits via inverse-Compton radiation.  If X-ray emission is synchrotron, then synchrotron radio emission should be expected.  To determine if our upper limits on the radio flux are consistent with synchrotron emission, we consider the emission from the brightest X-ray spot along the jet.  The X-ray spectrum was extracted from a circle with a radius of $28\\arcsec$.  This is slightly larger than the radio beam size at 2.4~GHz, so we can compare the X-ray flux density in the regions to the radio upper limit per beam, which we take as 3 times the rms noise level of 0.2~mJy.  Comparing the X-ray and radio flux density, we derive a lower limit on the radio/X-ray spectral index $\\alpha > -0.65$, i.e. the spectrum must be flatter than $\\alpha = -0.65$.   Therefore, our radio observations may simply be not sensitive enough to detect the radio synchrotron emission.  The limit on the spectral index implies that the exponent, $p$, of the electron energy distribution, $N(E) \\propto E^{-p}$, must be $p < 2.3$.  This is within the range that can be produced by relativistic shocks.   The spectral index is similar to that measured for the large-scale jets of the black hole candidate X-ray transient XTE J1550-564 of $\\alpha = -0.660 \\pm 0.005$ \\citep{corbel02} and consistent with the index of $\\alpha = -0.45$ measured for the eastern lobe of SS 433 \\citep{safiharb99}.  To investigate the energetics of a jet radiating via synchrotron emission, we calculate the equipartition magnetic field.  Using a spectral index $\\alpha = -0.6$ which is consistent with the bound above, a lower frequency cutoff of 2.4~GHz, and the same assumption for distance and volume as for the thermal bremsstrahlung case, we find a magnetic field of 36~$\\mu$G.  The electrons producing the X-ray emission must have Lorentz factors up to $6 \\times 10^7$, corresponding to energies of up to 60 TeV.  The radiative lifetime of these electrons is of order 320~years.  The minimum total energy required is  $2 \\times 10^{44} \\rm \\, erg$ and the number of electrons needed to produce the observed radiation is $1.2 \\times 10^{46}$.  Assuming that the jet is composed of normal matter (i.e.\\ roughly one proton per electron), then the required mass is $2 \\times 10^{22} \\rm \\, g$.  These numbers are relatively insensitive to distance.  For a distance of 9~kpc, the magnetic field decreases to 29~$\\mu$G, the required energy increases to $1.5 \\times 10^{45} \\rm \\, erg$, and the required mass increases to $1.3 \\times 10^{23} \\rm \\, g$.  For the larger distance, the energy required corresponds to 2 weeks accumulation at the Eddington rate for a 10~$M_{\\sun}$ black hole and the mass could be accumulated in less than one day.  Thus, the energy and material required for a synchrotron emitting jet could easily have been accumulated over the 20 year active phase of the X-ray source.  Using the decay rate calculated for the entire jet and the energy estimate from the synchrotron equipartition calculation, the energy loss rate is then $\\sim 4 \\times 10^{35} \\rm \\, erg \\, s^{-1}$ for a 4~kpc distance.  This is about 1\\% of the Eddington luminosity for a 10~$M_{\\sun}$ black hole and, therefore, the jet could have been energized by the conversion of a few percent of the energy into relativistic electrons.  The available data on the jet of 4U 1755--33 are consistent with the X-rays being synchrotron emission.  Synchrotron emission also is the most favorable mechanism in terms of the required mass and energy. However, significantly deeper radio observations are required to test if the predicted synchrotron radio emission is produced."
    },
    "0512/astro-ph0512076_arXiv.txt": {
        "abstract": "We use the Type Ia Supernova {\\it gold} sample to constrain the parameters of dark energy models namelly the Cardassian, Dvali-Turner (DT) and generalized Chaplygin gas (GCG) models. In our best fit analysis for these dark energy proposals we consider flat and the non-flat priors. For all models, we find that relaxing the flatness condition implies that data favors a positive curvature; moreover, the GCG model is nearly flat, as required by Cosmic Microwave Background (CMB) observations. ",
        "introduction": "Various proposals have been put forward to explain recent observations indicating that the universe is accelerating. A possible explanation is that the universe is filled with dark energy in the form of an exotic component, the generalized Chaplygin gas, with negative equation of state \\cite{Kamenshchik:2001cp,Bilic:2001cg,Bento:2002ps}. The striking feature of this model is that it allows for an unification of dark energy and dark matter \\cite{Bento:2002ps}.  Another possible explanation for the accelerated expansion of the Universe could be the infrared modification of  gravity one should expect from  extra dimensional physics, which would lead to a modification of the effective Friedmann equation at late times. A concrete model has been suggested by Dvali, Gabadadze and Porrati  \\cite{Dvali} and later generalized by Dvali and Turner \\cite{Dvali:2003rk}. Another possibility is the modification of the Friedmann equation by the introduction of an additional nonlinear term proportional to $\\rho^n$, the so-called Cardassian model \\cite{Freese:2002sq}.   Currently type Ia supernovae (SNe Ia) observations provide the most direct way to probe the dark energy component at low to medium redshifts.  Recently, supernovae data has been analysed by various groups and it was shown that it yields relevant constraints on some cosmological parameters. In particular, it is possible to conclude that, when one considers the full supernova data set, the decelerating model is ruled out with a significant confidence level \\cite{Choudhury:2003tj}. It is also shown that one can measure the current value of the dark energy equation of state with higher accuracy and the data prefers the phantom kind of equation of state, $w_X < -1$. Furthermore, the most significant result of that analysis is that, without a flat prior,  supernovae data does not favour a flat $\\Lambda$CDM model at least up to $68\\%$ confidence level, which is consistent with other cosmological observations. In what concerns the equation of state of the dark energy component,  it has been shown in Ref. \\cite{Alam:2003}, using the same set of supernovae data, that  the best fit equation of state of dark energy evolves rapidly from $w_X \\simeq 0$ in the past to $w_X \\sim -1$ at present, which suggests that a time varying dark energy fits the data better than the  $\\Lambda$CDM model.  In this paper, we analyze the Cardassian, the DT and the GCG models using so called {\\it gold} sample SNe Ia compilation of data by  Riess {\\it et al.} \\cite{Riess:2004nr}; we consider both flat and non-flat priors. For more details see Ref. \\cite{Bento:2004ym}. ",
        "conclusions": "We performed a likelihood analysis of the latest type Ia supernovae data for three  models: the modified gravity Cardassian and Dvali-Turner models and the generalized Chaplygin gas model of unification of dark energy and dark matter. We find that SNe Ia most recent data allows, in all cases, for non-trivial constraints on model parameters as summarized in Table  \\ref{table:best}. We find that, for all models, relaxing the flatness condition implies that data favors a positive  curvature and the GCG model is nearly flat in this case. The fact that the  {\\it gold} sample of supernovae data prefers a flat GCG model, which is consistent with CMB observations, leads us to conclude that  the GCG is a better choice among the three  alternative models that we have considered."
    },
    "0512/astro-ph0512595_arXiv.txt": {
        "abstract": "We develop a simple physical model to describe the most common type of low-frequency quasi-periodic oscillations (QPOs) seen in a number of accreting black hole systems, as well as the shape of the relativistically broadened iron emission lines that often appear simultaneously in such sources. The model is based on an inclined ring of hot gas that orbits the black hole along geodesic trajectories. For spinning black holes, this ring will precess around the spin axis of the black hole at the Lense-Thirring (``frame-dragging'') frequency. Using a relativistic ray-tracing code, we calculate X-ray light curves and observed energy spectra as a function of the radius and tilt angle of the ring, the spin magnitude, and the inclination of the black hole. The model predicts higher-amplitude QPOs for systems with high inclinations, as seen in a growing number of black hole binary systems. We find that the {\\it Rossi X-ray Timing Explorer} observations of low-frequency QPOs in GRS 1915+105 are consistent with a ring of radius $R \\approx 10M$ orbiting a black hole with spin $a/M \\approx 0.5$ and inclination angle of $\\iBH\\approx 70^\\circ$. Finally, we describe how future X-ray missions may be able to use simultaneous timing and spectroscopic observations to measure the black hole spin and probe the inner-most regions of the accretion disk. ",
        "introduction": "Over the past decade, X-ray timing and spectroscopy observations of accreting black hole binaries have provided us with increasingly sensitive measurements of the inner-most regions of the accretion disk. With the proper theoretical understanding of these observations, we should be able to probe the behavior of matter and energy in the strongest known gravitational fields. Recent observations by \\citet{mille05} with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) provide time-varying measurements of spectroscopic features for a black hole on sub-second time scales. They find a strong correlation between the phase of the light curve oscillation and the width of the iron K$\\alpha$ emission line.  We propose a simple model to explain this correlation and make a number of additional predictions for the spectral features and light curves of similar black hole binary systems. The quasi-periodic oscillations (QPOs) considered in this paper are so-called ``type C'' QPOs, following the naming convention described in \\citet{remil02}. These appear to be the most common type of QPO in black hole binaries; they are seen in the spectrally hardest states, are stronger and more coherent than types A and B, and are observed over a large range in frequency (typically between $\\sim 0.1$ Hz and $\\sim 10$ Hz), even within a single source. It is believed that this type of black hole QPO is related to the 10-60 Hz QPOs observed in many neutron star low-mass X-ray binaries \\citep{wijna99,casel05}.  \\citet{stell98} suggested for these 10--60 Hz neutron star QPOs that their frequencies might correspond to the Lense-Thirring frequency of test particles around a spinning neutron star. A later version of their model was also applied to low-frequency QPOs in black hole X-ray binaries \\citep{stell99}. In this work we expand on that model, using a fully relativistic ray-tracing model, and for the first time compute quantitative X-ray light curves and emission line spectra that may be compared more directly with observations. This is similar to the ray-tracing analysis described in \\citet{schni04} for the hot spot model of high frequency QPOs. However, that hot spot model, which was developed primarily to explain the 3:2 ratio in the frequencies of QPO peaks, requires that the emission comes from a specific radius in the disk where the coordinate frequencies have small integer ratios \\citep{abram01,abram03}. The precessing ring model presented in this paper requires no such special radius, consistent with observations, where the LFQPOs are seen to drift in frequency (and thus radius), while the pairs of high frequency QPOs are always seen at the same frequency within a given source \\citep{mccli05}.  One of the primary objections to the interpretation of QPOs in the context of gravitomagnetic frame-dragging is based on the Bardeen-Petterson effect \\citep{barde75}, which predicts that the inner regions ($r \\lesssim 100 M$) of a viscous accretion disk should align perpendicular to the black hole spin axis. If the disk is constrained to the black hole equatorial plane, then the Lense-Thirring precession will not produce a time-varying light curve, as the observed orientation of the disk does not change. However, subsequent theoretical studies of the inner accretion disk have led to a number of different scenarios in which the emitting gas could have a significant inclination relative to the black hole spin axis. These perturbations could be driven by trapped pressure gradients \\citep{kato91,ipser96}, radiation warping \\citep{iping90,pring96,malon96,malon97}, or gravitomagnetic forces \\citep{marko98}. \\citet{marko98} argue that the radiation warping modes will be strongly damped, while the lowest-order gravitomagnetic modes, corresponding to the Lense-Thirring precession frequency, will be weakly damped and should exhibit quality factors of $Q \\sim 2-50$, in agreement with observations.  In the context of the supermassive black hole in Sgr A$^\\ast$, \\citet{rocke05} have recently shown that for hot, magnetized disks with low Mach-numbers, internal pressure gradients can sustain large-scale, coherent precession of the entire inner-disk region. This picture of a hot, geometrically thick region in the inner disk is consistent with some of the earliest models of black hole accretion disks \\citep{thorn75}, as well as more recent general relativistic magneto-hydrodynamic simulations that predict a local density maximum, or ``inner torus'' just outside the inner-most stable circular orbit \\citep{devil03a,devil03b}. Additionally, the fact that most black hole QPOs appear predominately in the power-law dominant spectral states and usually show greater amplitudes at higher energies further suggests that they are originating from a hot, non-thermal region of the accretion flow. For the purposes of this paper, we are not concerned with the exact physical mechanism that produces the precessing ring, but rather we are interested in the observational manifestation of such a ring, assuming that it exists.  In Section \\ref{model} we describe in greater detail the features of the precessing ring model and the global disk geometry. We also show the relation between black hole spin and ring radius for a given QPO frequency. Section \\ref{results} presents the results of our ray-tracing calculation, giving light curves and iron line widths for a range of model parameters. In Section \\ref{data} we compare these results to a number of black hole systems with type C QPOs, in particular the GRS 1915+105 observations of \\citet{mille05}. In Section \\ref{discussion} we conclude with a discussion of how this model may be tested further with future X-ray missions and observations of active galactic nuclei (AGN). ",
        "conclusions": "\\label{discussion} We have developed a geodesic precessing ring model to help explain the properties of type C low-frequency QPOs in black hole binaries. In this model, the surrounding accretion disk may be slightly misaligned  with the black hole spin axis (as would generally be the case for compact binaries), or the inner-most region of the disk may get excited into an inclined orbit, forming a ring of hot gas precessing around the black hole spin axis at the Lense-Thirring frequency. For a given QPO frequency, we have shown the dependence of the ring's radius on the black hole spin, which gives a lower limit on the spin parameter, assuming the ring must be outside of the ISCO. Using a relativistic ray-tracing code, we have produced X-ray light curves and time-varying iron line emission spectra. Similar to the hot spot model described in \\citet{schni04}, the ring model predicts the QPO amplitudes should increase with increasing inclination, in agreement with observations. Another important feature predicted by this model is the existence of a third, intermediate peak in the broad iron line spectrum at the QPO phase corresponding to maximum intensity, formed by the gravitational lensing of emission from the far side of the black hole.  In the states where most of the type C QPOs are observed, the emission of black hole X-ray binaries in the {\\it RXTE} band is usually dominated by two spectral components: a soft thermal component that probably comes from the accretion disk and a hard power-law component. This hard component has often been associated with the inverse-Compton scattering of thermal photons through a corona of hot electrons, and more recently also with emission from the base of a jet \\citep{marko2003}. However, the existence of QPOs in the radio-faint observations of GRS 1915+105 suggests that they do not originate in a jet \\citep{mille05}. Type C QPOs, and in fact many types of QPOs in X-ray  binaries, have often been associated with this hard spectral component---not only because the fractional amplitude of the QPOs increases with energy, but also because the QPOs are detected in energy bands where the contribution from the disk is negligible. In view of these facts, it seems clear that while the disk may determine the seed frequencies, it is likely that some other geometry is also involved in producing the non-thermal component and the QPO behavior. For example, it has long been proposed that the inner-most section of a thin accretion disk may be very hot, optically thin, and thus radiatively inefficient \\citep{thorn75}. We assume this medium of hot gas is closely aligned with the inner edge of the accretion disk, which in our model takes the form of a geodesically precessing ring.  The GRS 1915+105 observations of \\citet{mille05} can be explained well with the precessing ring model, matching the observed amplitudes of the fundamental and harmonic peaks of the QPO power spectrum. The rms variations of $13-15\\%$, coupled with an independently measured inclination of $\\iBH = i_{\\rm bin} = 66^\\circ$ \\citep{fende99}, suggest a tilt angle of $\\Delta\\theta \\approx 15-20^\\circ$. The spectral line measurements suggest a moderate to high spin with $a/M \\gtrsim 0.5$ so that $\\dE \\lesssim 1.5$ keV, but the {\\it RXTE} spectral data does not have enough energy resolution to rule out lower values of $a/M$. In short, our ring model agrees quite well with the light curve modulations, giving good estimates for the overall geometry of the system, yet the poor spectral data still cannot sufficiently constrain the black hole spin $a/M$.  Until recently, only time-averaged iron emission line profiles, and lines drawn from flux windows during aperiodic variability, had been studied with moderate and high resolution spectrometers. The recently launched {\\it Suzaku} and future missions such as {\\it Constellation-X} and {\\it XEUS} promise high spectral resolution along with high effective area.  This is exactly the combination required to rigorously test the predictions of our precessing ring QPO model.  {\\it Suzaku} X-ray Imaging Spectrometer (XIS) observations of GRS~1915$+$105 for $\\sim$100~ksec in states similar to those considered by \\citet{mille05} should achieve the same flux sensitivity as obtained with {\\it RXTE} but with a spectral resolution ten times higher ($\\sim$100~eV versus $\\sim$1~keV at 6~keV).  This may be sufficient to clearly detect a third peak in the iron emission line profile using the same QPO-phase-resolved technique. Depending on their final configurations, {\\it Constellation-X} and {\\it XEUS} may be able to achieve the same flux sensitivity as the \\citet{mille05} result in only a few ksec (assuming similar QPO amplitudes and frequencies), with an energy resolution of a few eV at 6~keV.  With such observatories, a third peak in the iron emission line profile, if present, should be clearly resolved. Indeed, with {\\it Constellation-X} and {\\it XEUS}, it may be possible to detect a modulation of the iron line flux and profile in as few as 100 QPO cycles, which would enable detailed measurements of the inner disk ring evolution with time.  While this paper has focused primarily on stellar-mass black holes, and a single source in particular, the ring model could easily be applied to AGN observations as well. In fact, the much longer time scales associated with AGN would greatly improve our spectral resolution by using ``slower'' observatories like {\\it Chandra} and {\\it XMM-Newton}. Since the spin of an AGN is thought to come primarily from accretion, it is likely the accretion disk and black hole spin axes are closely aligned, potentially giving a small tilt $\\Delta\\theta$ and thus little variation in the light curve. Also, the massive accretion disk around an AGN is almost certainly optically thick, decreasing the relative effects of gravitational lensing (photons cannot pass through the disk and form Einstein rings). However, due to the larger signal-to-noise ratio for AGN timing observations, no QPOs have yet been unambiguously identified in AGN, making a direct comparison to galactic black holes more difficult.  \\vspace{0.25cm}\\noindent We thank Chris Reynolds for helpful discussions and Cole Miller for his extensive and constructive comments to an early version of the manuscript. JDS is grateful to Ed Bertschinger for his continued insights and encouragement. Support comes from NASA grant NAG5-13306."
    },
    "0512/astro-ph0512548_arXiv.txt": {
        "abstract": "Collisionless shocks are commonly argued to be the sites of cosmic ray (CR) acceleration.  We study the influence of CRs on weakly magnetized relativistic collisionless shocks and apply our results to external shocks in gamma-ray burst (GRB) afterglows.  The common view is that the transverse Weibel instability (TWI) generates a small-scale magnetic field that facilitates collisional coupling and thermalization in the shock transition.  The TWI field is expected to decay rapidly, over a finite number of proton plasma skin depths from the transition.  However, the synchrotron emission in GRB afterglows suggests that a strong and persistent magnetic field is present in the plasma that crosses the shock; the origin of this field is a key open question.  Here we suggest that the common picture involving TWI demands revision.  Namely, the CRs drive turbulence in the shock upstream on scales much larger than the skin depth.  This turbulence generates a large-scale magnetic field that quenches TWI and produces a magnetized shock. The new field efficiently confines CRs and enhances the acceleration efficiency. The CRs modify the shocks in GRB afterglows at least while they remain relativistic. The origin of the magnetic field that gives rise to the synchrotron emission is plausibly in the CR-driven turbulence.  We do not expect  ultrahigh energy cosmic ray production in external GRB shocks. ",
        "introduction": "\\setcounter{footnote}{0}  Collisionless shocks are observed on all astrophysical scales. The diffusive shock acceleration (DSA) mechanism is believed to accelerate cosmic rays (CRs) in these shocks \\citep{Bell:78,Blandford:78,Blandford:87}.  The CRs can carry a substantial fraction of the energy of the shock and thus the CR pressure can influence the structure of the shock \\citep{Eichler:79,Blandford:80,Drury:81,Ellison:84}. This picture has recently received support from  X-ray observations of supernova (SN) remnants \\citep{Warren:05}. Recently, \\citet{Bell:04,Bell:05} has shown that CRs in SNe can drive turbulence and amplify magnetic fields in the shock upstream, and \\citet{Dar:05} speculate that relativistic jets could do the same.   The current lore is that weakly magnetized relativistic collisionless shocks are mediated by the transverse Weibel instability (TWI; \\citealt{Weibel:59,Fried:59}). TWI produces a magnetic field near equipartition that provides collisional coupling in the shock transition layer \\citep{GruzinovWaxman:99,Medvedev:99}. TWI is universally observed in two- and three-dimensional particle-in-cell (PIC) simulations of unmagnetized colliding shells (e.g.,~\\citealt{Lee:73,Gruzinov:01,Silva:03, Frederiksen:04,Jaroschek:04,Nishikawa:05a,Medvedev:05,Kato:05}). The resulting field is small-scale, of the order of the proton plasma skin depth $\\lambda_{\\rm s}$.  The role of TWI in particle acceleration and in the downstream thermodynamics is controversial. The field is expected to decay rapidly, within a few $\\lambda_{\\rm s}$ from the shock transition \\citep{Gruzinov:01,Milosavljevic:06}. This decay is evident in three dimensional PIC simulations \\citep{Spitkovsky:05}. Although there are claims that the decay saturates at distances of $\\sim 100\\lambda_{\\rm s}$ from the shock \\citep{Silva:03,Medvedev:05}, survival of the field over larger distances has not been demonstrated.  According to the popular model of gamma-ray bursts (GRBs; e.g., \\citealt[and references therein]{Piran:05a}), the afterglow originates in a relativistic blast wave that propagates into an ambient $e^-p$ plasma. The afterglow emission is ascribed to synchrotron radiation from nonthermal electrons that gyrate in the shock-generated magnetic field. Detailed studies of GRB spectra and light curves \\citep{Panaitescu:02,Yost:03} have shown that the magnetic energy density in the emitting region (the downstream shocked plasma) is a fraction $\\epsilon_B\\sim 10^{-2}$ to $10^{-3}$ of the internal energy density.\\footnote{Under some circumstances $\\epsilon_B$ as low as $\\sim10^{-5}$ can fit the data \\citep{Eichler:05}.} This field must persist over at least a few percent of the width of the blast wave \\citep{Rossi:03}. Its origin has remained a key open question.  Compressional amplification of the weak pre-existing magnetic field of the interstellar medium (ISM) yields $\\epsilon_B\\sim 10^{-9}$ \\citep{Gruzinov:01} and does not explain the field in the emitting region. If TWI does develop in the transition layer, it generates a strong field in the vicinity of the shock transition. However, as discussed above, there is no evidence that it can persist over the required $\\sim10^9\\lambda_{\\rm s}$ from the shock (e.g.,~\\citealt{Piran:05b} and references therein).    X-ray observations of GRB afterglows are modeled as an optically thin synchrotron spectrum requiring nonthermal electrons with Lorentz factors as high as $\\sim 10^6$. The observations indicate that electrons, and thus protons as well, are efficiently accelerated in the  shock to produce a hard power-law spectrum. DSA can achieve such Lorentz factors if the circum-burst medium is magnetized at $\\sim 1\\ \\mu\\textrm{G}$ level, as expected in the ISM. The weak magnetic field, however,  does not directly affect the structure of the shock (e.g., TWI can still develop in the shock transition).  It seems that a self-consistent picture of relativistic astrophysical shocks includes high energy particles (CRs) and at least a weakly magnetized upstream. Here we explore the interaction between these components. In \\S~\\ref{sec:amplification}, we derive the conditions under which the CRs drive turbulence in the upstream; this turbulence generates a large-scale magnetic field that increases the acceleration efficiency.\\footnote{\\citet{Gruzinov:01} and \\citet{Waxman:04} speculate that CRs could amplify large-scale magnetic fields in external GRB shocks.} In \\S~\\ref{sec:grb}, we argue that this mechanism is a candidate solution of the origin of the field inferred from the afterglow emission in external GRB shocks. In \\S~\\ref{sec:discussion}, we discuss implications for TWI and for the production of ultrahigh energy cosmic rays (UHECRs). ",
        "conclusions": "\\label{sec:conclusions}  CRs accelerated in relativistic collisionless shocks excite large-scales turbulence and magnetic field generation in the shock upstream.  The field generated in the shock precursor has power on scales much larger than the proton plasma skin depth.  The propagation of CRs in the generated field is diffusive.  In external GRB shocks, CRs are accelerated to higher energies in the generated field than in the pre-existing field of the ISM. The generated field reaches equipartition with the energy density in the fluid.  The shock transition is turbulent with a hydrodynamic profile dominated by CR pressure. The commonly invoked TWI is probably quenched in relativistic collisionless shocks by the magnetic field and the turbulence generated in the shock precursor.  External GRB shocks do not accelerate UHECRs. PIC simulations of collisionless shocks must include a weak upstream magnetic field and simulate a spatial domain as large as the CR gyroradius in the upstream to observe the acceleration of particles beyond equipartition."
    },
    "0512/astro-ph0512581_arXiv.txt": {
        "abstract": "We estimate the polarized thermal dust emission from MHD simulations of protostellar collapse and outflow formation in order to investigate alignment of outflows with magnetic fields. The polarization maps indicate that alignment of an outflow with the magnetic field depends on the field strength inside the cloud core; the direction of the outflow, projected on the plane of the sky, is aligned preferentially with the mean polarization vector for a cloud core with a magnetic field strength of 80~$\\mu$G, while it does not tend to be aligned for 50~$\\mu$G as long as the 1000~AU scale is considered.  The direction of the magnetic field at the cloud center is probed by the direction of the outflow.  In addition, the magnetic field at the cloud center can be revealed by {\\it ALMA} even when the source is embedded deeply in the envelope.  The Chandrasekhar-Fermi formula is examined using the polarization maps, indicating that the field strength predicted by the formula should be corrected by a factor of $0.24 - 0.44$.  The correction factor has a tendency to be lower for a cloud core with a weaker magnetic field. ",
        "introduction": "Magnetic fields are believed to control not only protostellar collapse, but also formation of circumstellar disks and outflows.  Many observations have suggested that the outflow and jet axis of the young star is aligned preferentially along the cloud-scale magnetic field \\citep*[e.g.,][]{Cohen84,Strom86,Vrba86,Vrba88,Tamura89,Jones03}.  However, recent observations indicate a suggestion contrary to previous ones concerning the issue of alignment of outflows and jets with the magnetic fields.  High-resolution observations of submillimeter polarization have resolved the magnetic fields around young stars on a $\\sim 10^{3-4}$~AU scale, which is comparable to the outflow scale \\citep*{Momose01,Henning01,Wolf03,Vallee03}. \\citet{Wolf03} investigate alignment of outflows with magnetic fields for Bok globules associated with Class 0 protostars and Class I sources, suggesting that two Bok globules are associated with outflows parallel to the magnetic fields, while the two other globules are associated with outflows perpendicular to the magnetic fields.  For Classical T Tauri stars (CTTSs), \\citet{Menard04} estimate orientation of the symmetry axes of the disk-jet systems in the Taurus-Auriga region,  and indicate that CTTSs are oriented randomly with respect to the local magnetic field.   Recently, \\citet{Matsumoto04} (hereafter MT04) have performed MHD simulations of the collapse of magnetized cloud cores and reproduced outflow generation in order to investigate the directions of outflows, circumstellar disk, rotation, and magnetic fields.  The simulations show that the outflow tends to be aligned with a local magnetic field of a 10~AU scale irrespective of the magnetic field strength assumed, while the alignment depends on the field strength on the cloud core scale. A disk-outflow system is aligned with the cloud core scale magnetic field within $\\sim 5^\\circ$ and $\\sim 30^\\circ$ for the initial field strengths of 37.1 and $18.6\\,\\mu$G, respectively, because of the strong magnetic braking during the collapse.  When a weak field strength of $7.42\\,\\mu$G is assumed, the outflow is not aligned with the cloud core scale magnetic field. In this Letter, alignment of an outflow with magnetic field is discussed by constructing polarization maps from the MHD simulations of MT04. ",
        "conclusions": "\\label{sec:discussion}  Activity of an outflow may be discussed in terms of alignment of the outflow with the magnetic field. The cloud core with stronger magnetic field exhibits a faster outflow as shown in Figure 18 of MT04, and the outflow tends to be aligned with the polarization vector.  This indicates that the fast outflow is observed parallel to the magnetic field.  According to the observations toward CTTSs, \\citet{Menard04} suggest a similar tendency in spite of the different evolutionary stage from that considered here: CTTSs without bright and extended outflows have a tendency to be perpendicular to the magnetic field.   The outflows are extended up to only $150-200$~AU in the MHD simulation data used here, while the outflows observed by molecular line emission are extended up to a 1000~AU scale \\citep[e.g.,][]{Wolf03}. Moreover, the outflows presented here have considerably slower velocity than that observed around young stars. Therefore the cloud cores presented here are restricted to the very early evolutionary stage compared with the observed cloud cores. This restriction arises in response to computational cost for solving the launch mechanism of the outflow near the protostar. In further stages, the magnetic field hardly seems to affect the direction of the outflow on the scale of 1000~AU, because the outflow is accelerated at $r \\sim 10$~AU at a speed comparable to the Alfv\\'en velocity at this radius, and the magnetic field strength of the envelope decreases rapidly, proportional to $B \\propto r^{-1}$. Moreover, the protostar may be decoupled from the magnetic field of the envelope as a result of efficient ambipolar diffusion \\citep*{Nakano02} in further stages. The directions of the outflow will be fixed during the main accretion phase of the protostar."
    },
    "0512/astro-ph0512062_arXiv.txt": {
        "abstract": "We have developed an orbit-based method for constructing triaxial models of elliptical galaxies, which fit their observed surface brightness and kinematics. We have tested the method against analytical models with general distribution functions. Here, we present models that fit integral-field {\\tt SAURON} observations of NGC~3379 and NGC~4365. ",
        "introduction": "We have developed a method to construct triaxial models of elliptical galaxies, based on an extension of the Schwarzschild (1979) orbit superposition method. The models fit the observed surface brightness and (two-dimensional) stellar kinematics. The mass distribution and gravitational potential are three-dimensional and point-symmetric and hence include spherical, axisymmetric or triaxial geometries (van den Bosch et al., in prep). Extensive tests on axisymmetric and triaxial analytical models show that our method is able to recover general three-integral distribution functions (van de Ven et al., in prep). ",
        "conclusions": ""
    },
    "0512/astro-ph0512254_arXiv.txt": {
        "abstract": "We report spectroscopic observations of \\H B, the secondary in a visual binary in which the physically associated primary (separation $\\sim$19\\arcsec) is a \\ion{B9}{5} star. The secondary shows strong Li~$\\lambda$6708 absorption suggesting youth, and has attracted attention in the past as a candidate post-T Tauri star although this has subsequently been ruled out.  It was previously known to be a double-lined spectroscopic binary (F8+G6) with a period of 17.6 days, and to show velocity residuals indicating a more distant massive third companion with a period of at least 8 years.  Based on our radial velocity measurements covering more than two cycles of the outer orbit, along with other measurements, we derive an accurate triple orbital solution giving an outer period of $9.447 \\pm 0.017$ yr. The third object is more massive than either of the other two components of \\H B, but is not apparent in the spectra. We derive absolute visual magnitudes and effective temperatures for the three visible stars in \\H.  Isochrone fitting based on those properties gives an age of $200 \\pm 50$~Myr for the system. We infer also an inclination angle of $\\sim53\\fdg3$ for the inner orbit of \\H B. We detect a near-infrared excess in \\H B which we attribute to the third star being a close binary composed of late-type stars. This explains its large mass and lack of a visible signature. Modeling of this excess allows us to infer not only the masses of the components of the unseen companion, but also the inclination angle of the outer orbit ($\\sim$73\\arcdeg). The \\H\\ system is thus at least quintuple. ",
        "introduction": "\\label{sec:introduction}  \\H\\ (also known as $\\chi$ Tau, 59 Tau, HIP 20430, HR 1369, $\\alpha = 4^{\\rm h} 22^{\\rm m} 34\\fs94$, $\\delta = +25\\arcdeg 37\\arcmin 45\\farcs5$, J2000, $V = 5.395$, SpT = \\ion{B9}{5}) is the brighter component in a visual binary system (ADS 3161, STF 528) with an angular separation of about 19\\arcsec. The relative position of the companion, \\H B ($V = 8.423$, SpT = \\ion{G2}{5}), was first recorded by William Herschel in 1782 \\citep[see][]{Lewis:06} and has not changed much since, indicating the physical association between the stars. An investigation by \\cite{Murphy:69} included this and many other physical pairs composed of a B-type and a late-type component to establish the absolute magnitudes of the primaries by reference to the better known magnitudes of the secondaries. Since the B stars must be relatively young, the \\H\\ system was included also in a study by \\cite{Lindroos:86} that concluded that many of the secondaries in such pairs are post-T Tauri stars, and show other indicators of youth such as \\ion{Ca}{2} H and K emission, H$\\alpha$ emission, strong Li~$\\lambda$6708 absorption, strong X-rays, infrared excess, or a location in the H-R diagram above the Zero Age Main Sequence (ZAMS). The age of \\H A was estimated to be 123 Myr by comparison with stellar evolution models.  \\H B was reported to be a double-lined spectroscopic binary by \\cite{Martin:92}, and independently by \\cite{Pallavicini:92}. Both teams detected significant Li~$\\lambda$6708 absorption. The spectroscopic orbit of the binary with a period of 17.6 days and an eccentricity of 0.3 was first derived by \\cite{Tokovinin:99}, who pointed out also that the residuals clearly indicated the presence of a third, rather massive star in the system with a period of a few years. The elements of the double-lined orbit were published by \\cite{Tokovinin:01}, along with a preliminary solution for the outer orbit for which $\\sim$70\\% of the estimated 8-year cycle was covered. This fit allowed those authors to confirm that the mass of the third star is larger than either of the objects in the inner pair, despite there being no sign of it in their spectra.  We report here our own spectroscopic observations of \\H B giving full coverage of the outer orbit over more than two cycles. We model the properties of the system and find compelling evidence that the third star is itself a close binary, making the \\H\\ system at least quintuple. ",
        "conclusions": "\\label{sec:discussion}  Our modeling of the infrared excess of \\H B shows that the unseen and over-massive third star is well explained as an equal-mass binary composed of late-type stars of $M \\approx 0.70$~M$_{\\sun}$ (spectral type approximately K4).  As indicated above, the alignment of the major axes of the inner (Ba--Bb) and outer (Bab--Bc) orbits in this qquadruple system as evidenced by the nearly identical longitudes of periastron is a sign that dynamical interactions are at play.  Our estimates of the inclination angles of the two orbits ($i_{\\rm Bab}$ and $i_{\\rm Bab-c}$) provide some information on the relative inclination angle $\\phi$, which is of considerable dynamical importance. That angle is given by $\\cos\\phi = \\cos i_{\\rm Bab} \\cos i_{\\rm Bab-c} + \\sin i_{\\rm Bab} \\sin i_{\\rm Bab-c} \\cos(\\Omega_{\\rm Bab}-\\Omega_{\\rm Bab-c})$ \\citep[e.g.,][]{Fekel:81}.  The position angles of the nodes ($\\Omega_{\\rm Bab}$ and $\\Omega_{\\rm Bab-c}$) are unknown, so we can only set limits to $\\cos(\\Omega_{\\rm Bab}-\\Omega_{\\rm Bab-c})$ between $-1$ and +1, which leads to $i_{\\rm Bab-c} - i_{\\rm Bab} \\leq \\phi \\leq i_{\\rm Bab-c} + i_{\\rm Bab}$. A lower limit to $\\phi$ can thus be placed at $\\phi_{\\rm min} = 20\\arcdeg \\pm 6\\arcdeg$, which appears to exclude coplanarity.  With our estimates of $i_{\\rm Bab}$ and $i_{\\rm Bab-c}$ the total mass of \\H B is 3.6~M$_{\\sun}$. The angular semimajor axis of the wide orbit in this system is therefore $a = 83.6$~mas, corresponding to 6.86~AU. Given the eccentricity and orientation of the orbit, the angular separation can be as large as 90~mas at times. The combined brightness of the stars in Bc is expected to be approximately 3.3~mag fainter than Ba+Bb in $V$, which should make it feasible to resolve Bc, e.g., by the speckle interferometry technique in the visible on a 4-m class telescope. The brightness difference in the $K$ band is even more favorable ($\\Delta K \\approx 1.5$~mag)\\footnote{\\cite{Tokovinin:01} reported an attempt made in October 1997 by I.\\ I.\\ Balega and collaborators to resolve \\H B with the speckle technique in the $K$ band using the 6-m telescope at the Special Astrophysical Observatory in Zelenchuk, Russia. The companion was not resolved: the resolution limit was 90~mas, and the predicted separation according to our orbit was about 60~mas.}.  Perhaps one of the most intriguing properties of \\H B is the presence of strong Li~$\\lambda$6708 absorption, as reported originally by \\cite{Gahm:83}, suggesting the star is young. This motivated a number of studies to explore the possibility that it is a post-T Tauri star. The equivalent width of the Li line was found to be $156 \\pm 4$~m\\AA\\ by \\cite{Martin:92} and 152~m\\AA\\ by \\cite{Pallavicini:92}, although no account of the double-lined nature of the star was made in these measurements.  \\cite{Pallavicini:92} also reported that no \\ion{Ca}{2} H and K emission is seen in \\H B, and that the H$\\alpha$ line is in absorption, which is somewhat unusual for a pre-main sequence object. Neither of the two visual components have been detected in X-rays \\citep{Schmitt:93, Huelamo:00}. Other than the near-infrared excess discussed in the present work, which is adequately explained by binary nature for Bc, no additional excess at longer wavelengths has been seen in \\H B. Although the IRAS satellite detected flux at 12$\\mu$m (but not at 25$\\mu$m, 60$\\mu$m, or 100$\\mu$m), that flux is consistent with being of photospheric origin and does not suggest any substantial amount of dust \\citep{Wyatt:03}. More sensitive observations by these authors at 850$\\mu$m yielded no detection of \\H B. With this evidence and the fact that the age of the system as estimated from the B-type primary is more typical of ZAMS stars, most studies have concluded that \\H B is not a post-T Tauri star. We support this conclusion, and our age is in fact even older than previous estimates. The presence of a strong Li line is a necessary but not sufficient condition for youth, as pointed out by \\cite{Pallavicini:92}. In fact, this feature is sometimes strong in evolved RS~CVn binaries and in other old stars \\citep[e.g.,][and others]{Duncan:81, Pallavicini:87, doNascimento:03}, which is considered to be a result of the interplay between rotation and the properties of the convective envelopes in these stars, as described in the previous references. Thus, while typical of very young stars, the strength of the Li line in \\H B is not exceptional for older objects.  From the evidence presented here at least five components are known in the \\H\\ system.  The B-type star itself has been examined for duplicity by the lunar occultation technique \\citep{Meyer:95} but was found to be unresolved. We note, however, that the Hipparcos Catalog lists the star as a `suspected non-single' \\citep{ESA:97} based on the quality of the astrometric solution, although no convincing non-single star solution was found. The radial velocity measurements of \\H A are too few in number and too poor in quality to be of much help in this regard. Thus, it is still possible that additional components are present, making this a very interesting multiple system."
    },
    "0512/astro-ph0512124_arXiv.txt": {
        "abstract": "The  Dvali, Gabadadze and Porrati (DGP) model has a self-accelerating solution, the positive  branch, where the brane is asymptotically de Sitter. A de Sitter space-time can be seen as a boundary between quintessence-like behaviour and phantom-like behaviour. We show that in a 5D dilatonic bulk, where the dilaton has an exponential potential,  with an induced gravity term on the brane, whose matter content corresponds only to vacuum energy,  the positive branch solution  undergoes a phantom-like stage where it faces a curvature singularity in its infinite future. The singularity can be interpreted as the  ``big rip'' singularity pushed towards an infinite future cosmic time. The phantom-like behaviour on the brane occurs without violating the null energy condition. There is another solution, the negative branch, where the brane can undergo an early-epoch (transient) inflationary phase  induced by the dilaton field. ",
        "introduction": "The supernova Ia (SNIa) observations \\cite{SNIa} and the cosmic microwave background (CMB) anisotropy data \\cite{Spergel:2003cb} suggest  that the expansion of our universe seems to be accelerating. A possible explanation for this evolution is the usual vacuum energy represented by a cosmological constant providing a negative pressure \\cite{3b}. However, the observational value of $\\Lambda$ is about $120$ orders of magnitude smaller than that established from field theory methods \\cite{3b}.  So far, alternative phenomenological models have been proposed to describe the late-time acceleration of the universe \\cite{Copeland:2006wr}.  One approach is to consider an  effective dark energy component  in the energy momentum tensor. For example, this component can be described by a scalar field  as  in quintessence models \\cite{Wetterich:fm} or tachyonic models \\cite{Gibbons:2002md}. Dark energy  can also be described effectively by a perfect fluid with a linear equation of state or a more general barotropic equation of state like in Chaplygin gas models \\cite{Kamenshchik}.  Motivated initially by SNIa observations, phantom energy  models \\cite{Caldwell:2003vq,Caldwell:1999ew,Bouhmadi-Lopez:2004me,singularity} have also been proposed to account for the late-time acceleration of the universe. A recent analysis \\cite{Rapetti:2004aa} of the current equation of state for dark energy  based on CMB anisotropies, SNIa and X-ray galaxy cluster data concluded that the current equation of state is compatible with a phantom-like behaviour of dark energy; i.e. $w$, the ratio of  the pressure and the energy density of dark energy, could be less than $-1$. A similar conclusion is reached in \\cite{Percival:2007yw}; i.e. observational data do not seem incompatible with phantom-like behaviour of dark energy. In phantom energy models the null energy condition is not satisfied. Hence, the energy density is an increasing function of the scale factor in an expanding Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) universe. This may lead to the occurrence of a big rip singularity in the future evolution of a phantom energy  dominated universe \\cite{Caldwell:2003vq,Bouhmadi-Lopez:2004me,singularity}.  An alternative approach to account for the late-time acceleration of the universe is to consider a generalised Einstein theory of gravity like brane-world models (for  reviews, see Ref.~\\cite{review}) where the observable four-dimensional (4D) universe is a brane (hyper-surface) embedded in a higher-dimensional space (bulk). For example, the Dvali, Gabadadze and Porrati (DGP) model \\cite{Dvali:2000hr}  has a self-accelerating solution at late-time which is asymptotically de Sitter \\cite{Deffayet,brane}. In this model the bulk is a 5D Minkowski space-time and the brane action contains an induced gravity term \\cite{Dvali:2000hr,Deffayet,brane,other1,lambdadgp,ktt,Papantonopoulos:2004bm,Maeda:2003ar,Bouhmadi-Lopez:2004ys,Bouhmadi-Lopez:2004ax} which is proportional to the 4D Ricci scalar curvature of the brane. Dilatonic brane-world models \\cite{Chamblin:1999ya,Maeda:2000wr,dilatonicbw,Feinstein:2001xs} can also account for late-time acceleration of the brane through a quintessence-like behaviour  driven by a bulk dilaton \\cite{Kunze:2001ji}. One aim of this paper is to show that a dilatonic brane-world model with an induced gravity term in the brane can mimic a phantom-like behaviour without including matter on the brane that violates the null  energy condition.   In an induced gravity brane-world model the Friedmann equation  has two solutions (depending on the embedding of the brane in the bulk \\cite{Deffayet}). One of these solutions (the self-accelerating solution or the positive branch) can account for the late-time evolution of the universe \\cite{Deffayet,brane}.  The other solution (the negative branch) can describe the early-time evolution of the universe\\footnote{For an alternative use of the negative branch of  the DGP scenario see Ref.~\\cite{lambdadgp}.} \\cite{ktt,Papantonopoulos:2004bm,Maeda:2003ar,Bouhmadi-Lopez:2004ys,Bouhmadi-Lopez:2004ax} and in particular corresponds to a correction to Randall-Sundrum (RS) model \\cite{RS}. In the dilatonic brane-world model with induced gravity that we will analyse, the negative branch may undergo a transient inflationary epoch.  The layout of this paper is as follows. In section 2, we present our dilatonic brane-world model with induced gravity. The bulk scalar field potential is a Liouville potential. The matter content of the brane is coupled to the dilaton field. We deduce the modified Friedmann equation for both branches, the junction condition of the dilaton across the brane, which constrains the brane tension, and the energy balance on the brane. In section 3, we analytically derive the  solutions of a vacuum brane for both branches; i.e. whose matter content is described through the brane tension. We show that the brane tension has a phantom-like energy density behaviour on the positive branch in the sense that the brane tension grows as the brane expands. The brane hits a singularity in its future evolution which may be interpreted as a ``big rip'' singularity pushed towards an infinite cosmic time. The negative branch can undergo an early-epoch (transient) inflationary phase  induced by the dilaton field through the brane tension. Finally, we conclude and summarise in section 4. ",
        "conclusions": "In this paper we study the behaviour of a dilatonic brane-world model with an induced gravity term on the brane with a constant induced gravity parameter $\\alpha$. We assume a  $\\mathbb{Z}_2$-symmetry across the brane. The dilatonic potential is an exponential function of the bulk scalar field and the matter content of the brane is coupled to the dilaton field. We deduce the modified Friedmann equation for the positive and negative branch (which specifies the way the brane is embedded in the bulk), the junction condition for the scalar field across the brane and the energy balance on the brane. We obtain the vacuum solutions Eqs.~(\\ref{+solution1})-(\\ref{+solution}); i.e. the matter content of the brane is specified by the brane tension, for a FLRW brane.  In the vacuum positive branch, the brane tension is a growing function of the scale factor and, consequently, mimics the behaviour of a phantom energy component on the brane. This phantom-like behaviour is obtained without including a phantom fluid on the brane. In fact, the brane tension does not violate the null energy condition. The expansion of the brane is super-inflationary; i.e. the Hubble parameter is a growing function of the cosmic time. At high energy (small scale factors), the brane is asymptotically de Sitter  for $c>0$, where the parameter $c$ is proportional to the initial Hubble parameter. There is a another solution which is a  bouncing solution ($c=0$) around a minimum non-vanishing scale factor. The brane faces a curvature singularity in its future evolution, where the Hubble parameter, brane tension and scale factor diverge. The singularity happens in an infinite cosmic time. Therefore, the singularity can be interpreted as a ``big rip'' singularity pushed towards an infinite future cosmic time.  In the vacuum negative branch, the brane tension is a decreasing function of the scale factor. Unlike the positive branch, the branch is not super-inflating. However, it always undergoes an inflationary expansion (see Eq.~(\\ref{adotdot})). The inflationary expansion can be eternal ($k^2\\leq 3$) or transient ($k^2> 3$). For large values of the scale factor, the negative branch is asymptotically Minkowski.   It remains to be seen how the  results obtained in this paper are modified  when the brane has  matter contents in addition to  the brane tension. On the other hand, we have chosen a specific dilatonic bulk and consequently we do not know how general are the brane properties we have found. We leave these interesting questions for future work."
    },
    "0512/astro-ph0512638_arXiv.txt": {
        "abstract": "Arguments are presented in favor of the idea that the solar dynamo may operate not just at the bottom of the convection zone, i.e.\\ in the tachocline, but it may operate in a more distributed fashion in the entire convection zone. The near-surface shear layer is likely to play an important role in this scenario. ",
        "introduction": "The issue of the location of the solar dynamos has been discussed and reviewed in a number of recent papers. Over the past 25 years a general consensus has been developing to place the solar dynamo at the bottom of the convection zone or even beneath it in the overshoot layer. This location also coincides with the tachocline, where the latitudinal differential rotation in the convection zone turns into rigid rotation in the radiative interior. A number of arguments in favor and against both distributed and overshoot dynamos have been collected in Brandenburg (2005). Which of the two scenarios is more viable cannot yet be decided conclusively until more realistic turbulence simulations of the solar dynamo become available.  From a dynamo-theoretic point of view it appears rather difficult to produce $\\sim100\\kG$ fields that are required in the standard scenario of an overshoot dynamo (D'Silva \\& Choudhuri 1993, Sch\\\"ussler et al.\\ 1994, Caligari, Moreno-Insertis, \\& Sch\\\"ussler 1995). Looking at a mixing length model of the solar convection zone, the equipartition field strength at the bottom of the convection zone is less than $1\\kG$, so the dynamo would need to produce a field in excess of a hundred times the equipartition value; see \\Tab{SolarModel}, where we have used data from stellar envelope models of Spruit (1974). Also, the idea of flux tubes ascending without disrupting through 20 pressure scale heights all the way from the bottom of the convection zones to the top seems nearly impossible.  \\begin{table*}[t!]\\caption{ Solar mixing length model of Spruit (1974). The equipartition field strength obeys $B_{\\rm eq}^2/4\\pi=\\rho u_{\\rm rms}^2$. }\\vspace{5pt}\\centerline{\\begin{tabular}{ccccccc} $\\quad z\\,\\mbox{[Mm]}\\quad$ & $\\quad H_p\\,\\mbox{[Mm]}\\quad$ & $\\quad u_{\\rm rms}\\,\\mbox{[m/s]}\\quad$ & $\\quad\\tau\\,\\mbox{[d]}\\quad$ & $\\quad\\nu_{\\rm t}\\,\\mbox{[cm$^2$/s]}\\quad$ & $\\quad2\\Omega_0\\tau$ & $B_{\\rm eq}\\,[\\G]$ \\\\ \\hline 24 &  8 & 70 & 1.3 & $1.5\\times10^{12}$ & 0.6 & 1600 \\\\ 39 & 13 & 56 & 2.8 & $2.0\\times10^{12}$ & 1.3 & 2000 \\\\ 155 & 48 & 25 &  22 & $3.2\\times10^{12}$ & 10  & 3100 \\\\ 198 & 56 &  4 & 157 & $0.6\\times10^{12}$ & 70  &  650 \\\\ \\label{SolarModel}\\end{tabular}}\\end{table*}  By contrast, distributed dynamos operating in the entire convection zone would be expected to have sub-equipartition field strengths of around $300\\G$ for the mean field. An important ingredient is the presence of shear; recent simulations (Brandenburg 2005) indicated that not even helicity is essential for producing large scale fields. Occasionally, such simulations produce what looks like bi-polar regions. So, the typical picture of $\\Omega$-shaped loops tied to the bottom of the convection zone (Parker 1979) may not be quite accurate, and the whole sunspot phenomenon may be rather more shallow that suggested by the standard picture. Examples of synthetically produced magnetograms are shown in \\Fig{pmagnetogram}.  In the present scenario the peak fields that emerge at the surface are thought to be the result of local concentrations. According to work by Kitchatinov \\& Mazur (2000), sunspots are actually the result of an instability of the mean-field equations of radiation magnetohydrodynamics, possibly assisted by negative turbulent magnetic pressure effects (Kleeorin, Mond, \\& Rogachevskii 1996). These ideas are in some ways similar to the convective collapse of magnetic fibrils (Zwaan 1978, Spruit \\& Zweibel 1979).  The usual argument against dynamos working in the convection zone proper is that magnetic buoyancy would bring the field to the surface on too short a time scale (Moreno-Insertis 1983). Indeed, buoyant loss of magnetic fields were anticipated when the first compressible simulations of convective dynamo action came out (Nordlund et al.\\ 1992, Brandenburg et al.\\ 1996). The lack of evidence for buoyant loss of magnetic field was explained by the stronger effect of turbulent downward pumping. This idea has recently been studied in much more detail (Tobias et al.\\ 1998, 2001, Dorch \\& Nordlund 2001, Ossendrijver et al.\\ 2002, Ziegler \\& R\\\"udiger 2003).  \\begin{figure}[t!]\\begin{center} \\includegraphics[width=\\columnwidth]{pmagnetogram} \\end{center}\\caption[]{ Magnetograms of the radial field at the outer surface on the northern hemisphere at different times. Light shades correspond to field vectors pointing out of the domain, and dark shades correspond to vectors pointing into the domain. The elongated rings highlight the positions of bipolar regions. Note the clockwise tilt relative to the $y$ (or toroidal) direction, and the systematic sequence of polarities (white left and dark right) corresponding to $\\meanB_y>0$. [Adapted from Brandenburg (2005).] }\\label{pmagnetogram}\\end{figure}  \\begin{table*} \\caption{Summary of arguments for and against tachocline and distributed dynamos, some of which are discussed in the text. [Adapted from Brandenburg (2005).]} \\vspace{5pt} \\label{ArgumentsSummary} \\centerline{\\begin{tabular}{|l|l|l|} \\hline arguments & tachocline dynamos & distributed/near-surface dynamos\\\\ \\hline in favor & flux storage & negative surface shear yields equatorward migration\\\\ & turbulent distortions weak & correct phase relation \\\\ & correct butterfly diagram with mer.\\ circ.\\ & strong surface shear at latitudes where the spots are \\\\ & size of active regions naturally explained & $\\max(\\Omega)/2\\pi=473\\nHz$ agrees with $\\Omega(\\mbox{youngest spots})$ \\\\ & & active zones move with $\\Omega(0.95)$ \\\\ & & $11\\yr$ variation of $\\Omega$ seen in the outer $70\\Mm$ \\\\ & & even fully convective stars have dynamos \\\\ \\hline against & $100\\kG$ field hard to explain & strong turbulent distortions \\\\ & flux tube integrity during ascent & rapid buoyant losses \\\\ & too many flux belts in latitude & too many flux belts if dynamo only in shear layer\\\\ & maximum radial shear at the poles & not enough time for shear to act \\\\ & no radial shear where sunspots emerge & long term stability of active regions\\\\ & quadrupolar parity preferred & profile of $\\Omega(\\mbox{youngest})$ by $4\\nHz$ above $\\Omega(0.95)$\\\\ & wrong phase relation & possible anisotropies in supergranulation \\\\ & $1.3\\yr$ variation of $\\Omega$ at base of CZ & \\\\ & coherent mer.\\ circ.\\ pattern required & \\\\ \\hline \\end{tabular}} \\end{table*}  A more complete list of arguments both in favor and against distributed dynamos versus tachocline dynamos is given in \\Tab{ArgumentsSummary}. For a more complete discussion of the various points see Brandenburg (2005).  An important aspect that requires some appreciation is simply the fact that mean (toroidally averaged) fields close to equipartition strength can actually be produced. This is an important result because there is a long history of arguments about the very possibility of producing large scale magnetic fields by the famous $\\alpha$ effect, starting with the work of Vainshtein \\& Cattaneo (1992) and Cattaneo \\& Hughes (1996). Again, this is not the place to attempt reviewing the vast amount of literature that has emerged over the past few years. An excellent review has been given by Ossendrijver (2003). For yet more recent aspects see the review by Brandenburg \\& Subramanian (2005a).  \\begin{figure}[t!]\\begin{center} \\includegraphics[width=\\columnwidth]{pmean_comp} \\end{center}\\caption[]{ Evolution of the energies of the total field $\\bra{\\BB^2}$ and of the mean field $\\bra{\\meanBB^2}$, in units of $B_{\\rm eq}^2$, for runs with non-helical forcing and open or closed boundaries; see the solid and dotted lines, respectively. The inset shows a comparison of the ratio $\\bra{\\meanBB^2}/\\bra{\\BB^2}$ for nonhelical ($\\alpha=0$) and helical ($\\alpha>0$) runs. For the nonhelical case the run with closed boundaries is also shown (dotted line near $\\bra{\\meanBB^2}/\\bra{\\BB^2}\\approx0.07$). Note that saturation of the large scale field occurs on a dynamical time scale; the resistive time scale is given on the upper abscissa. [Adapted from Brandenburg (2005).] }\\label{pmean_comp}\\end{figure}  At the heart of the problem with the $\\alpha$ effect is the fact that this and a few other related effects produce large scale magnetic helicity. On the other hand, the total magnetic helicity obeys a conservation law. However, since the {\\it total} magnetic helicity is the sum of large scale magnetic helicity and small scale helicity, the production large scale magnetic helicity of one sign must imply the production of a similar amount of small scale helicity of the opposite sign. It is this small scale helicity of the opposite sign that acts are to quench and suppress the original $\\alpha$ effect (Pouquet, Frisch, \\& L\\'eorat 1976). In the absence of magnetic helicity fluxes, this leads to a resistively controlled slow-down toward the final saturation of the dynamo (Brandenburg 2001). This behavior is now well reproduced in the framework of the dynamical quenching model (Field \\& Blackman 2002, Blackman \\& Brandenburg 2002, Subramanian 2002).  A possible way out of this was suggested first by Blackman \\& Field (2000a,b) who proposed that small scale magnetic helicity could leave the sun through the surface so as to allow the dynamo to saturate unimpededly; see also Kleeorin et al.\\ (2000, 2002, 2003) for similar work on the galactic dynamo. However, this does not happen just automatically; what is required is an active driving of magnetic helicity flux within the domain toward the boundaries. One such flux was identified by Vishniac \\& Cho (2001). Their flux works only in the presence of shear; see Subramanian \\& Brandenburg (2004, 2005), and Brandenburg \\& Subramanian (2005b). Another important flux would be due to simple advection; see Shukurov et al.\\ (2005). The way the sun could dispose of its excess small scale magnetic helicity might be through coronal mass ejections (Blackman \\& Brandenburg 2003). \\FFig{pmean_comp} shows the dramatic difference between simulations with and without open boundaries. This simulation does have strong shear, which is important for driving the Vishniac \\& Cho (2001) flux. ",
        "conclusions": "In this short paper we have summarized just a few of the aspects that appear crucial in determining the location of the solar dynamo. As we have said in the beginning, a full account of these ideas is given in Brandenburg (2005), and have been reviewed in Brandenburg (2006). The main reason is that a distributed dynamo appears quite plausible, i.e.\\ previous problems have largely been ruled out. Furthermore, from a dynamo-theoretic viewpoint, dynamos operating only in a narrow shell at the bottom of the convection zone appear rather implausible. As far as observational evidence is concerned, one can say that the distributed dynamo scenario is at least not in conflict with observations. Moreover, as expected, the magnetic field drives cyclic variations of the toroidal flow speed (so-called torsional oscillations) with the 11 year cycle period (Howe et al.\\ 2000a, Vorontsov et al.\\ 2002). The amplitude of these flow variations decreases with depth, which is mainly due to the larger mass to be swung around at greater depth. However, if the dynamo really produced $100\\kG$ fields in the overshoot layer, one would eventually expect corresponding flow variations at that depth. Such variations may currently still be below the detection limit, but what is seen are variations with a typical period of around 1.3 year at the base of the convection zone (Howe et al.\\ 2000b).  Another aspect concerns the proper motion of sunspots: young sunspots are know to rotate faster than older ones (Tuominen 1962). This suggests that sunspots may be anchored in the layer where the angular velocity is maximum (Tuominen \\& Virtanen 1988, Balthasar, Sch\\\"ussler, \\& W\\\"ohl 1982, Nesme-Ribes, Ferreira, \\& Mein 1993, Pulkkinen \\& Tuominen 1998). The rotational velocity of very young sunspots (age less than 1.5 days) is $14.7^\\circ/{\\rm day}$ at low latitudes (Pulkkinen \\& Tuominen 1998), corresponding to $473\\nHz$, which is about the largest angular velocity measured with helioseismology anywhere in the sun. This corresponds to the helioseismologically determined angular velocity at a radius $r/R=0.95$, which is $35\\Mm$ below the surface. Similar conclusions can be drawn from the apparent angular velocity of old and new magnetic flux at different latitudes (Benevolenskaya et al.\\ 1999).  There is still a problem in understanding why the cycle period is 22 years, and not 3 years, which would be the natural frequency for distributed dynamos (K\\\"ohler 1973). This is in principle also a problem for overshoot dynamos and it is traditionally ``solved'' by postulating an overall decrease of the electromotive force. This is obviously not satisfactory. A plausible ``excuse'' for such an overall decrease of the electromotive force might be a partial alleviation of catastrophic quenching due to magnetic helicity fluxes, mediated by coronal mass ejections. However, at the moment there is no dynamo model taking seriously into account the magnetic helicity losses due to coronal mass ejections. However, this would be a major goal for future work."
    },
    "0512/hep-ph0512282_arXiv.txt": {
        "abstract": "A dynamical model for the dark energy is presented in which the ``quintessence'' field is the axion, $a_Z$, of a spontaneously broken global $U(1)_{A}^{(Z)}$ symmetry whose potential is induced by the instantons of a new gauge group $SU(2)_Z$. The $SU(2)_Z$ coupling becomes large at a scale $\\Lambda_Z \\sim 10^{-3}\\,eV$ starting from an initial value $M$ at high energy which is of the order of the Standard Model (SM) couplings at the same scale $M$. A perspective on a possible unification of $SU(2)_Z$ with the SM will be briefly discussed. We present a scenario in which $a_Z$ is trapped in a false vacuum characterized by an energy density $\\sim (10^{-3}\\,eV)^4$. The lifetime of this false vacuum is estimated to be extremely large. Other estimates relevant to the ``coincidence issue'' include the ages of the universe when the $a_Z$ potential became effective, when the acceleration ``began'' and when the energy density of the false vacuum became comparable to that of (baryonic and non-baryonic) matter. Other cosmological consequences include a possible candidate for the weakly interacting (WIMP) Cold Dark Matter as well as a scenario for leptogenesis. A brief discussion on possible laboratory detections of some of the particles contained in the model will also be presented. ",
        "introduction": "The nature of the dark energy (responsible for an accelerating universe \\cite{acceleration}) is one of the deepest problem in contemporary cosmology. Supernovae observations at redshifts $1.25 \\leq z \\leq 1.7$ when combined with cosmic microwave background (CMB) and cluster data gave an equation of state $w=p/\\rho = -1.02+0.13-0.19$ \\cite{riess} and are consistent with a generic $\\Lambda\\,CDM$ model where  $w=-1$ independently of $z$. Most recently, distance measurements of 71 high redshift Type Ia supernovae by the Supernova Legacy Survey (SNLS) up to $z=1$ combined with measurements of baryon acoustic oscillations by the Sloan Digital Sky Survey also fits a flat $\\Lambda\\,CDM$ with constant $w = -1.023 \\pm 0.090 \\pm 0.054$ \\cite{snls}. Future proposed measurements to test whether or not $w$ is time-varying will be of crucial importance. Various forms of Quintessence had been proposed to describe the present accelerating universe \\cite{sahni}. A generic feature of these models is the presence of a time-varying $w$. However, it is a known fact that the dark energy is subdominant at higher values of redshift which makes it much harder to detect the $z$-dependence of $w$ \\cite{doran}. Until this is resolved, it is practically impossible to distinguish the class of quintessence models with time-varying $w$ from one in which $w$ is practically constant and is equal to $-1$. However, one should keep in mind that several quintessence models typically predict $w>-1$ now with many of them having $w \\agt -0.8$. In fact, one can try to reconstruct the quintessence potential as had been done recently by \\cite{sahlen} whose analysis of recent data appeared to favor a cosmological constant.  Is there a quintessence scenario in which $w=-1$ for a large range of $z$ and which mimics the $\\Lambda\\,CDM$ model? Can such a scenario make predictions that go beyond the accelerating universe issue and that can be tested experimentally? These are the types of questions we wish to address in this paper.  There exists a well-known phenomenon that can be readily applied to the search for models that mimic $\\Lambda\\,CDM$: The idea of the false vacuum. It has been used in the construction of models of early inflation (although a ``standard model'' is yet to be found) \\cite{kolb}. In its simplest version, the potential of a scalar field (whose nature depends on a given model) develops two local minima: a ``false'' and a ``true'' one, as the temperature drops below a certain critical value. In this class of models, the universe is trapped in the false vacuum and the total energy density of the universe is soon dominated by the energy of this false vacuum, leading to an exponential expansion. For the early inflation case, models have been constructed to deal with the so-called graceful exit problem, i.e. how to go from the false vacuum to the true vacuum without creating gross inhomogeneities, resulting in a class of so-called new inflationary scenarios (see \\cite{kolb} for an extensive list of references).  Is the fact that present measurements appear to be consistent with a flat $\\Lambda\\,CDM$ model with a constant $w=-1$ an indication that we have been and are still living in a false vacuum with an energy density $\\rho_{vac} \\sim (10^{-3}\\,eV)^4$? If that is the case, when did we get trapped in that false vacuum and when are we getting out of it? And where does this false vacuum come from?   In this paper, we would like to explore the above possibility and present a model for the false vacuum scenario. First, we will postulate the existence of an unbroken gauge group $SU(2)_Z$ \\cite{su2} and show that, starting with a gauge coupling comparable in value to the Standard Model (SM) couplings at some high energy scale ($\\sim 10^{16}\\,GeV$), it becomes strongly interacting at a scale $\\sim 10^{-3}\\,eV$. This new gauge group $SU(2)_Z$ \\cite{zophos} can be seen to come from the breaking $E_6$ \\cite{bj} into $SU(2)_Z \\otimes SU(6)$, where $SU(6)$ can, as one possible scenario, first break down to $SU(3)_c \\otimes SU(3)_L \\otimes U(1)$ and then to $SU(3)_c \\otimes SU(2)_L \\otimes U(1)_Y$, the details of which will be dealt with in a separate paper \\cite{hung2}. Next, we will list the particle content of our model and present an argument showing how the $SU(2)_Z$ instanton-induced axion potential can provide a model for the aforementioned false vacuum with the desired energy density \\cite{jain}. We will compute the transition rate to the true vacuum and show that it is plausible that the universe was trapped in this false vacuum and will be accelerating for a very, very long time. We then show that the particle spectrum of the model contains fermions which have the necessary characteristics of being candidates for a WIMP Cold Dark Matter. Finally, we will briefly discuss the possibility of SM leptogenesis in our model where the SM lepton number violation comes from the asymmetry in the decay of a ``messenger'' scalar field which carries both $SU(2)_Z$ and $SU(2)_L$ quantum numbers. A more detailed version of this leptogenesis scenario will appear in a separate paper \\cite{hung3}. We will end with a brief discussion of the possibility of detection for the messenger field and the $SU(2)_Z$ fermions (CDM candidates). ",
        "conclusions": "We have presented a model involving a new unbroken gauge group $SU(2)_Z$ which becomes strongly interacting at a scale $\\Lambda_Z \\sim 10^{-3}\\,eV$, starting with a value for the gauge coupling, at a high scale $\\sim 10^{16} \\,GeV$, which is close to that of a typical SM coupling at a similar scale. This similarity in gauge couplings at high energies is suggestive of a unification between $SU(2)_Z$ and the SM. A possible scenario for such a unification is briefly discussed here.  There are several cosmological implications of the $SU(2)_Z$ model. The most important one is a quintessence model for dark energy in which the quintessence field is the Peccei-Quinn-like axion $a_Z$ whose potential is induced by the $SU(2)_Z$ instantons. Unlike other quintessence models, our scenario involves the existence of a false vacuum where the $SU(2)_Z$ axion is trapped as the $SU(2)_Z$ plasma is cooled to the temperature $T \\sim \\Lambda_Z$. This occured when the age of the universe is $t_z \\approx 125 \\pm 14 Myr$ (at redshift $z \\sim 25$). The age when the acceleration began was computed to be $t_a \\approx 7.2 \\pm 0.8\\,Gyr$ (redshift $z \\sim 0.67$). The energy density of the false vacuum started to dominate the (baryonic and non-baryonic) matter density at around $t_{eq} = 9.5 \\pm 1.1 \\,Gyr$ (redshift $z \\sim 0.33$). Since the universe is trapped in the false vacuum, the equation of state $w(a_Z) \\approx -1$. This means that the quintessence scenario presented here {\\em effectively mimics} the flat $\\Lambda\\,CDM$ model! The most recent supernovae results (up to redshift $z=1$) when combined with those from the Sloan Digital Sky Survey fits a flat $\\Lambda\\,CDM$ model with $w \\approx -1$.  There are two other cosmological consequences of our model: 1) The $SU(2)_Z$ fermions $\\psi^{(Z)}_{1,2}$ as candidates of Weakly Interacting Massive Particles (WIMP) cold dark matter; 2) The decay of the messenger scalar field $\\tilde{\\bm{\\varphi}}_{1}^{(Z)}$ into $\\psi^{(Z)}_{1,2}$ plus a SM lepton generating a SM lepton asymmetry which transmogrifies into a baryon asymmetry through the electroweak sphaleron process. For (1), we showed that, with the masses of $\\psi^{(Z)}_{1,2}$ of $\\mathcal{O}(100\\,GeV)$, not only one obtains the initial (high energy) value of the $SU(2)_Z$ gauge coupling to be close in value to those of the SM couplings at a similar scale, one also finds that, when the temperature drops below their masses, the annihilation cross section is typically of the size of a weak cross section which is what is usually required in order for the relic abundances of these particles to be of the order of the ``observed'' CDM abundance. For (2), we showed that the interference between the tree-level and one-loop decay rates of $\\tilde{\\bm{\\varphi}}_{1}^{(Z)}$ into $\\psi^{(Z)}_{1,2}$ plus a SM lepton gives rise to a non-vanishing SM lepton asymmetry, which can be subsequently transformed into a baryon asymmetry. We then showed that, in order for this to happen, $\\tilde{\\bm{\\varphi}}_{1}^{(Z)}$ has to be lighter than $\\sim 1\\,TeV$. Since $\\tilde{\\bm{\\varphi}}_{1}^{(Z)} = (2,3)$ under $SU(2)_L \\otimes SU(2)_Z$, this mass constraint opens up the possibility of detecting the messenger fields at the LHC (or other future colliders). The details of the leptogenesis scenario are presented in a companion article \\cite{hung3}.  Finally, we end the paper with a brief discussion of the detectability of the messenger scalar field as well as other processes involving the CDM candidates $\\psi^{(Z)}_{1,2}$. In particular, we showed that the production and subsequent decay of the messenger field shows characteristic signals in terms of the decay geometry as well as the length of the charged tracks. The possible detection of $\\psi^{(Z)}_{1,2}$ as CDM matter as well as its contribution to a process such as $\\mu-e$ conversion present interesting phenomenological challenges which are under investigation.  Note added: After this present paper was completed, I learned from James ({\\em bj}) Bjorken that an earlier paper by Larry Abbott \\cite{abbott} contained some ideas which are similar in spirit to those presented here. It would be interesting to see if one can apply our model to the idea of a ``compensating field'' presented in \\cite{abbott}."
    },
    "0512/astro-ph0512197_arXiv.txt": {
        "abstract": "This work deals with cosmological models driven by real scalar field, described by standard dynamics in generic spherical, flat, and hyperbolic geometries. We introduce a first-order formalism, which shows how to relate the potential that specifies the scalar field model to Hubble's parameter in a simple and direct manner. Extensions to tachyonic dynamics, and to two or more real scalar fields are also presented. ",
        "introduction": "The recent discovery of cosmic acceleration \\cite{a1,a2,a3} has deeply affected modern cosmology. One of the main mysteries of the cosmic acceleration is the so-called dark energy, which makes up a very significant portion of the total energy of the universe. The simplest way to deal with dark energy includes a cosmological constant, but we can also consider standard Friedmann-Robertson-Walker (FRW) models described by real scalar fields, as in quintessence \\cite{q1,q2,q3,q4}. For the interested reader we recommend Refs.~{\\cite{prr,pr,tc,li}} for some recent reviews on the subject.  In this Letter we focus on dark energy, that is, we turn attention to FRW models described by real scalar fields. Our investigations inspect Einstein's equation and the equation of motion for the scalar field in a very direct way. We consider models described by real scalar field in generic space-time, displaying the usual spheric, flat or hyperbolic spatial profile. However, we follow a very specific route, in which we use the potential of the scalar field to infer how the scale factor or, alternatively, Hubble's parameter evolves in time.  The power of the method that we develop in this Letter is related to an important simplification, which leads to models governed by scalar field potential of very specific form, depending on a new function, $W=W(\\phi),$ usually named superpotential when supersymmetry is present, which very much remind us of supergravity, although we do not deal with supersymmetry in the present work. As we show below, we directly relate the function $W(\\phi)$ with Hubble's parameter, and this leads to scenarios of current interest to cosmology, unveiling a new route to investigate the subject. Almost all calculations are done as directly as possible, and we illustrate the main results with examples of current interest to modern cosmology.  Models for FRW cosmology with a single real scalar field are described by the standard action \\be\\label{model} S=\\int\\,d^4x\\;{\\sqrt{-g}\\;\\left(-\\frac14\\,R+{\\cal L(\\phi,\\partial_\\mu\\phi)}\\right)} \\ee where $\\phi$ describes a real scalar field and we are using ${4\\pi G}=1.$ The line element is $ds^2=dt^2-a^2(t)d{\\vec r}^{\\,2}$, and $a(t)$ is the scale factor. In general, the energy-momentum tensor is given by $T^\\mu_{\\;\\;\\nu}=(\\rho,-p,-p,-p),$ where $\\rho$ and $p$ represent energy density and pressure. We use Einstein's equation to get \\bes\\label{fe} \\ben H^2&=&2\\rho/3-k/a^2 \\\\ {\\ddot a}/{a}&=&-(\\rho+3p)/3 \\een \\ees where $k$ is constant: $k=1,0,$ or $-1,$ for spherical, flat, or hyperbolic geometry, respectively.  The equation of motion for the scalar field depends on ${\\cal L}(\\phi,\\partial_\\mu\\phi),$ which has the standard form \\ben\\label{sm} {\\cal L}=\\frac12\\partial_\\mu\\phi\\partial^\\mu\\phi-V(\\phi) \\een The energy density and pressure are given by $\\rho=\\dot\\phi^2/2+V$ and $p=\\dot\\phi^2/2-V,$ and the equation of motion for the scalar field has the form \\be \\ddot\\phi+3H\\dot\\phi+V_{\\phi}=0\\label{em1} \\ee where $H=\\dot a/a$ stands for Hubble's parameter, and $V_{\\phi}$ represents $dV/d\\phi.$ ",
        "conclusions": ""
    },
    "0512/astro-ph0512474_arXiv.txt": {
        "abstract": "\\noindent Various formation channels for the puzzling ultra-compact dwarf galaxies (UCDs) have been proposed in the last few years. To better judge on some of the competing scenarios, we present spectroscopic [Fe/H] estimates for a sample of 26 compact objects in the central region of the Fornax cluster, covering the magnitude range of UCDs and bright globular clusters ($18<V<21$ or $-13.4<M_V<-10.4$ mag). We find a break in the metallicity distribution of compact objects at $M_V\\simeq -11$ mag ($\\simeq$ 3$\\times 10^6$ M$_{\\sun}$): for $M_V<-11$ mag the mean metallicity is [Fe/H]$=-$0.62 $\\pm$ 0.05 dex, 0.56 $\\pm$ 0.15 dex higher than the value of $-$1.18 $\\pm$ 0.15 dex found for $M_V>-11$ mag. This metallicity break is accompanied by a change in the size-luminosity relation for compact objects, as deduced from HST-imaging: for $M_V<-11$ mag, $r_h$ scales with luminosity, while for $M_V>-11$ mag, $r_h$ is almost luminosity-independent. In our study we therefore assume a limiting absolute magnitude of $M_V=-11$ mag between UCDs and globular clusters. The mean metallicity of five Fornax dE,N nuclei included in our study is about 0.8 dex lower than that of the UCDs, a difference significant at the 4.5$\\sigma$ level. This difference is marginally higher than expected from a comparison of their  $(V-I)$ colors, indicating that UCDs are younger than or at most coeval to dE,N nuclei. Because of the large metallicity discrepancy between UCDs and nuclei we disfavor the hypothesis that most of the Fornax UCDs are the remnant nuclei of tidally stripped dE,Ns. Our metallicity estimates for UCDs are closer to but slightly below those derived for young massive clusters (YMCs) of comparable masses. We therefore favor a scenario where most UCDs in Fornax are successors of merged YMCs produced in the course of violent galaxy-galaxy mergers. It is noted that in contrast to that, the properties of Virgo UCDs are more consistent with the stripping scenario, suggesting that {\\it different} UCD formation channels may dominate in either cluster. ",
        "introduction": "\\subsection{Discovery of Ultra-compact Dwarf Galaxies} \\label{intro} \\noindent In their spectroscopic studies of the Fornax region, Hilker et al. \\cite{Hilker99} and later Drinkwater et al. (\\cite{Drinkw00} and \\cite{Drinkw03}) reported on the discovery of six isolated compact stellar systems in the Fornax cluster, having $-13.4<M_V<-12$ mag and half-light radii $r_h$ between 15 and 30 pc. These objects had been interpreted as galactic stars in previous imaging surveys because of their star-like appearance on photographic plates. Due to their compactness at luminosities comparable to average dwarf galaxies, they were dubbed ``ultra-compact dwarf galaxies`` (UCDs) (Phillips et al.~\\cite{Philli01}). High resolution imaging and spectroscopy (Drinkwater et al.~\\cite{Drinkw03}) place the Fornax UCDs between the sequence of globular clusters (GCs) and dwarf elliptical galaxies (dEs) in the fundamental plane of stellar systems. Their M/L ratios are in the range 2-4, at the upper end of values obtained for globular clusters. The luminosities and sizes of UCDs are comparable to those of bright nuclei of nucleated dwarf ellipticals (dE,Ns) (Lotz et al.~\\cite{Lotz04}).  It is not yet clear how objects as compact and luminous as the UCDs may have formed. One much discussed hypothesis is the stripping scenario (Bekki et al.~\\cite{Bekki03}), which proposes that UCDs are remnant nuclei of dwarf galaxies stripped in the potential well of their host cluster. It has also been proposed that most UCDs are simply very bright globular clusters (Mieske et al.~\\cite{Mieske02}), although we note that the sizes of the brightest Fornax UCDs are much larger than those of globular clusters (Drinkwater et al.~\\cite{Drinkw03}, Jord\\'an et al.~\\cite{Jordan05a}). UCDs could also be explained by stellar super-clusters created in gas-rich galaxy galaxy mergers (Fellhauer \\& Kroupa~\\cite{Fellha02} and~\\cite{Fellha05}). Related to the latter scenario, Bastian et al.~\\cite{Bastia05a} propose that young massive clusters (YMCs) created in recent galaxy mergers may resemble UCD properties after a Hubble time. Maraston et al.~\\cite{Marast04} show that the young massive cluster W3 in the merger remnant galaxy NGC 7252 will resemble UCD fundamental plane properties if aged to several Gyrs.  Recently, UCDs have also been detected in other environments: UCD candidates brighter than those in Fornax have been found in the massive galaxy cluster Abell 1689 (Mieske et al.~\\cite{Mieske05a}), consistent with the idea that tidal forces play a major role in the creation of UCDs. Hasegan et al. \\cite{Hasega05} report on the discovery of several UCDs in the Virgo cluster, some of which have high M/L ratios between 6 and 9. Those authors suggest a high M/L ratio as a criterion to tell UCDs from globular clusters. Also Jones et al.~\\cite{Jones05} have discovered 9 UCD candidates in Virgo, in a survey analogous to the one performed by Drinkwater et al.~\\cite{Drinkw00} in Fornax. Like the investigation of Hasegan et al., they find that UCDs in Virgo are close to the blue peak of the GC color distribution, in contrast to Fornax where they coincide with the red peak (Mieske et al.~\\cite{Mieske04a}). The data of Jones et al.~\\cite{Jones05} give a $(b_j - r)$ colour difference of 0.33 $\\pm$ 0.08 mag between Fornax and Virgo UCDs, significant at the 4$\\sigma$ level. This color difference hints at possible different formation scenarios, which will be discussed in the course of this paper. De Propris et al.~\\cite{deprop05} show from HST imaging that the surface brightness profile of five Fornax UCDs is only matched by the brightest dE,N nuclei in Virgo, while the UCDs are significantly more extended than fainter nuclei. That is not necessarily inconsistent with the stripping scenario, because UCDs have luminosities comparable to the brightest nuclei. Fainter UCDs might have different scale sizes. Furthermore, stripping processes can alter the structural parameters of embedded nuclei (Bekki et al.~\\cite{Bekki03}). Therefore, a direct comparison of surface brightness profiles might be a biased way of judging on differences between UCDs and embedded nuclei.  A note regarding nomenclature: the term ``ultra-compact dwarf galaxies'' (Phillipps et al.~\\cite{Philli01}) should be understood as a purely {\\it morphological} description, referring to astronomical objects of luminosities comparable to average dwarf galaxies and sizes much smaller than those. This term does not imply a specific UCD {\\it formation scenario}, be it as remnant of more extended dwarf galaxies or as merged super-clusters. Furthermore, in order to call a source UCD, it should clearly be distinguishable from the bulk of ``normal'' GCs, based on properties like mass, size or stellar populations. In our studies, whenever a separation between GCs and UCDs is not clear, we therefore use the term ``compact objects''. \\subsection{The Fornax Compact Object Survey} \\noindent In the Fornax Compact Objects Survey FCOS (Mieske et al.~\\cite{Mieske02}~and~\\cite{Mieske04a}) we have analyzed the photometric and kinematic properties of compact objects in the Fornax cluster. Our sample of compact objects was defined via radial velocity cluster membership, spanning a range of $18<V<21$ ($-13.4<M_V<-10.4$) mag. At the bright end, these are the UCDs, while at the faint end the sample is dominated by GCs.  A main finding of the FCOS was that compact objects with $V<20$ mag have somewhat different kinematics, spatial distribution and photometric properties than compact objects with $V>20$ mag. Furthermore, the latter sample was more consistent with the GC system of NGC 1399 than the former one. This was interpreted such that UCDs dominate the sample for $V<20$ mag, while GCs dominate for $V>20$ mag. Specifically, UCDs (i.e. compact objects with $V<20$ mag) are found to be redder by about 0.10 mag than the overall GC population of NGC 1399. Their colors are shifted about 0.2 mag red-wards of the well known color magnitude relation for dEs (Hilker et al.~\\cite{Hilker03}, Karick et al.~\\cite{Karick03}) and consistent with a color magnitude trend of comparable slope. The nuclei of present-day dE,Ns are bluer than UCDs by ~0.10 to 0.15 mag (Lotz et al.~\\cite{Lotz04}), suggesting that UCDs either have higher metallicities or older integrated stellar populations than the present-day nuclei. The magnitude distribution of compact objects observed in the FCOS shows a soft transition between UCDs and GCs. There is a 2$\\sigma$ overpopulation with respect to the extrapolated bright end of NGC 1399's GC luminosity function.  The findings obtained from the FCOS are consistent both with the stripping scenario (Bekki et al.~\\cite{Bekki03}) and the stellar super-cluster scenario (Fellhauer \\& Kroupa~\\cite{Fellha02}) as a source of UCDs. The major uncertainty of the data interpretation for the FCOS is the origin of the color difference between Fornax UCDs, the nuclei of dE,Ns and the main body of GCs. The redder color of UCDs as compared to nuclei either means older integrated ages or higher metallicities. Within the scenario of Bekki et al.~\\cite{Bekki03}, the M/L ratio of the stripping remnant (nucleus + remainder of envelope) is lower by a factor of two compared to the progenitor M/L, while the nucleus itself loses about 20\\% of its mass. This mass loss and the associated weakening of the gravitational potential will lead to expulsion of a substantial gas fraction out of the stripping remnant, a process conceivably supported by ram-pressure stripping. This would lower the efficiency of or completely halt subsequent star formation events in the naked nucleus. Therefore one would expect the UCDs to have older integrated ages and lower or at most equal metallicity as compared to dE,N nuclei (see also Jones et al.~\\cite{Jones05}). One would {\\it not} expect in the stripping scenario that the UCDs are on average more metal-rich than still existing dE,N nuclei. \\subsection{Aim of this paper} \\noindent The aim of this paper is to lift the age-metallicity degeneracy inherent in the broad-band color measurements for Fornax UCDs, GCs and dE,N nuclei, in order to better assess the validity of the various UCD formation scenarios. To this end we present new spectroscopic data for compact objects in Fornax, obtained with the 6.5m Magellan telescope at Las Campanas Observatory, Chile. The paper is structured as follows: in Sect.~\\ref{data} we present the data and their reduction. Sect.~\\ref{cmrel} re-assesses the color-magnitude relation for Fornax compact objects proposed in Mieske et al.~\\cite{Mieske04a}.  In Sect.~\\ref{results} we present the results of the spectroscopic metallicity measurements for UCDs, GCs and dE,N nuclei. The results are discussed in Sect.~\\ref{discussion}. We finish this paper with the conclusions in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} \\noindent We have presented spectroscopic [Fe/H] estimates for a sample of 26 compact objects in the central Fornax cluster, spanning the luminosity range $18<V<21$ ($-13.4<M_V<-10.4$) mag. Our sample includes UCDs at the bright end and is dominated by GCs at the faint end. We furthermore present [Fe/H] measurements for the nuclear regions of five Fornax dE,Ns.  We draw two main results from our measurements:  A. The mean metallicity of bright compact objects ($M_V<-11$ mag) is [Fe/H]$=-$0.62 $\\pm$ 0.05 dex, which is 0.56 $\\pm$ 0.15 dex higher than the mean metallicity  of $-$1.18 $\\pm$ 0.15 dex for faint compact objects ($M_V>-11$ mag). The difference is significant at the 3.7$\\sigma$ level. This metallicity break at $M_V=-11$ mag ($\\simeq 3\\times 10^6 M_{\\sun}$) is accompanied by a change in the size-luminosity relation for compact objects, as deduced from ACS-imaging: for $M_V<-11$ mag, $r_h$ scales with luminosity, while for $M_V>-11$ mag, $r_h$ is luminosity-independent and of typical value for GCs. Due to these characteristic features in the metallicity and size distribution, we identify compact objects with $M_V<-11$ mag as UCDs and those with $M_V>-11$ mag as GCs. The metallicity difference between UCDs and GCs is consistent with their color difference in $(V-I)$ under the assumption of coeval populations.  B. The mean metallicity of the five investigated dE,N nuclear regions is [Fe/H]$=-$1.38 $\\pm$ 0.15 dex, which is 0.75 $\\pm$ 0.17 dex lower than that of the UCDs, a difference significant at 4.5$\\sigma$. From the balance of two counteracting observational biases, we expect the mean metallicity of the actual nuclei to be about 0.1-0.2 dex lower, increasing the disagreement with the UCDs to at least 0.8 dex. The metallicity difference between UCDs and nuclei is slightly higher but still marginally consistent with that expected from their color difference. This implies that UCDs are coeval with or slightly younger than the nuclei. We find our metallicity estimates for UCDs close to but slightly below those derived for young massive clusters (YMCs).  From these findings and incorporating additional knowledge from the literature, we draw the following conclusions regarding the origin of UCDs:  1. We disfavor the hypothesis that most Fornax UCDs are the remnant nuclei of tidally stripped dE,Ns, given the significantly higher metallicity of UCDs compared to nuclei.  2. Although the Fornax UCDs match the metallicity and color of the red, metal-rich peak of the GC distribution, they are unlikely to be single, super-massive GCs due to their larger average sizes. Their sizes are more consistent with those predicted for merged super-clusters.  3. We propose that most Fornax UCDs are successors of merged YMCs produced in earlier star formation bursts in the course of gas rich galaxy-galaxy mergers. In contrast, for the Virgo cluster we suggest that the stripping scenario is the more important UCD formation channel. Four facts favor this distinction of formation channels: A) The Virgo UCDs are - in contrast to the Fornax UCDs -  on average metal-poor and of comparable color to dE,N nuclei. B) Some Virgo UCDs possess high M/L ratios that may require dark matter, while Fornax UCDs have lower M/L ratios. C) The Fornax cluster has a higher current galaxy merger rate than Virgo, favoring the YMC scenario. D) The Virgo cluster has a higher mass and deeper potential well, favoring the tidal stripping scenario.\\\\  In the future, increasing the available sample of Fornax compact objects by a factor of two to three may allow us to check whether there is a {\\it continuous trend} in metallicity with luminosity, as suggested from the color distribution. This will give important clues on the detailled formation histories of the Fornax UCDs. Furthermore, one must increase the baseline for environmental comparisons of UCD properties by surveying galaxy clusters with a broad variety of intrinsic properties such as merger history, total mass and density."
    },
    "0512/astro-ph0512642_arXiv.txt": {
        "abstract": "Recent theoretical work has solidified the viability of the collisional runaway scenario in young dense star clusters for the formation of very massive stars (VMSs), which may be precursors to intermediate-mass black holes (IMBHs).  We present first results from a numerical study of the collisional runaway process in dense star clusters containing primordial binaries.  Stellar collisions during binary scattering encounters offer an alternate channel for runaway growth, somewhat independent of direct collisions between single stars.  We find that clusters with binary fractions $\\mathpunct{\\gtrsim}10\\%$ yield {\\em two} VMSs via collisional runaways, presenting the exotic possibility of forming IMBH--IMBH binaries in star clusters.  We discuss the implications for gravitational wave observations, and the impact on cluster structure. ",
        "introduction": "\\label{sec:intro}  Observations hinting at the possible existence of intermediate-mass black holes (IMBHs) have mounted in recent years.  Ultra-luminous X-ray sources---point X-ray sources with inferred luminosities $\\mathpunct{\\gtrsim} 10^{39}\\,{\\rm erg}\\,{\\rm s}^{-1}$---may be explained by sub-Eddington accretion onto BHs more massive than the maximum mass of $\\mathpunct{\\sim}10M_\\sun$ expected via core collapse in main sequence stars, although viable alternative explanations exist \\citep{2004IJMPD..13....1M}. Similarly, the cuspy core velocity dispersion profiles of the globular clusters M15 and G1 may also be explained by the dynamical influence of a central IMBH \\citep{2002AJ....124.3255V,2002AJ....124.3270G,2005ApJ...634.1093G}, although theoretical work suggests that the observations of M15 may be equally-well explained by a collection of compact stellar remnants in the cores of the clusters \\citep{2003ApJ...582L..21B}.  At least three distinct IMBH formation mechanisms have been discussed in the literature.  The first, and possibly simplest, is core collapse of a massive Pop III star \\citep{2001ApJ...551L..27M}. The very low metallicity of Pop III stars ($Z\\lesssim 10^{-5}Z_\\sun$) allows much larger main-sequence stars to form, limits mass loss during stellar evolution, and increases the fraction of mass retained in the final BH (for stars more massive than $\\mathpunct{\\sim}250M_\\sun$) \\citep{2001ApJ...550..372F,2004IJMPD..13....1M}.  The second is the successive merging of stellar mass BHs via dynamical interactions, which may occur in star clusters that do not reach deep core collapse before $\\mathpunct{\\sim}3\\,{\\rm Myr}$, when the most massive cluster stars have become BHs \\citep{1993Natur.364..421K,1993Natur.364..423S,2000ApJ...528L..17P,2002MNRAS.330..232M,oleary2006}.  The process is relatively inefficient---in terms of the amount of mass added to the growing BH per BH ejected from the cluster---requiring BH seeds $\\mathpunct{\\gtrsim}500 M_\\sun$ to create $10^3 M_\\sun$ IMBHs \\citep{2004ApJ...616..221G}, even when aided by the Kozai mechanism \\citep{2002MNRAS.330..232M} or gravitational wave losses during close approaches \\citep{gultekin2005}. Introducing a mass spectrum for the BHs decreases the required seed mass, although growth is still rare \\citep{oleary2006}. The third is the runaway merging of main-sequence stars via direct physical collisions to form a very massive star (VMS), which may then collapse to form an IMBH \\citep{1999A&A...348..117P,2001ApJ...562L..19E,2002ApJ...576..899P,2004ApJ...604..632G}. Recent work shows that runaway growth of a VMS occurs generically in clusters with deep core collapse times shorter than $\\mathpunct{\\sim}3\\,{\\rm Myr}$ \\citep{freitag2005b}.  With the exception of one simulation \\citep{2004Natur.428..724P}, all simulations of runaway collisional growth in clusters have ignored the effects of primordial binaries, which are known to exist in clusters in dynamically significant numbers \\citep{1992PASP..104..981H}.  Indeed, some numerical results suggest that the primordial binary fraction ($f_b$) may have to be nearly 100\\% to explain the currently observed binary fractions in cluster cores \\citep{2005MNRAS.358..572I}. Primordial binaries are an important piece of the runaway collisional growth puzzle, since they introduce two effects which may strongly affect the process. On the one hand, binaries generate energy via dynamical scattering interactions in cluster cores, supporting the core against deep collapse and limiting the maximum stellar density attainable, and hence limiting the direct stellar collision rate \\citep{2003gmbp.book.....H}.  On the other hand, stellar collisions are much more likely in dynamical interactions of binaries, since the interactions are typically resonant \\citep{1996MNRAS.281..830B,2004MNRAS.352....1F}. Since these two effects of primordial binaries act in opposite senses with respect to the collision rate, it is not clear {\\em a priori} how they affect the collisional runaway scenario.  Before appealing to numerical methods, however, one can gain insight into the effects of primordial binaries by considering the coagulation equation \\citep{1993ApJ...418..147L,2000Icar..143...74L,2001Icar..150..314M}---a simplification of which is presented in \\citet{freitag2005a}---which describes the evolution of a spectrum of masses due to mergers.  For growth to occur in a runaway fashion, the coagulation equation requires that the cross section for collisions with the runaway object scales sufficiently rapidly with its mass: $S_{\\rm coll} \\propto M^\\eta$, with $\\eta > 1$.  For single--single star collisions in star clusters in which the central velocity dispersion is less than the escape speed from the surface of a typical star (so that the cross section is dominated by gravitational focusing), this corresponds to the constraint $R \\propto M^\\alpha$ with $\\alpha > 0$ on the main-sequence mass--radius relationship, which is satisfied by main-sequence stars of any mass or metallicity. With some approximations, the coagulation equation analysis \\citep{1975MNRAS.173..729H,1983ApJ...268..319H,1993ApJ...415..631S,gultekin2005} can be applied to collisions occurring in binary scattering interactions.  From Fig.~4 of \\citet{gultekin2005}, the cross section for close approach distances of $r_{\\rm min}$ in binary--single scattering encounters scales as $(S_{\\rm coll}/\\pi a^2) (v_\\infty/v_c)^2 \\propto (r_{\\rm min}/a)^\\gamma$, with $0.3 \\lesssim \\gamma \\lesssim 1$, where $a$ is the binary semimajor axis, $v_\\infty$ is the relative velocity between the binary and single star at infinity, and $v_c$ is the critical velocity \\citep[see {eq.~[12]} of][]{gultekin2005}.  Using the radius of the runaway star for $r_{\\rm min}$, $R \\propto M^\\beta$ with $0.5 \\lesssim \\beta \\lesssim 1$ for the scaling of the mass--radius relation, and assuming that the binding energy of the binary is roughly preserved during the encounter so that $a \\propto M$, we find: \\begin{equation}\\label{eq:coag} S_{\\rm coll} \\propto M^{2 + \\gamma (\\beta - 1)}\\,, \\end{equation} with $0.3 \\lesssim \\gamma \\lesssim 1$ and $0.5 \\lesssim \\beta \\lesssim 1$.  The minimum of the exponent is $\\mathpunct{\\approx} 1.5$.  Thus according to the coagulation equation, collisions induced in binary--single scattering interactions should yield runaway growth of a VMS.  This result says nothing about the {\\em rate} of growth of the runaway object.  In other words, it is still not clear whether binary interactions will limit the cluster core density such that the runaway timescale is longer than the massive star main-sequence lifetime of $\\mathpunct{\\approx} 3\\,{\\rm Myr}$, in which case the process would be halted.  Assuming that the cluster core density reached is high enough for the runaway to proceed, it would appear that a single binary is sufficient for a binary interaction-induced runaway to occur. This is, of course, not the case, since binary scattering interactions tend to destroy binaries.  Thus we expect that for sufficiently low $f_b$, the runaway will be primarily mediated by single--single collisions.  For sufficiently large $f_b$, the runaway will be primarily mediated by binary--binary interactions (the analysis above assumes binary--single interactions---for binary--binary the value of $\\gamma$ will be smaller, still allowing a runaway by eq.~[\\ref{eq:coag}]).  For intermediate $f_b$ it is possible that a binary interaction induced-runaway could proceed until the core binary population is sufficiently depleted that the cluster's core collapses.  If the first runaway is far enough from the center of the cluster when the core collapses, it is possible that a second runaway will be formed during core collapse, mediated by single--single collisions.  The exotic possibility of forming two VMSs in a cluster, and thus two IMBHs, is a tantalizing one, with implications for gravitational wave observations and cluster dynamics.  In this letter we present first results from a study of the runaway collisional scenario for the formation of VMSs in young dense clusters with primordial binaries. We briefly describe our numerical method, present results showing the growth of two runaways, and discuss the implications. ",
        "conclusions": "\\label{sec:discussion}  Although the process is somewhat uncertain, it is likely that the VMSs formed in young clusters via collisional runaways will undergo core-collapse supernovae and become IMBHs on a timescale of $\\mathpunct{\\sim}1\\,{\\rm Myr}$ \\citep{freitag2005b}. Our results show no evidence of the VMSs merging prior to becoming IMBHs.  After their separate formation, the two IMBHs will quickly exchange into a common binary via dynamical interactions.  The IMBH--IMBH binary will then shrink via dynamical friction due to the cluster stars, to the point at which the stellar mass enclosed in the binary is comparable to the binary mass.  This occurs on a timescale $\\mathpunct{\\sim}t_r \\langle m \\rangle / M_{\\rm IMBH}$, where $t_r$ is the local relaxation time and $\\langle m \\rangle$ is the local average stellar mass.  Since $\\langle m \\rangle / M_{\\rm IMBH} \\lesssim 10^{-2}$, this timescale is likely to be $\\mathpunct{\\lesssim}10\\,{\\rm Myr}$.  The binary will then shrink via dynamical encounters with cluster stars, the rate of which is governed by loss-cone physics.  The timescale for the binary to shrink to the point at which it merges quickly via gravitational radiation energy loss is likely to be $\\mathpunct{\\lesssim}1\\,{\\rm Gyr}$ \\citep{2003ApJ...599.1129Y,2005ApJ...618..426M}.  Although IMBH--IMBH binaries do not merge in the LISA band---the gravitational wave frequency at merger is $\\mathpunct{\\sim}1\\,{\\rm Hz}$---they do represent bright sources that take at least $\\mathpunct{\\sim}10^6\\,{\\rm yr}$ to cross the LISA band. Their inspiral (chirp) signals should be easily detectable by LISA out to a few tens of Mpc. Thus the number of detectable IMBH binary sources may be quite large, since most clusters are probably born with $f_b \\gtrsim 0.1$, and any cluster with mass $\\mathpunct{\\gtrsim}10^6 M_\\sun$ and central relaxation time $\\mathpunct{\\lesssim}20\\,{\\rm Myr}$ will lead to a double runaway.  Significant core rotation (with rotational speed comparable to the local velocity dispersion) is observed in the clusters M15, $\\omega$ Cen, 47 Tuc, and G1 \\citep[e.g.][]{2002AJ....124.3255V,2005ApJ...634.1093G}. This rotation suggests the presence of a core angular momentum source, such as an IMBH binary \\citep{2005MNRAS.364.1315M}. Similarly, observations of a millisecond pulsar in the halo of NGC 6752, and two others in the core with high negative spin derivatives, hint at the existence of an IMBH binary in the core \\citep{2003ApJ...599.1260C}. Since the IMBH--IMBH binaries formed via collisional runaways will merge within $\\mathpunct{\\sim}1\\,{\\rm Gyr}$ after formation, any angular momentum imparted to the cluster by the IMBH--IMBH binary will quickly diffuse out of the core on a core relaxation time.  An alternate mechanism must be at work in creating the core rotation seen in some globular clusters today."
    },
    "0512/astro-ph0512368_arXiv.txt": {
        "abstract": "{In this paper we derive a generic expression, which is valid for scales larger than Hubble radius and contains only the local terms, for the second order curvature perturbations for more than one field, provided the expansion is sourced by the energy density of a single field. As an application, motivated by our previous paper~\\cite{asko1}, we apply our formalism to two fields during preheating, where the inflaton oscillations are sourced by $\\lambda\\varphi^4$ potential which is governing the expansion of the Universe. A second field $\\sigma$, coupled to the inflaton through $g^2\\varphi^2\\sigma^2$, is excited from the vacuum fluctuations. The excited modes of $\\sigma$ amplify the super-Hubble isocurvature perturbations, which seed the second order curvature perturbations to give rise to a significantly large non-Gaussianity. Our results show that within $3$ inflaton oscillations for a range of parameters, $1< g^2/\\lambda < 3$, the non-Gaussianity parameter becomes: $f_{NL}\\geq {\\cal O}(1000)$, which is already ruled out by the current WMAP observation.}  \\preprint{NORDITA-2005-79} ",
        "introduction": "Inflation stretches the fluctuations outside the Hubble radius which later on enter inside the Hubble radius to seed the structure formation~\\cite{Liddle-Lyth}.  A single field inflationary model generates Gaussian fluctuations which is in an excellent agreement with the current observations~\\cite{WMAP}. These Gaussian fluctuations are due to the adiabatic spectrum of the curvature perturbations~\\cite{Mukhanov}, for a review see~\\cite{Brandenberger}, which in a single field case is conserved after the wavelength of a perturbed mode becomes larger than the Hubble radius~\\footnote{Perturbations whose wavelength is larger than the Hubble radius are often abused in the literature as ``super-horizon perturbations'', instead they should be called as ``super-Hubble perturbations.  Note that due to inflation the perturbations are always well inside the causal horizon. We particularly thank Robert Brandenberger for stressing this point time and again.}.  For multiple fields there is an iso-curvature component which can source the curvature perturbations, i.e., as a result the curvature perturbations are no longer conserved on large scales.   Although negligible but a finite calculable departure from Gaussianity can be obtained within a single field inflationary model, with a self interacting inflaton potential.  However the measure for non-Gaussianity, defined as $f_{NL}$ (see the precise definition below), is bounded by the slow roll parameters~\\cite{Acquaviva,Maldacena}, i.e., $f_{NL}\\sim |\\eta-3\\epsilon|$, where the slow roll parameters, $\\epsilon,~\\eta \\ll 1$. This trend remains true even for multi field case, see for example ~\\cite{Enqvist,Lidsey,Lyth1}. Unless the slow roll conditions are violated during inflation, it is hard to imagine a reasonably large non-Gaussianity can be generated, which would really challenge the common lore that slow roll inflation produces mainly Gaussian fluctuations. For a review on non-gaussianity, see \\cite{Bartolo:2004if}.  The only exceptional moment is right at the time when inflation ends, when the slow roll parameters become large, i.e., $\\epsilon,~\\eta \\sim {\\cal O}(1)$, nevertheless this still leaves $f_{NL}\\sim {\\cal O}(1)$ \\cite{Lyth2}. Large $f_{NL}\\gg 1$, however, can be obtained during preheating after inflation as we illustrated in our previous paper~\\cite{asko1}~\\footnote{See also other examples considered in Refs.~\\cite{asko2,asko3} and \\cite{Finelli:2001db}.}.   Particularly a scalar preheating~\\cite{some} is a phase where after inflation the coherent oscillations of the inflaton, $\\varphi$, dominates the energy density of the Universe and another scalar field, $\\sigma$, coupled to the inflaton through, $g^2\\varphi^2\\sigma^2$, is excited from the vacuum fluctuations. In this respect preheating is a non-perturbative phenomena.  Although preheating does not lead to a complete thermalization~\\cite{averdi,averdi1}, nevertheless, in many models preheating can be just a possibility before the perturbative decay of the inflaton, for example see the last reference of~\\cite{some}, if there exists favorable conditions such as significantly large coupling.  In reality it is hard to predict exactly how the inflaton couples to the Standard Model degrees of freedom, due to lack of a model where the inflaton is not an absolute gauge singlet~\\cite{Enqvist-anu}~\\footnote{An exceptional case of assisted inflation~\\cite{assist} with gauge invariant supersymmetric flat directions~\\cite{Enqvist-anu} can possibly address this crisis~\\cite{asko-anu}.}, therefore, it becomes difficult to pin down thermalization time scales, etc., which can be tested through present cosmological experiments.  Nevertheless, very recently as we pointed out that preheating can give rise to a large non-Gaussianity which can lead to constraining certain parameter regions, such as a coupling constant~\\cite{asko1,asko2,asko3}, we regard this as a window of opportunity to test the initial stage of preheating.    In order to understand why preheating is the test bed for a large primordial non-Gaussianity. Let us first note that in general we require:  \\begin{itemize}  \\item{Significant interaction which leads to curved potential that breaks slow roll conditions.}   \\item{Some small amplitude isocurvature fluctuations which could be amplified later on via non-adiabatic evolution.}   \\end{itemize}   During preheating it is possible to have both the criteria fulfilled. The first one can be obtained mainly because during preheating the slow roll conditions are violated. The second condition is also satisfied because the $\\sigma$ field, which carries isocurvature fluctuations~\\footnote{The fluctuations in $\\sigma$ field will be scale invariant and the amplitude of the fluctuations in $\\sigma$ will be similar to that of the inflaton, i.e., $\\delta\\sigma\\sim H_{inf}/2\\pi$ (during inflation and after the first few inflaton oscillations). With these comparatively large initial amplitude and scale invariant fluctuations it is possible to amplify the fluctuations in $\\delta \\sigma_{k}$ on super Hubble scales.} , undergoes non-adiabatic evolution.  The quantum modes of $\\sigma$ field are excited because of the time dependence on the effective mass of $\\sigma$ field (due to coupling between $\\varphi$ and $\\sigma$ where $\\varphi$ is oscillating) changes sign, therefore violates adiabatic evolution of the vacuum.    Still a reader might wonder why amplifying $\\delta\\sigma_{k}$ leads to a large non-Gaussianity? After all non-Gaussianity is related to the second order curvature fluctuations. Two specific questions could be:  \\begin{itemize}  \\item{Can a sub-Hubble process influence the super-Hubble second order curvature perturbations?}   \\item{Can existing super-Hubble modes amplify themselves to influence the second order curvature perturbations?}  \\end{itemize}  Note that the second order perturbations take into account of the interactions in leading order of the perturbation theory, namely the quadratic combinations of the first order perturbations act as sources for the second order. In momentum space the sources become convolutions and thus couple different scales. To answer the first question, we should also keep in mind that inflation makes the {\\it causal horizon} exponentially large compared to the Hubble radius. Therefore causality allows various modes to mix, i.e., super-Hubble modes in second order can be influenced by either convoluting the first order sub-Hubble modes or the first order super-Hubble modes. In paper~\\cite{asko1} we studied preheating related to $m^2\\varphi^2$-potential for the inflaton where the resonance affects the super-Hubble modes in second order through sub-Hubble perturbations, i.e., it relates to the first question.   In this paper we provide a generic expression for the second order curvature perturbation which are fed by the super-Hubble fluctuations in the first order. As an example we consider preheating in $\\lambda\\varphi^4$-potential for the inflaton where the second order super-Hubble modes are affected.     Let us now highlight the role of isocurvature fluctuations. It is a well known fact that the isocurvature fluctuations source the adiabatic fluctuations and therefore the curvature perturbations.  The adiabatic perturbations cannot increase by themselves on super-Hubble scales without iso-curvature perturbations.  This is true in first order perturbations, in second order case there are additional sources which feed the second order metric and curvature perturbations, we shall see them below. Without having isocurvature fluctuations it is therefore not possible to amplify the first order and subsequently the second order curvature perturbations. However note that non-Gaussianity is a ratio between the second order curvature perturbation with respect to square of the first order curvature perturbations. If both first and second order curvature perturbations grow due to large isocurvature fluctuations then the ratio need not be large always.   Therefore in order to get large non-Gaussianity we need a setup where only the second order curvature perturbations grow while the first order curvature perturbations do not grow at all. Can this happen in reality?  Yes indeed, this can happen when the isocurvature fluctuations do not seed the first order perturbations at all, but it only seeds the second order perturbations. We achieved this modest goal in our first paper,~\\cite{asko1}, where the background motion of a $\\sigma$ field were absent, only the coherent oscillations of $\\varphi$ field were present. In which case the perturbations in $\\sigma$ decouple from the first order perturbations, but not in the second order and in higher order perturbation theory. Therefore exponentially large growth in $\\delta\\sigma_{k}$ can feed the second order curvature perturbations to give rise to a very large non-Gaussianity.   The aim of the present paper is to lay down the formalism for second order curvature perturbations, where there is a single field dominating the energy density of the Universe while other multi-fields can leave their imprints via fluctuations~\\footnote{In our previous papers~\\cite{asko1,asko2,asko3} we mainly studied the first order and second order metric perturbations, which need not be conserved on large scales. Here our formalism is more robust as we study the perturbations in terms of conserved quantities, $\\zeta^{(1)},~\\zeta^{(2)}$, for the definition see below.  In this paper we verify our earlier claim that preheating can really boost large primordial non-Gaussianity.}.  We also show that the second order curvature perturbation becomes {\\it free} from the {\\it non-local contributions} on super-Hubble scales when we neglect the gradient terms, which is true in a generic background (all the scalar fields can have non-vanishing VEVs), see the detailed discussion in~\\ref{NLT}.    As an application we will consider a massless preheating, where the $\\sigma$ field does not have a bare mass term and the inflaton potential is quartic, i.e., $V(\\varphi,\\sigma)\\sim \\varphi^4+g^2\\varphi^2\\sigma^2$. This potential has certain nice properties which we will highlight in the course of our discussion, i.e., the instability band for the $\\delta\\sigma_k$ mode depends on the ratio, $g^2/\\lambda$, cosmologically interesting solutions can be obtained for the first instability band, $1< g^2/\\lambda < 3$.  In order to obtain the results analytically, we assume that the background value of $\\sigma $ is vanishing. A non-vanishing case has to be studied separately with an aid of numerics~\\footnote{If $\\sigma$ field is also oscillating and contributing to the energy density of the Universe then again we would have to address the problem numerically.}.   A large non-Gaussianity also depends on the number of oscillations, if the number of oscillations grow the non-Gaussianity grows exponentially at an alarming rate, see the Plots.~[\\ref{lphi4f1},\\ref{lphi4f2}]. A simple explanation can be given as follows, fluctuations of the $\\sigma$ field grow as $\\sim e^{2\\mu x}$, where $\\mu\\sim 0.2$ is a physical parameter governing the production of $\\sigma$ particles ( for a detailed discussion, see below) and $x$ is time in dimensionless units which measures the period of oscillations, $T\\sim 7$, so that the time elapsed is the number of oscillations, $N\\sim 3$, times the period, i.e., $x=NT$. With these numbers the growth factor is amazingly large, \\begin{equation} e^{2\\mu x}\\sim e^{8.4}\\sim 4000\\,. \\end{equation}\\label{growth} Barring accidental cancellations between various source terms or dampening the initial amplitude for the isocurvature fluctuations this leads to a large non-Gaussianity.   However we should remind the readers that the phase of preheating cannot last forever. Besides the metric and curvature perturbations there is a backreaction on energy momentum components due to exponentially large particle creation, this inevitably stops preheating. Depending on the couplings the backreaction can stop preheating in dozens of oscillations. When the backreaction becomes important our assumption with a vanishing VEV of $\\sigma$ field also breaks down.   Unless the time evolution of $\\sigma$ field really brings down the non-Gaussianity parameter to the level that is consistent with the current observations, the model considered in our paper with $1< g^2/\\lambda<3$ is deemed to be ruled out. For other instability bands the effective mass for the $\\sigma $ field becomes heavy, i.e., $m_{\\sigma,~eff}\\geq H_{inf}$, which decreases the amplitude of the isocurvature fluctuations and leads to a smaller non-Gaussianity parameter. In this respect non-Gaussianity acts as a discriminatory and it should be considered as a useful observable to distinguish various physical situations.    We begin with the definitions, then we describe the first and second order curvature perturbations. We then discuss the background evolution where we obtain the final expression for the second order curvature perturbations for more than one field, provided the background evolution is governed by one field alone, for an example this could be the inflaton. We then turn our attention to a massless preheating and derive an expression for the non-Gaussianity parameter. In order to be self consistent, we added an appendix, where useful relationships have been derived explicitly. ",
        "conclusions": "Following our previous claim~\\cite{asko1}, we pointed out in this paper that preheating in general can produce very large primordial non-Gaussianity.  We substantiate our point by providing a formalism which can take into account of the fluctuations for more than one field, therefore accounting for the isocurvature fluctuations which can source the first and second order curvature perturbations. The final expression for the second order curvature perturbation is given by Eq.~(\\ref{solzeta2fin}).  We studied a particular case of massless preheating as an example, where the second order curvature perturbations are fed by the the isocurvature fluctuations but not the first order.  This gives rise to a large non-Gaussianity which already exceeds the currently observed limit. The non-Gaussianity parameter grows exponentially and the final number depends on the number of oscillations before the backreaction kicks in to halt preheating completely, see the Plots~[\\ref{lphi4f1},\\ref{lphi4f2}].  Based on our result we rule out completely massless preheating on the parameter range, $1<g^2/\\lambda < 3$.   Our results are encouraging and obviously demands thorough study of non-Gaussianity generated during preheating. One of the goal will be to setup a numerical code where more than one field can participate in the expansion of the Universe. This would definitely improve the non-Gaussianity bound set in the current paper."
    },
    "0512/astro-ph0512532_arXiv.txt": {
        "abstract": "We provide a description of the SNSPH code---a parallel 3-dimensional radiation hydrodynamics code implementing treecode gravity, smooth particle hydrodynamics, and flux-limited diffusion transport schemes.  We provide descriptions of the physics and parallelization techniques for this code.  We present performance results on a suite of code tests (both standard and new), showing the versatility of such a code, but focusing on what we believe are important aspects of modeling core-collapse supernovae. ",
        "introduction": "SNSPH is a particle-based code, using smooth particle hydrodynamics (SPH) to model the Euler equations and a flux-limited diffusion package to model radiation transport.  Its tree algorithm is designed for fast traversal on parallel systems (Warren \\& Salmon 1993,1995) and has been shown to scale well up to high processor number on a wide variety of computer architectures.  The increased synchronization required for the transport algorithm limits this scalability to roughly 512 processors on current parallel computers for the 1-10 million particle core-collapse calculations that are feasible with our current codes and current computing power.  SNSPH has now been used in several papers appearing in the literature studying a range of problems from stellar collapse (Fryer \\& Warren 2002,2004; Fryer 2004; Fryer \\& Kusenko 2005) to supernova explosions and supernova remnants (Hungerford et al. 2003, 2005; Young et al. 2005; Rockefeller et al. 2005a; Fryer et al. 2005) to models of massive binaries and winds from these binaries (Fryer \\& Heger 2005; Fryer, Rockefeller, \\& Young 2005) to models of both the gas and disk evolution in the Galactic Center (Rockefeller et al. 2004, 2005b, 2005c).  But none of these papers has provided a detailed description of the SNSPH code itself.  Here we describe these details: \\S 2 describes much of the physics implementation from gravity to hydrodynamics to radiation transport; and \\S 3 describes many of the numerical techniques used in this particle-based code, focusing on many of the computational techniques required to make such a scheme scalable.  In \\S 4, we show the results of SNSPH for a number of tests (both standard and new) used to confirm the validity of SNSPH for core-collapse simulations.  Astrophysics problems tend to be too complex for most standard code tests to completely confirm the validity of the code for that specific astrophysics problem.  We present a small suite of tests to show the broad range of physics that must be tested to model core-collapse.  These tests demonstrate the wide applicability of this SNSPH code, but further testing is required to truly test all the physics in the code.  We conclude with a list of strengths and weaknesses of SNSPH. ",
        "conclusions": "The SNSPH code is built upon the parallel tree code of Warren \\& Salmon (1993,1995).  This code has proven extremely portable and scalable (beyond 1000 processors) on a number of computer architectures from large supercomputers --- the Advanced Simulation and Computing Program (e.g. Red, Pink, Q) and National Energy Research Scientific Computing Center (e.g. Seaborg) --- to local Beowulf clusters --- e.g. Space Simulator ---.  SNSPH uses the Benz version (Benz 1989) of Smooth Particle Hydrodynamics to model the Euler equations. A number of equations of state have been implemented to close these equations, from ideal and perfect gas equations of state to planetary equations of state (Benz, Slattery, \\& Cameron 1986) to the equation of state cocktail used to model core-collapse (Herant et al. 1994). SNSPH also includes an explicit flux-limited diffusion scheme to moderate neutrino transport (Herant et al. 1994, Fryer \\& Warren 2002).  This transport scheme is even now being modified to model photon transport.  Although the focus of SNSPH has been to study stellar collapse (Fryer \\& Warren 2002,2004; Fryer 2004) it has also been adapted to model stellar explosions (Hungerford et al. 2003, 2005) and the Galactic center (Rockefeller et al. 2004, 2005) with a list of ongoing projects that spreads even further to planet formation and cosmology.  But, as with any computational technique, this code has both advantages and disadvantages.  We have presented a number of tests of this SNSPH code.  These test problems go beyond the standard ``shock'' tests (Sedov, Sod problem, shock tube) common in most test suites and include tests of the broad range of physics modelled in core-collapse.   Our tests mark just the beginning set of tests SNSPH must pass if it is to become a multi-purpose radiation hydrodynamics code.  Even so, we can already outline the strengths and weaknesses of the SNSPH technique.  The particle based scheme allows modellers to preserve many advantages of a Lagrangian scheme in turbulent flows where typical grid-based Lagrangian cells become tangled.  We have shown that the hydrodynamics scheme conserves both total energy and momentum.  Our scheme is easily modified to include a wide variety of equations of state and external forces and the transport scheme can be modified to model the transport of photons as well as neutrinos.  The advantages include: \\begin{itemize} \\item[1:] Lagrangian technique allows the resolution to follow the mass.  This is ideal for problems that cover a large range of spatial scales, but that focus the mass in one place (e.g. supernova explosions).  It is also ideal for problems where the area of interest moves with time  (e.g. in neutron star kicks and in binaries). \\item[2:] The Lagrangian scheme also avoids any numerical diffusion (either heat or elemental abundances) that plague any Euler scheme during the advection step.  Such artificial diffusion can lead to disastrous results in problems where the abundances must be tracked exactly (e.g. nucleosynthesis issues in supernova explosions) and abundance gradients are large.  The same is true for problems where the temperature gradient is large. \\item[3:] SPH conserves angular momentum and linear momentum at the same time. \\item[4:] The explicit flux-limited diffusion transport scheme scales well (for transport schemes). \\end{itemize} The disadvantages of the particle based method are few and can be summed up in three limitations with respect to core-collapse supernovae: \\begin{itemize} \\item[1:] SPH initial conditions tend to have small perturbations on the 5-10\\% level that will artificially seed convection.  Such perturbations are difficult to reduce completely and their effects must be understood when interpreting simulation results. \\item[2:] The artificial viscosity term added to model shocks does not work as well as Riemann methods when the shock moves along the grid. However, when the shock is not so well behaved, we believe (and propose a test for comparison) that SPH behaves as well as most grid codes.  The artificial viscosity also leads to both an artificial angular momentum and artificial heating term.  In shear flows, such numerical effects can lead to unrealistic heating.  We have shown (Fryer \\& Warren 2004) that this is not an issue for even our fast-rotating supernova cores. \\item[3:] Transport schemes are generally designed for grid-based codes. The flux-limited transport scheme in this code works at a basic level, but more sophisticated schemes have not been incorporated into SPH, and much more work must be done to prove that more sophisticated schemes can be added smoothly. \\end{itemize}  No code is ideal for all problems.  Without the development of a more sophisticated transport scheme for SPH (whether it be a direct discretization method such as $S_n$ or a monte-carlo technique), our SNSPH code will not be able to model the detailed neutrino transport many believe is necessary to get a final solution to the core-collapse supernova problem.  New transport schemes applicable for SPH are being pursued and the current dearth of schemes may not prove a long-term limitation.  There are many problems that are most easily solved with a Lagrangian technique: from supernova explosions to binary mergers to galaxy collisions.  SNSPH is ideally suited for these problems and has the versatility to adapt to these problems quickly.  With other problems, such as black hole accretion disks, the gas accretion during the formation of Jovian planets or the complex shock structure in the Galactic center, SNSPH has its advantages and disadvantages over grid-based codes.  SNSPH is a powerful tool in developing a solution to these outstanding astrophysical problems.  {\\bf Acknowledgments} This work was funded under the auspices of the U.S.\\ Dept.\\ of Energy, and supported by its contract W-7405-ENG-36 to Los Alamos National Laboratory, by a DOE SciDAC grant number DE-FC02-01ER41176 and by NASA Grant SWIF03-0047-0037.  The simulations were conducted on the Space Simulator and ASC PINK at Los Alamos National Laboratory."
    },
    "0512/astro-ph0512018_arXiv.txt": {
        "abstract": "Recently a lot of attention has been given to building dark energy models in which the equation-of-state parameter $w$ can cross the phantom divide $w=-1$. However, to our knowledge, these models with crossing the phantom divide only provide the possibility that $w$ can cross $-1$. They do not answer another question: {\\em why crossing phantom divide occurs recently?} Since in many existing models whose equation-of-state parameter can cross the phantom divide, $w$ undulates around $-1$ randomly,  {\\em why are we living in an epoch $w<-1$?} This can be regarded as the second cosmological coincidence problem. In this note, we propose a possible approach to alleviate this problem within a hybrid dark energy model.\\\\  \\noindent PACS numbers: 95.36.+x, 98.80.Cq, 98.80.-k ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512420.txt": {
        "abstract": "Conservative numerical schemes for general relativistic magnetohydrodynamics (GRMHD) require a method for transforming between ``conserved'' variables such as momentum and energy density and ``primitive'' variables such as rest-mass density, internal energy, and components of the four-velocity.  The forward transformation (primitive to conserved) has a closed-form solution, but the inverse transformation (conserved to primitive) requires the solution of a set of five nonlinear equations.  Here we discuss the mathematical properties of the inverse transformation and present six numerical methods for performing the inversion.  The first method solves the full set of five nonlinear equations directly using a Newton-Raphson scheme and a guess from the previous timestep.  The other methods reduce the five nonlinear equations to either  one or two nonlinear equations that are solved numerically.  Comparisons between the methods are made using a survey over phase space, a two-dimensional explosion problem, and a general relativistic MHD accretion disk simulation.  The run-time of the methods is also examined.  Code implementing the schemes is available for download on the web. ",
        "introduction": "It is commonly thought that many astrophysical systems contain relativistic  plasmas with a dynamically significant magnetic field. Examples include accreting black holes in black hole binaries, galactic nuclei, gamma-ray bursts, the cores of massive stars undergoing core collapse, isolated neutron stars, and neutron stars in binary systems.  As a result, there is currently considerable interest in numerical methods for integrating the equations of general relativistic magnetohydrodynamics (GRMHD).  Within the last few years about six GRMHD schemes have been deployed \\citep{kom05,koide,harm,dvh,fragile,duezetal,anton}.  Some of these authors \\citep{kom05,koide,harm,duezetal,anton} have adopted a conservative scheme.  This means that the integrated equations are of the form \\begin{equation} \\del_t \\bU(\\bP) = -\\del_i \\bF^i(\\bP) + \\mathbf{S}(\\bP) \\label{conservative-eq} \\end{equation} Here $\\bU$ is a vector of ``conserved'' variables, such as particle number density, or energy or momentum density in the coordinate frame, the $\\bF^i$ are the fluxes, and $\\mathbf{S}$ is a vector of source terms that do not involve derivatives of $\\bP$ and therefore do not affect the characteristic structure of the system. $\\bU$ is conserved in the sense that, if $\\mathbf{S} = 0$, the rate of change of the integral of $\\bU$ over the volume depends only on fluxes at the boundaries, by the divergence theorem.  The vector $\\bP$ is composed of ``primitive'' variables such as rest-mass density, internal energy density, velocity components, and magnetic field components.  The fluxes and conserved quantities depend on $\\bP$. Conservative numerical schemes advance $\\bU$, then depending on the order of the scheme, calculate $\\bP(\\bU)$ once or twice per timestep.  In nonrelativistic conservative MHD schemes the conserved quantities are trivially related to the primitive variables; both the forward transformation $\\bP \\rightarrow \\bU$ and the inverse transformation $\\bU \\rightarrow \\bP$ have a closed-form solution.  In GRMHD (or even SRMHD) $\\bU(\\bP)$ is a complicated, nonlinear relation.  The inverse transformation has no closed-form solution and must be performed numerically.  This numerical operation must be robust, accurate, and fast---it is at the heart of all conservative GRMHD schemes.  In this paper we investigate several schemes for the inversion $\\bP(\\bU)$ and test each in an axisymmetric simulation of accretion onto a rotating black hole.  \\S~\\ref{sec:definitions} covers definitions and notational matters.  \\S~\\ref{sec:inversion-schemes} describes five distinct formulations of the algebraic equations to be solved numerically, and \\S~\\ref{sec:numer-impl} describes how each method is implemented numerically.   \\S~\\ref{sec:tests} describes the performance of these methods in the context of a survey over a range of primitive variable values and in two typical applications. \\S~\\ref{sec:conclusion} summarizes our findings and contains a guide to our results for those wishing to make their own implementation.  We will assume throughout that $G = c = 1$.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusion}  We have outlined a compact derivation of the equations for calculating $\\mathbf{P}(\\mathbf{U})$.  This formulation suggests a variety of possibilities for performing the calculation numerically.  We have implemented a small subset of the possible methods, and compared these implementations through (1) a survey over a large subset of possible $\\bU$ values; (2) embedding the inverter in a SRMHD code and evolving a cylindrical explosion problem due to Komissarov; (3) embedding the inverter in a GRMHD code and evolving a turbulent, magnetized disk around a rotating black hole.  Several key points emerge from the comparison.  First, the implementation can be made accurate enough that variable inversion does not make a significant contribution to the code error budget.  Second, some implementations can be made fast enough that they occupy only a few percent of the total cycles used by our GRMHD code.  Since we are using a particularly simple GRMHD algorithm, this is likely an upper limit to the fractional cost of the inversion in all GRMHD schemes.  Third, the inversion routine originally used in HARM \\citep{harm}, called ``5D'' here, is particularly slow and inaccurate.  We recommend the ``2D'' scheme, described in \\S 3.2.  A version of this scheme is available for download on the web at {\\tt http://rainman.astro.uiuc.edu/codelib/codes/pvs.tgz}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0512/astro-ph0512108_arXiv.txt": {
        "abstract": "We present results from a campaign of multiple epoch echelle spectroscopy of relatively faint ($V = 9.5-13.5$ mag) red giants observed as potential astrometric grid stars for the Space Interferometry Mission (SIM PlanetQuest). Data are analyzed for 775 stars selected from the Grid Giant Star Survey spanning a wide range of effective temperatures ($T_{eff}$), gravities and metallicities. The spectra are used to determine these stellar parameters and to monitor radial velocity (RV) variability at the 100 m s$^{-1}$ level. The degree of RV variation measured for 489 stars observed two or more times is explored as a function of the inferred stellar parameters. The percentage of radial velocity unstable stars is found to be very high -- about 2/3 of our sample. It is found that the fraction of RV-stable red giants (at the 100 m s$^{-1}$ level) is higher among stars with $T_{eff} \\sim 4500$ K, corresponding to the calibration-independent range of infrared colors $0.59 < (J-K_s)_0 < 0.73$. A higher percentage of RV-stable stars is found if the additional constraints of surface gravity and metallicity ranges $2.3< \\log~g < 3.2$ and -0.5 $<$ [Fe/H] $<$ -0.1, respectively, are applied. Selection of stars based on only photometric values of effective temperature ($4300\\,K < T_{eff} < 4700\\, K$) is a simple and effective way to increase the fraction of RV-stable stars. The optimal selection of RV-stable stars, especially in the case when the Washington photometry is unavailable, can rely effectively on 2MASS colors constraint $0.59 < (J-K_s)_0 < 0.73$. These results have important ramifications for the use of giant stars as astrometric references for the SIM PlanetQuest. ",
        "introduction": "The Grid Giant Star Survey \\citep[GGSS;][]{patterson01} is a partially-filled, all-sky survey to identify giant stars from $9 \\lesssim V \\lesssim 17.5$ mag using the Washington $M, T_2 + DDO51$ filter photometric prescription outlined in \\citet{M00}.  A primary motivation for the survey and a driver of its design is the selection of stars suitable for the Astrometric Grid of the Space Interferometry Mission (SIM PlanetQuest). These Grid stars, which serve as astrometric references against which the motions of SIM targets are measured, must themselves be as astrometrically stable as possible. Thus, they must be free of stellar or significant planetary companions as well as atmospheric activity (spotting/flaring) that will induce photocenter wobbles at the several microarcsecond level on of order a decade-long timescale (approximately the duration of SIM).  G and K giants were selected by the SIM project as the primary stellar constituent of the SIM Astrometric Grid because these stars are the most luminous, common stellar type found in all directions of the sky.  High intrinsic luminosity places giant stars at greater distances for a given apparent magnitude, and this increased distance decreases the angular scale of any fixed linear astrometric wobble, thus making these particular stars less likely to be problematical as references. The GGSS project has been undertaken with the specific aim of identifying {\\it subsolar} metallicity giant stars, which are intrinsically more luminous than solar metallicity giants, and also less likely to have planets \\citep[see][and references therein]{fischer05}.  Despite consideration of such effects for careful selection, all SIM Astrometric Grid candidates require pre-mission monitoring to assess their likelihood of astrometric stability during the mission. One method is to search for variability in other properties, such as brightness and radial velocity (RV).  Indeed, we are presently involved, along with other groups, in a large campaign of echelle high-resolution RV monitoring of Astrometric Grid candidates. Fortunately, not all variability in RV (or luminosity) necessarily translates to detrimental astrometric variability. However, some intrinsic photometric and/or RV variability is due to processes that, although benign in terms of photocenter wobble (e.g., radial pulsations), can mask otherwise detectable signatures of problematical sources of variability, like astrometric wobbles due to planetary, brown dwarf or stellar companions. Therefore, because campaigns to monitor radial velocity variability to the precision needed for vetting SIM targets (i.e. $\\sim100$ m s$^{-1}$) are both expensive and time consuming, it is useful to understand how intrinsic RV variability of giant stars depends on other stellar properties (like temperature, metallicity, and surface gravity) to optimize efficient selection of giant stars for the Astrometric Grid.  Although it has long been known that stars at the top of the RGB are both brightness and RV-variable \\citep[see e.g.][]{pryor88,cote96}, there has been little systematic study of this variability as a function of stellar atmospheric parameters, especially at fainter absolute magnitudes. Monitoring of radial velocities for bright red giants has indicated a high probability of their RV variations at the level of 100 m s$^{-1}$ \\citep[see][and references therein]{hatzes94}. \\citet{jorissen97} found that red giants with spectral types late G to early K are stable and giants with later spectral types are all variable.  At the same time, study of Hipparcos red clump giants \\citep{adelman01} reveals photometric stability in this evolutionary phase. Radial velocity jitter at large amplitude is expected for metal-poor, luminous red giants ($M_V \\le$ -1.4) \\citep[see][]{carney03}.  In their assessment of the suitability of giant stars as SIM Astrometric Grid members, \\citet{frink01} studied a proxy sample of 86 nearby giant stars selected from the Hipparcos catalogue.  Their study included three to eight epochs of high precision (5-8 m s$^{-1}$) echelle RV measurements on timescales from days to a year. \\citeauthor{frink01} found that most (73 of 84) of the giant stars they investigated have a stable RV at the level of $<$ 100 m s$^{-1}$, and the peak of the distribution of measured dispersions in velocity occurs at 20 m s$^{-1}$. However, \\citeauthor{frink01} do report that RV variation is more probable for redder stars ($(B-V) >$ 1.2 in their sample). While the closest match to the survey presented here, the \\citeauthor{frink01} survey contains stars that are much brighter than the expected SIM Grid sample, and, moreover, they did not probe variability as a function of metallicity, which is a key aspect of the stars being found in the GGSS.  Our goal here is to carry out an initial RV-stability assessment of stars more like those expected to fill the Astrometric Grid by focusing on stars taken directly from the GGSS itself. The hope is that RV stability may correlate with some intrinsic stellar property, such as effective temperature, surface gravity, or metallicity. Although the precision of our RV measurements is an order of magnitude less than that of \\citet{frink01}, it is appropriate for finding RV wobbles at the level needed to identify astrometrically-detrimental grid candidates for the fainter, $>1$ kpc distant GGSS candidates (and, indeed, is the precision at which monitoring campaigns of SIM Grid stars are being conducted). Moreover, our sample is almost an order of magnitude larger than that in the \\citeauthor{frink01} survey. The observations studied and discussed here were obtained under an initial JPL-sponsored program in which the internal RV accuracy was set at about 100 m s$^{-1}$. ",
        "conclusions": "We are carrying out high-resolution spectroscopic observations of an all-sky sample of giants which was pre-selected based on the GGSS program of Washington photometry. The GGSS was intended to identify stars with the highest likelihood of astrometric stability --- namely, subsolar metallicity red giants.  The main purpose of the present study is to make an initial assessment of the RV-variability properties of GGSS stars as a function of their atmospheric characteristics. The RV variation was estimated for 489 stars with spectroscopy taken on both Northern and Southern telescopes. Basic stellar parameters were estimated for the northern subsample from high-resolution spectroscopy.  A surprisingly high fraction of investigated giant stars have unstable radial velocities at 100 m s$^{-1}$ level: about 2/3 of our sample. Both of the samples, northern and southern, although having been observed independently and with different instruments and techniques, show similar distributions of RV variability.  Although a number of obvious spectroscopic binaries are included in this sample, much of the low-amplitude RV variability is probably due to atmospheric motions and temperature inhomogeneities.  A higher fraction of stars with $\\sigma < 100$ m s$^{-1}$ can be found among the objects with $4300\\,K < T_{eff} < 4700\\,K$, which corresponds to $(J-K)_0$ = 0.59 -- 0.73. If we incorporate spectroscopic metallicity and surface gravity information as well, and select only stars with -0.5 $<$ [Fe/H] $<$ -0.1 and $2.3 <  \\log~g < 3.2$, it will help to identify RV-stable candidates more effectively, but narrows the sample of potential candidates substantially. From the point of view of minimizing expensive observations, the optimal way of preliminarily selecting RV-stable stars is to use the effective temperature, or more exactly, the calibration independent range of NIR colors as $0.59 < (J-K)_0 < 0.73$.  This range encompasses 44\\% of stars with $\\sigma < 100 \\, m\\, s^{-1}$ in both our southern and northern samples; 38.5\\% of the initial GGSS sample of candidate Astrometric Grid stars (1710 of 4440) lies in this range. The result obtained here will aid in the continuing efforts to define the astrometric reference grid for SIM."
    },
    "0512/astro-ph0512564_arXiv.txt": {
        "abstract": "The hadronic interaction model BBL implements the ideas of gluon saturation due to large densities. When approaching the black body limit at high energies, leading partons acquire large transverse momenta which breaks up their coherence. This leads to a suppression of forward scattering, and is therefore important for air showers. We discuss some general aspects of this new approach and their influence on air shower properties as seen by fluorescence and surface detectors: The position of the shower maximum is reduced due to stronger absorption in the atmosphere. The lateral distribution functions become flatter for the same reason. Muons are produced abundantly due to high multiplicities in the mid-rapidity region. The response of water Cherenkov detectors and comparisons to other interaction models are shown. ",
        "introduction": "In this paper, we discuss the high energy limit of hadron nucleus scattering and their influence on air shower properties, as described in Ref. \\cite{Drescher:2004sd}. When approaching highest energies, we expect the differential scattering amplitude to become close to unity. The partons acquire a large transverse momentum of the order of the saturation scale, which leads to an steeper spectrum of forward scattered particles, the most important phase space region for air shower properties.  All air showers in this paper have been computed with the Seneca model \\cite{seneca} using UrQMD~1.3.1 \\cite{urqmd} as low-energy model below 100 GeV. ",
        "conclusions": ""
    },
    "0512/astro-ph0512278_arXiv.txt": {
        "abstract": "Most, perhaps all, stars go through a phase of vigorous outflow during formation. We examine, through 3D MHD simulation, the effects of protostellar outflows on cluster formation. We find that the initial turbulence in the cluster-forming region is quickly replaced by motions generated by outflows. The protostellar outflow-driven turbulence (``protostellar turbulence'' for short) can keep the region close to a virial equilibrium long after the initial turbulence has decayed away. We argue that there exist two types of turbulence in star-forming clouds: a primordial (or ``interstellar'') turbulence and a protostellar turbulence, with the former transformed into the latter mostly in embedded clusters such as NGC 1333. Since the majority of stars are thought to form in clusters, an implication is that the stellar initial mass function is determined to a large extent by the stars themselves, through outflows which individually limit the mass accretion onto forming stars and collectively shape the environments (density structure and velocity field) in which most cluster members form. We speculate that massive cluster-forming clumps supported by protostellar turbulence gradually evolve towards a highly centrally condensed ``pivotal'' state, culminating in rapid formation of massive stars in the densest part through accretion. ",
        "introduction": "The origin of the stellar initial mass function (IMF) is a fundamental unsolved problem in star formation. An important clue comes from millimeter and submillimeter observations of the nearby cluster-forming region, rho Oph cloud core, which uncovered prestellar condensations with a mass spectrum that resembles the Salpeter IMF (Motte et al. 1998; Johnstone et al. 2000; Stanke et al. 2005; see Testi \\& Sargent 1998 for a similar result for the Serpens core). The implication is that the IMF may be determined to a large extent by the core mass distribution from cloud fragmentation. An attractive scenario is turbulent fragmentation. Padoan \\& Nordlund (2002) showed analytically that, under plausible assumptions, an IMF-like core mass distribution can arise naturally from random turbulent compression. This result is corroborated by numerical simulations that include either driven (Li et al. 2004) or decaying turbulence (Tilley \\& Pudritz 2005; see, however, Ballesteros-Paredes et al.  2005). In this picture, the IMF is determined largely by the properties of the turbulence. Since most stars are thought to form in clusters (Lada \\& Lada 2003), the origin of the majority of stars boils down to the origin of turbulence in cluster-forming regions.  It is established that supersonic turbulence decays quickly, on a time scale comparable to the turbulence crossing time on the dominant energy-carrying scale, with or without a strong magnetic field (Mac Low et al. 1998; Stone, Ostriker \\& Gammie 1998; Padoan \\& Nordlund 1999). In a turbulent cluster-forming region, if the dissipated energy is not replenished quickly, the cloud would be in a state of low level of turbulence and/or global free-fall collapse, neither of which is commonly observed (Evans 1999; Garay 2005). The turbulence must therefore be replenished somehow. The most likely mechanism is through (proto)stellar outflows, which are observed in abundance in nearby embedded clusters such as NGC 1333 (Knee \\& Sandell 2000; Bally et al. 1996).  The idea of outflow-regulated star formation goes back to Norman \\& Silk (1980). They envisioned that star-forming clouds are constantly stirred up by the winds of optically revealed T Tauri stars. This scenario was strengthened by the discovery of molecular outflows, which point to even more powerful winds from the stellar vicinity during the protostellar phase of star formation (Lada 1985; Bontemps et al. 1996). The protostellar outflows were incorporated into a general theory of photoionization-regulated star formation in magnetic clouds by McKee (1989). Their effects on cluster formation were examined by Matzner \\& McKee (2000) using a parametrized rate of turbulence decay and a time dependent form of the virial theorem that treats the cloud as a whole. The cloud internal dynamics, such as mass distribution and velocity field, remain to be quantified.  In this letter, we present a 3D MHD simulation of cluster formation in turbulent, magnetized clouds including protostellar outflows. In agreement with previous work, we find that the initially imposed turbulence decays away quickly. We demonstrate that it is replaced by outflow-generated motions, which keep the cluster-forming region in an approximate virial equilibrium and allow for gradual star formation long after the decay of initial turbulence (\\S~2). We argue in \\S~3 that it is the protostellar outflow-generated turbulence (``protostellar turbulence'' hereafter), as opposed to the primordial (or ``interstellar'') turbulence that prevails in regions of molecular clouds that are little affected by local star formation activities, that is directly relevant to the formation of the majority of stars, including massive stars. ",
        "conclusions": ""
    },
    "0512/hep-th0512259_arXiv.txt": {
        "abstract": "{The classical dynamics of the tachyon scalar field of cubic string field theory is considered on a cosmological background. Starting from a nonlocal action with arbitrary tachyon potential, which encodes the bosonic and several supersymmetric cases, we study the equations of motion in the Hamilton--Jacobi formalism and with a generalized Friedmann equation, appliable in braneworld or modified gravity models. The cases of cubic (bosonic) and quartic (supersymmetric) tachyon potential in general relativity are automatically included. We comment the validity of the slow-roll approximation, the stability of the cosmological perturbations, and the relation between this tachyon and the Dirac--Born--Infeld one.} ",
        "introduction": "Although string theory is still far from being completely understood, there has been compelling progress in the past decade, chiefly as regards its nonperturbative objects, the $Dp$-branes, and their role in determining the stable vacuum of the theory. In particular, the study of the open string tachyon mode of unstable $Dp$-branes has been pioneered by Sen \\cite{sen5,sen6,sen7} via conformal field theory techniques. These and other methods (reviewed, e.g., in \\cite{sen04}) all agree in the main results and show that as the tachyon rolls down towards the asymptotic minimum of the potential, the brane it lives on decays into a lower-dimensional brane or the closed string vacuum. In general, the tachyon dynamics can be described by a Dirac--Born--Infeld (in short, DBI) effective action  \\cite{gar00,bers0,klu00,GHY}.  A particularly interesting approach is based on an attempt, called string field theory (SFT), to second-quantize the open bosonic string in a nonperturbative way.\\footnote{Since the single string is a collection of modes which correspond to \\emph{particle} fields, the theory of \\emph{string} fields may be regarded as a `third quantization' scheme.} The starting point is the Chern--Simons-like action \\cite{wi86a} \\be\\label{SFT} \\cS=-\\frac{1}{g_o^2}\\int \\left(\\frac{1}{2\\alpha'} \\Phi* Q_B\\Phi+\\frac13\\Phi*\\Phi *\\Phi\\right), \\ee where $g_o$ is the open string coupling constant (with dimension $[g_o^2]=E^{6-D}$), $\\int$ is the path integral over matter and ghost fields, $Q_B$ is the BRST operator, * is the star product, and the string field $\\Phi$ is a linear superposition of states in the Fock-space representation, whose coefficients correspond to the particle fields of the string spectrum. We refer to this as bosonic cubic SFT, or bosonic CSFT. (There is another string field theory, called boundary or background-independent SFT, which we will not consider here). The reader can find three of many excellent reviews on CSFT in \\cite{sen04,ohm01,ABGKM}.  At the lowest truncation level, all particle fields in $\\Phi$ are neglected except the tachyonic one, labeled $\\phi(x)$ and depending on the center-of-mass coordinate $x$ of the string. That is to say, the Fock-space expansion of the string field is truncated so that, by virtue of the state-vertex operator isomorfism, $\\Phi \\cong |\\Phi\\rangle=\\phi(x)|\\!\\downarrow\\rangle$, where $|\\!\\downarrow\\rangle$ is the ghost vacuum with ghost number $-1/2$. Then at level $(0,0)$,\\footnote{We recall that the \\emph{level} of a state is the sum of the level numbers of the creation operators acting on $|\\!\\downarrow\\rangle$. At level $(L,M)$, the string field includes terms up to level $L$, while the action includes terms up to level $M$.} the action becomes \\cite{KS1,KS2}, in $D=26$ dimensions and with metric signature $({-}{+}{\\dots}{+})$, \\be \\bar{\\cS}_\\phi=\\frac{1}{g_o^2}\\int d^D x \\left[\\frac{1}{2\\alpha'}\\phi(\\alpha'\\p_\\mu\\p^\\mu+1)\\phi-\\frac{\\l}{3}\\left(\\l^{\\alpha'\\p_\\mu\\p^\\mu/3}\\phi\\right)^3-\\Lambda\\right],\\label{tactmin} \\ee where $\\l=3^{9/2}/2^6\\approx 2.19$, $\\alpha'$ is the Regge slope, and Greek indices run from 0 to $D-1$ and are raised and lowered via the Minkowski metric $\\eta_{\\mu\\nu}$. The tachyon field is a real scalar with dimension $[\\phi]=E^2$. The extra constant $\\Lambda$ does not contribute to the equation of motion for the scalar field but it does determine its dynamical behaviour. In particular, it corresponds to the $D$-brane tension which sets the height of the tachyon potential at the (closed-string vacuum) minimum to zero. This happens when $\\Lambda=(6\\l^2)^{-1}$, which is around $68\\%$ of the brane tension; this value is lifted up when taking into account higher-level fields in the truncation scheme.  In order for CSFT to be a sensible candidate for a nonperturbative formulation of string theory, it has to reproduce the results of the conformal or boundary field theory approaches. In other words, the CSFT tachyon should be the same object as that in the latter cases (which for definiteness we shall identify with the DBI picture), sharing similar characteristics as regards rolling solutions, brane decay, and so on. With `same' object we mean that there should exist a field redefinition which maps tachyonic solutions obtained with different methods. Therefore it is important to study directly the CSFT action and its rolling tachyon solutions, which were inspected in \\cite{MZ,yan02,FH1,FH2,FGN,CST} both analytically and numerically.  The DBI model has been considered also in cosmology \\cite{gib02}, where the rolling homogeneous tachyon field can be regarded as an (or the) inflaton or dark energy field. The generalization to a cosmological context is natural because of the kind of evolution the DBI tachyon undergoes during its rolling down the potential: calling $\\rho$ and $p$ its energy density and pressure, the field can behave as a dynamical cosmological constant, $-\\rho\\lesssim p < 0$. However, since its equation of state $p=w\\rho$ is such that $-1\\lesssim w<0$, the DBI scalar cannot decay faster than matter ($\\rho\\sim a^{-3}$, $w\\sim 0$) after reheating, regardless the features of the latter. Therefore it is difficult to adopt the DBI tachyon as a model of dark energy. The problems worsen when imposing the tachyon to describe both the inflaton and quintessence field, because of the severe constraints from nucleosynthesis. The literature on the subject is rather extensive; for an overview see, for instance, \\cite{CGT,PhD} and references therein.  The rolling CSFT solutions in flat spacetime have the key property to allow for positive values of the pressure after a pressureless phase similar to that of the DBI tachyon, which may open up the possibility to describe viable cosmological scenarios. Then, as for the DBI tachyon, it would be instructive to study the cubic tachyonic action on a curved background and, in particular, in a Friedmann--Robertson--Walker (FRW) spacetime.  The goal of this paper is to present general arguments describing the behaviour of the CSFT cosmological tachyon. The action and covariant equations of motion on an arbitrary metric background are settled in sections \\ref{setup} and \\ref{eoms}, respectively, where the energy-momentum tensor for a CSFT tachyon with general potential is computed. In section \\ref{cosmo} we consider several important topics which arise when the background metric is the FRW one. Section \\ref{coseq} presents the cosmological equations, also in their Hamilton--Jacobi formulation, for either the general relativistic and braneworld case. The integration formula for the number of $e$-foldings is shown in section \\ref{cosef}. Although we do not construct exact solutions of the equations of motion, some remarks on them are collected in section \\ref{cosso}. The slow-roll appoximation is discussed in section \\ref{cossr}, where it is shown that nonlocality requires a nontrivial refinement of the standard assumptions. Approximate solutions can be found in a quasi de Sitter limit, section \\ref{cosds}, where however the action becomes purely local. Section \\ref{cosst} is devoted to the issue of stability and the presence of ghost modes. We focus on the relation between the DBI and CSFT tachyons in section \\ref{cosfr}. Discussion and prospects are in section \\ref{disc}.  Before concluding the section, let us review the kind of potentials arising in CSFT. From now on we set $\\alpha'=1$. When integrating out the other particle fields in the string field $\\Phi$, at a given truncation level the effective tachyon potential acquires new terms \\cite{ohm01}. In the bosonic CSFT, the tachyon potential $V$ at zero momentum has also quartic and higher-order contributions: \\be V(\\phi)=-\\frac{\\phi^2}{2}+\\sum_{n\\geq 4} a_n\\phi^n. \\ee In the general off-shell case, exponential differential operators will act on $\\phi$.  In superstring theory, tachyon modes appear on non-BPS branes (or $D$-$\\bar{D}$ configurations, in which case the tachyon field is complex). So the dimension $D$ in the integral measure of the action is interpreted as the dimension $D=p+1$ of an unstable $Dp$-brane.\\footnote{In the braneworld case, we assume that moduli fields from the extra directions, compactified or not, are stabilized. The volume of the compactified dimensions will be factored out from the action.} The supersymmetric version of open CSFT (henceforward CSSFT) is harder to formulate and there have been several attempts in this direction. The introduction of picture-changing operators allows for more or less natural generalizations of the noncommutative action, but also poses a number of problems not always easily solvable (we refer again to \\cite{ohm01} for a review). Here we only quote the results for the tachyon effective potential in some CSSFT's. The first proposal, by Witten \\cite{wi86b}, predicts a level 0 truncated tachyon potential which is quadratic and negative definite. At level $(1,2)$,\\footnote{All the following expressions are at zero momentum and integration of the other particle fields is understood.} higher-order terms appear in the potential but do not improve its shape, $V(\\phi)\\sim -\\phi^2/2-\\phi^4/(1-16\\phi^2)+\\dots$ \\cite{DR}. This potential has singularities and no minimum, and in general the theory suffers from tree-level contact divergences.  For these and other reasons, other two candidates seem more promising. The first one is a modified (0-picture) CSSFT with a nonchiral, bilocal double-step operator \\cite{PTY}, where contact divergences disappear. In this case the tachyon potential has two global minima \\cite{AKBM}: \\be\\label{bilV} V(\\phi)= -\\frac{\\phi^2}{4}+\\frac{\\l^{4/3}}{36}\\phi^4, \\ee at level (1/2,1). The coefficient of the quartic term has been computed also at level (2,6) but its magnitude ($\\sigma=5053/69120\\approx 0.073$) does not change significantly ($\\l^{4/3}/36\\approx 0.079$). A modified CSSFT with a chiral, local double-step operator \\cite{AMZ1,AMZ2} gives the same $S$-matrices and different off-shell effective potentials, but is less suitable for calculations.\\footnote{The inequivalent dynamics in the Neveu--Schwarz sector of local and bilocal CSSFT's is perhaps an artifact of the level truncation approximation, which should not affect the full equations of motion in a `smooth' gauge \\cite{are06}.}  The second example is Berkovits' CSSFT \\cite{ber95,ber99}. Again, the tachyon potential has a double-well shape \\cite{ber00}: \\be\\label{berV} V(\\phi)= -\\frac{\\phi^2}{4}+\\frac{\\phi^4}{2}+\\dots, \\ee where dots stand for higher level corrections which make the minima deeper. Since the qualitative features of this potential do not change when considering these terms, we shall not write them explicitly and refer to the original results \\cite{BSZ,DR2,IN}.  Among other string field theories of interest we mention closed bosonic SFT \\cite{zwi93} (see \\cite{HMT} for an overview) and heterotic SFT \\cite{OZ}. Presently, the above examples already give a good outlook of the CSFT zoology.  The cosmology of the CSFT tachyon is not a completely new subject and has been recently studied in \\cite{are04,AJ} for an approximation of the modified CSSFT with bilocal double-step operator, where approximated solutions which are asymptotically de Sitter (i.e., with constant Hubble parameter) were found in four dimensions and for a standard Friedmann equation; below we shall discuss those results and point out some related caveats. ",
        "conclusions": "\\label{disc}  The cosmological behaviour of the tachyon field in the framework of cubic string field theory is potentially rich of new consequences and worth of investigation. A priori there are several problems affecting the field. First, it turns out that the nonlocal structure of the action is of difficult treatment and sometimes obscure interpretation, in particular when linked to inflation. Second, in the Minkowskian case the rolling solutions, after the above mentioned period of matter behaviour, undergo a  wild oscillatory phase and the interpretation as a matter field breaks down. This might forbid to view the CSFT tachyon as a physical field in cosmology, if neither the Hubble friction nor other mechanisms intervene in smoothing the unstable phase. There is a number of related issues which we have tried to understand and surely would deserve further attention. We list those we believe to be more important and interesting to address. \\begin{itemize} \\item To find approximate cosmological solutions valid even outside the SR approximation. From the lessons gained in the Minkowski case, this task is likely to be carried out with strong help from numerical methods. On the other hand, an exact, complete analytical solution seems, at the present, a goal beyond hope. There are several formalisms which might lead to the right path, as the perturbative one or the $1+1$ Hamiltonian formulation. These should also help to quantize the field in a consistent way. \\item Once one or more homogeneous solutions are found, to discuss the stability of the system around these solutions. This is related to the problem of finding the cosmological perturbation equations. At this stage, is not clear whether the CSFT tachyon would generate isocurvature perturbations or not. In general, the nonadiabatic pressure perturbation $\\delta p_\\text{nad}\\equiv \\dot{p}\\,[(\\delta p/\\dot{p})-(\\delta\\rho/\\dot{\\rho})]$ vanishes identically for a fluid with a well-defined equation of state $p=p(\\rho)$. In the normal or DBI scalar field case, the presence of the inflationary attractor determines univocally the solution (one of the modes decays and can be neglected very early), and the above condition of adiabaticity is satisfied \\cite{WMLL}. The existence of such an attractor is guaranteed in an extreme extended SR regime, but as we have shown this case is trivial; in a genuine nonlocal regime it is still to be proved. \\item From the perturbation equations, extract the cosmological observables. \\item Generalize to the case of open and closed universes. \\end{itemize} All these points concern the above phenomenological model, but there are others of more essential nature, and not less urgent. A solution of the equations of motion will of course depend on the choice for the potential. In Minkowski spacetime, the tachyon solution rapidly converges as the higher string modes are included (and integrated out giving corrections to the tachyon effective potential), so that one has good accuracy already at level 2 \\cite{CST}. As noted in \\cite{ohm01}, the corrections from the higher-level modes do not typically change the qualitative shape of the effective potential obtained at 0 level, and at most deepen the bottom of the potential. We expect a similar behaviour in the presence of cosmological friction. One might also have the opposite attitude and reconstruct the potential from a given scale factor evolution.  It should be noted that at $(0,0)$ level the tachyonic action is gauge invariant but the corrections to the tachyon potential depend on the gauge choice (while the difference of the potential values between the false and true vacuum is gauge invariant). The results quoted in the introduction are based upon the Siegel gauge, which has been shown to be a reasonable choice for fixing the gauge freedom of the string field action \\cite{MS}. When selecting a specific form of the tachyon potential, it would be interesting to look also among other possibilities, some of which might be more suitable for analytic purposes. The ${\\cal B}_0$ gauge of \\cite{sch05} might be a promising choice in this direction.  Another possibility is to do an analysis without integrating the other particle fields, and include some of them (for instance, the massless gauge boson) as cosmological fields. This step is presumably subdued to the understanding of the tachyon case first.  Let us conclude by quoting Woodard \\cite{woo00}: \\emph{``In the long struggle of our species to understand the universe it has never once proven useful to invoke a theory which is nonlocal on the most fundamental level.''} Here we have not yet said anything definite about the usefulness of such theories. Nevertheless, they are interesting from a mathematical point of view and, in the context of cosmology, might give rise to radically different phenomena testable through large-scale observations.\\footnote{We should note that some nonlocal theories do provide physical pictures that can be constrained by observations. An example is given by noncommutative inflationary models (see \\cite{PhD} for related literature, in particular for Brandenberger--Ho models \\cite{BH}), which predict a blue-tilted CMB spectrum at large scales. However, there nonlocality acts at the quantum level and in an effective theory, in other words, not at the ``most fundamental level'' of Woodard's statement.\\\\ By the way, in Brandenberger--Ho models the action measure $\\sqrt{-g}$ ($=a^3$ for the FRW metric) is coupled via a * product to the Lagrangian density, which is not the case considered here.} We hope to fulfill the above agenda in the near future."
    },
    "0512/astro-ph0512034_arXiv.txt": {
        "abstract": "{According to a generally accepted paradigm, small intrinsic sizes of Compact Steep Spectrum (CSS) radio sources are a direct consequence of their youth, but in later stages of their evolution they are believed to become large-scale sources. However, this notion was established mainly for strong CSS sources.} {In this series of papers we test this paradigm on 60 weaker objects selected from the VLA FIRST survey. They have 5-GHz flux densities in the range $150 < S_{5{\\rm GHz}} < 550$\\,mJy and steep spectra in the range $0.365 \\le \\nu \\le 5$\\,GHz. The present paper is focused on sources that fulfill the above criteria and have angular sizes in the range $\\sim$$0\\farcs2$ -- $1\\arcsec$.} {Observations of 19~such sources were obtained using MERLIN in ``snapshot'' mode at 5\\,GHz. They are presented along with 1.7-GHz VLBA and 5-GHz EVN follow-up snapshot observations made for the majority of them. For one of the sources in this subsample, 1123+340, a full-track 5-GHz EVN observation was also carried out.} {This study provides an important element to the standard theory of CSS sources, namely that in a number of them the activity of their host galaxies probably switched off quite recently and their further growth has been stopped because of that. In the case of 1123+340, the relic of a compact ``dead source'' is particularly well preserved by the presence of intracluster medium of the putative cluster of galaxies surrounding it.} {The observed overabundance of compact sources can readily be explained in the framework of the scenario of ``premature'' cessation of the activity of the host galaxy nucleus. It could also explain the relatively low radio flux densities of many such sources and, in a few cases, their peculiar, asymmetric morphologies. We propose a new interpretation of such asymmetries based on the light-travel time argument.} ",
        "introduction": "Following the first paper defining the class of Compact Steep Spectrum (CSS) sources \\citep{pw82}, two almost identical surveys of CSS sources \\citep{spencer89,fanti90} based on the 3CR catalogue and the \\citet{pw82} list have been constructed. CSS sources collected in these samples -- they are often labelled jointly as the so-called 3CRPW sample -- are as powerful as large-scale Fanaroff-Riley type\\,II \\citep[FR\\,II,][]{fr74} radio sources, yet their angular sizes are of the order of a few arcseconds. Their interpretation was given by \\citet{fanti90} and has become a paradigm \\citep[see][for a review]{odea98}. According to this paradigm, the large majority of CSS sources are intrinsically small objects -- their true linear sizes are $\\la 20h^{-1}$\\,kpc\\footnote{For consistency with earlier papers in this field, the following cosmological parameters have been adopted throughout this paper: $H_0$=100${\\rm\\,km\\,s^{-1}\\,Mpc^{-1}}$ and $q_0$=0.5. Wherever in the text we refer to linear sizes we introduce $h^{-1}$. Whenever the redshift for a particular source described in this paper is not available, a default value $z=1.25$ providing the maximum linear size for a given angular size ($1\\arcsec$ corresponds then to $4.3h^{-1}$\\,kpc) and the adopted cosmological model is assumed.} -- so that their small angular sizes do not result from projection.  As their linear sizes are subgalactic, CSS sources are immersed in their host galaxies and it is obvious that the environment responsible for the observable effects, such as strong and often asymmetric depolarisation of the radio emission caused by Faraday rotation \\citep{saikia85, saikia87, ag95, ludke98, f2004}, must have an influence on morphologies and the evolution of CSS sources. A strong interaction with the host galaxy interstellar medium (ISM) is a basis of the so-called ``frustration scenario''. According to this model, the small size of CSS sources may be attributed to the presence of dense clouds of ISM \\citep{bmh84}. However, to date there is no proof that such clouds are dense {\\em enough} to impede further growth of the sources, so it is very plausible that the internal pressure of a CSS source is sufficient to allow expansion in ram pressure balance with the ambient ISM. Moreover, the infrared properties of GHz-Peaked Spectrum (GPS) and CSS sources are consistent with those of large-scale radio galaxies \\citep{f2000}. Hence, the existence of a dense medium that enshrouds galaxies hosting CSS sources and hampers the ``normal'' growth of a radio source remains unproven.  The notion that CSS sources are not frustrated but instead young has been present in the literature since the publication of the papers by \\citet{pm82} and \\citet{c85}. \\citet{f95} provided a comprehensive theoretical background of the ``youth scenario'' which has subsequently been supported observationally. In particular, \\citet{murgia99} showed that the spectral ages of CSS sources can be up to $10^5$ years. Consequently, the youth scenario currently prevails over the frustration scenario; see the review by \\citet{fanti00}. Support for this view has been brought by \\citet{siem05} who observed diffuse X-ray emission around 3C186, a CSS identified with a QSO, using the {\\it Chandra} X-ray Observatory. The discovery of a distant ($z=1.063$) cluster associated with this CSS source provides direct observational evidence that at least this CSS source is not thermally confined, as required by the frustration scenario. Instead, it appears that 3C186 may be young and is observed at an early stage of its evolution.  If a CSS source lies close to the sky plane, it is perceived as a Medium-sized Symmetric Object (MSO), which is a scaled-down version of an FR\\,II. Given that CSS sources are thought to be young, this similarity can be explained by an evolutionary scenario. Indeed, it has been argued \\citep{r96} that Compact Symmetric Objects (CSOs), MSOs and Large Symmetric Objects (LSOs) make up an evolutionary sequence. Double-lobed radio sources continually increase their size with time and LSOs begin their active phase as very compact sources and must have passed the CSO and MSO phases. However, the converse may not be true: it is not obvious that every CSO must become an LSO.  It also appears that there is no compelling reason to confine the CSS class only to the most powerful objects fulfilling the compactness and spectrum steepness criteria. In other words, the high radio power, typical for bright CSS sources, is not an essential ingredient of the definition of the CSS class. Indeed, the existence of weak CSS sources has been confirmed observationally; see Sect.~\\ref{s_new_samp}. Three questions relating to these sources arise:  \\begin{enumerate} \\item Are there any fundamental morphological differences between strong and weak CSS sources analogous to e.g. the well-known FR\\,I/FR\\,II division for LSOs which is clearly correlated with the radio luminosity \\citep{fr74}? \\item Do weak CSS sources fit into the CSO--MSO--LSO evolutionary scheme developed by \\citet{r96}? \\item What are the actual causes of the low power output? \\end{enumerate}  This paper, along with its companion papers of the series, attempts to address the above questions based on new, high resolution observations of weak CSS sources selected from the {\\it Faint Images of Radio Sky at Twenty} (FIRST) survey \\citep{wbhg97}\\footnote{Website: http://sundog.stsci.edu}. It is organised as follows. In Sect.~\\ref{s_new_samp} the selection procedure of the sample is presented. The technical details of the VLBI observations are given in Sect.~\\ref{s_obs} and the results of these observations supplemented with the earlier results of MERLIN survey are discussed in Sect.~\\ref{s_notes}. Their astrophysical implications are discussed in Sect.~\\ref{s_discuss} and summarised in Sect.~\\ref{s_concl}.  \\begin{table*}[t] \\caption[]{FIRST coordinates, flux densities and spectral indices of the sources in Subsample Two.} \\begin{tabular}{l|c c|c c c|r r|c} \\hline & RA & Dec & $F_{365\\mathrm{MHz}}$& $F_{1.4\\,\\mathrm{GHz}}$& $F_{4.85\\,\\mathrm{GHz}}$& $\\alpha_{365\\,\\mathrm{MHz}}^{1.4\\,\\mathrm{GHz}}$& $\\alpha_{1.4\\,\\mathrm{GHz}}^{4.85\\,\\mathrm{GHz}}$& $F_{1.655\\,\\mathrm{GHz}}$\\\\ \\multicolumn{1}{c|} {Source name} & \\multicolumn{2}{c|} {(J2000)} & \\multicolumn{3}{c|} {[Jy]} & & & \\multicolumn{1}{c} {[Jy]} \\\\ \\multicolumn{1}{c|} {(1)} & (2) & (3) & {(4)} & {(5)} & {(6)} & \\multicolumn{1}{c} {(7)} & \\multicolumn{1}{c|} {(8)} & {(9)}\\\\ \\hline \\hline 0744+291$^\\ast$ & 07 48 05.332 & 29 03 22.54 & 1.25 & 0.48 & 0.17 & $-$0.71 & $-$0.83 &  0.41\\\\ 0747+314 &        07 50 12.318 & 31 19 47.52 & 2.62 & 1.05 & 0.35 & $-$0.68 & $-$0.89 &  0.90\\\\ 0811+360 &        08 14 49.068 & 35 53 49.70 & 1.71 & 0.58 & 0.17 & $-$0.80 & $-$0.99 &  0.49\\\\ 0853+291 &        08 56 01.169 & 28 58 35.06 & 2.28 & 0.63 & 0.19 & $-$0.95 & $-$0.97 &  0.54\\\\ 0902+416 &        09 05 22.197 & 41 28 39.65 & 1.10 & 0.48 & 0.17 & $-$0.61 & $-$0.86 &  0.42\\\\ 0922+322$^\\ast$ & 09 25 32.727 & 31 59 52.87 & 1.80 & 0.53 & 0.20 & $-$0.91 & $-$0.79 &  0.47\\\\ 1123+340 &        11 26 23.674 & 33 45 26.64 & 3.80 & 1.32 & 0.38 & $-$0.79 & $-$1.01 &  1.11\\\\ 1232+295 &        12 34 54.387 & 29 17 43.79 & 1.10 & 0.46 & 0.16 & $-$0.65 & $-$0.84 &  0.40\\\\ 1242+364 &        12 44 49.679 & 36 09 25.53 & 2.75 & 0.78 & 0.21 & $-$0.94 & $-$1.07 &  0.65\\\\ 1251+308 &        12 53 25.750 & 30 36 35.03 & 1.03 & 0.45 & 0.20 & $-$0.61 & $-$0.66 &  0.41\\\\ 1343+386 &        13 45 36.948 & 38 23 12.62 & 1.81 & 0.85 & 0.44 & $-$0.56 & $-$0.53 &  0.78\\\\ 1401+353$^\\ast$ & 14 03 19.238 & 35 08 11.88 & 2.27 & 0.63 & 0.18 & $-$0.96 & $-$1.01 &  0.53\\\\ 1441+409 &        14 42 59.307 & 40 44 28.79 & 1.51 & 0.97 & 0.30 & $-$0.33 & $-$0.94 &  0.83\\\\ 1601+382$^\\ast$ & 16 03 35.150 & 38 06 42.93 & 0.99 & 0.43 & 0.20 & $-$0.62 & $-$0.61 &  0.39\\\\ 1619+378 &        16 21 11.290 & 37 46 04.94 & 1.62 & 0.64 & 0.20 & $-$0.69 & $-$0.94 &  0.55\\\\ 1632+391 &        16 34 02.910 & 39 00 00.18 & 1.98 & 0.93 & 0.37 & $-$0.56 & $-$0.74 &  0.82\\\\ 1656+391 &        16 58 22.172 & 39 06 25.55 & 1.34 & 0.65 & 0.24 & $-$0.53 & $-$0.80 &  0.57\\\\ 1709+303 &        17 11 19.939 & 30 19 17.67 & 1.83 & 1.03 & 0.37 & $-$0.43 & $-$0.83 &  0.90\\\\ 1717+315 &        17 19 30.062 & 31 28 48.12 & 1.06 & 0.45 & 0.15 & $-$0.63 & $-$0.87 &  0.39\\\\ \\hline \\end{tabular} \\vspace{0.2cm}  $^\\ast$ Sources with names marked with an asterisk were not followed-up using VLBI.  Description of the columns: Col.~~(1): Source name in the IAU format; Col.~~(2): Source right ascension (J2000) extracted from FIRST; Col.~~(3): Source declination (J2000) extracted from FIRST; Col.~~(4): Total flux density at 365\\,MHz extracted from Texas Catalogue; Col.~~(5): Total flux density at 1.4\\,GHz extracted from FIRST; Col.~~(6): Total flux density at 4.85\\,GHz extracted from GB6; Col.~~(7): Spectral index ($S\\propto\\nu^{\\alpha}$) between 365 and 1400\\,MHz calculated using flux densities in cols. (4) and (5); Col.~~(8): Spectral index between 1.4 and 4.85\\,GHz calculated using flux densities in cols. (5) and (6); Col.~~(9): Estimated total flux density at 1.655\\,GHz interpolated from data in cols. (5) and (6).  \\label{t_radio} \\end{table*}  \\begin{table*}[t] \\caption[]{Redshifts and photometric data of the sources in Subsample Two.} \\begin{tabular}{l l c l|c c|c c c c c} \\hline \\multicolumn{1}{c|} {Source} & \\multicolumn{1}{c} {$z$} & {Ref.} & {Opt.} & \\multicolumn{2}{c|} {APM} & \\multicolumn{5}{c} {SDSS/DR4} \\\\ \\multicolumn{1}{c|} {name} & & & \\multicolumn{1}{c|} {ID} & $R$ & $B$ & $u$ & $g$ & $r$ & $i$ & $z$\\\\ \\multicolumn{1}{c|} {(1)} & \\multicolumn{1}{c} {(2)} & (3) & \\multicolumn{1}{c|} {(4)} & (5) & (6) & (7) & (8) & (9) & (10) & (11) \\\\ \\hline \\hline 0744+291 &       &     & q    & 19.70 & 19.79 & 21.04 & 20.91 & 20.98 & 20.64 & 20.14\\\\ 0747+314 &       &     & G    &       &       & 22.55 & 23.55 & 20.87 & 19.02 & 19.24\\\\ 0811+360 &       &     & EF   &       &       &       &       &       &       &      \\\\ 0853+291 & 1.085 & (5) & Q    & 18.45 & 18.79 & 18.99 & 18.91 & 18.67 & 18.69 & 18.79\\\\ 0902+416 &       &     & G    &       &       & 22.55 & 22.67 & 21.77 & 21.07 & 20.27\\\\ 0922+322 &       &     & G    &       &       & 23.55 & 22.71 & 22.32 & 22.08 & 21.70\\\\ 1123+340 & 1.247 & (1) & G    &       &       &       &       &       &       &      \\\\ 1242+364 &       &     & G    &       &       & 23.87 & 21.94 & 21.57 & 21.23 & 20.30\\\\ 1251+308 &       &     & --   & 19.67 & 21.14 &       &       &       &       &      \\\\ 1343+386 & 1.844 & (2) & Q    & 17.69 & --    & 18.42 & 18.24 & 18.05 & 17.60 & 17.52\\\\ 1401+353 & 0.45$\\dagger$ & (3) & q    & 19.40 & 20.32 & 21.09 & 20.35 & 19.86 & 19.54 & 19.02\\\\ 1441+409 &       &     & EF   &       &       &       &       &       &       &      \\\\ 1601+382 &       &     & G    & 18.45 & 20.96 & 21.42 & 19.60 & 18.23 & 17.66 & 17.24\\\\ 1619+378 & 1.271 & (4) & Q    & 18.89 & 21.29 & 21.29 & 19.90 & 18.81 & 18.29 & 18.00\\\\ 1632+391 & 1.085 & (5) & Q    & 17.93 & 18.33 & 18.95 & 18.68 & 18.34 & 18.21 & 18.04\\\\ 1656+391 &       &     & G    &       &       & 24.08 & 22.04 & 20.26 & 19.56 & 19.13\\\\ 1709+303 &       &     & EF   &       &       &       &       &       &       &      \\\\ \\hline \\end{tabular} \\vspace{0.2cm}  Description of the columns: Col.~~(1): Source name in the IAU format; Col.~~(2): Redshift; Col.~~(3): Reference for the redshift given in col. 2; Col.~~(4): Optical identification: G -- galaxy, Q -- QSO, EF -- ``empty field'', q -- a star-like object with no known redshift; Cols.~(5-6): Magnitudes for the two POSS/APM colours; Cols.~(7-11): Magnitudes for the five SDSS colours.\\\\ $\\dagger$Photometric redshift.\\\\ References for redshifts: (1) -- \\citet{rel01}, (2) -- \\citet{hb89}, (3) -- \\citet{mach98}, (4) -- \\citet{kock96}, (5) -- SDSS.\\\\ Sources 1232+295 and 1717+315 are not listed since there is no data available for them to date and their possible optical counterparts are too faint to be measured using APM.\\\\ Note for optical identifications: ``G'' should be treated with caution as these designations mostly come from the ``automated classifier'' of SDSS.  \\label{t_optical} \\end{table*} ",
        "conclusions": "The conclusions of the study of 19~weak CSS sources with sizes below $\\approx 1\\arcsec$ presented here are as follows.  \\begin{enumerate}  \\item The sources in Subsample Two have a variety of morphological types as shown in Table~\\ref{t_types}. Such a variety remains in agreement with the general picture drawn by \\citet{fanti90} that CSS sources are intrinsically small and can be randomly oriented with regard to the sky plane. However, there are only 3 quasars among the objects investigated. Therefore, distortions caused by beaming perhaps play a secondary role compared to those resulting from the asymmetries in the distribution of host galaxy ISM (but see also point 3. below).  \\item The well established youth scenario remains valid. However, an important factor has been added to it: the activity of AGNs hosting young radio sources can be intermittent and so they can ``die early''. At least three sources with clear signs of ceased activity have been found. The relic object in 1123+340 field is of a particular importance. It is most likely a member of a cluster of galaxies and as such it is immersed in the intracluster gas which slows down the dispersion of the lobes and so helps to preserve their shape.  \\item According to RB97, the timescale of the drop in the radio luminosity of ``dying'' CSS sources is of the order of $10^5$ years. It could therefore be comparable to the light-travel time across the source itself. Unless the source does not lie in the sky plane, the apparent differences in the evolutionary stages of coasting lobes can be visible and take the form of high asymmetries in the images of its lobes which, intrinsically, can be quite symmetric.  \\item None of the sources has an FR\\,I morphology, which is in agreement with the notion that the FR\\,I-like structures develop only above the atmosphere of the host galaxy. It is not excluded, however, that weak CSS sources {\\em can} eventually become FR\\,Is  but the transition from FR\\,II-type morphology typical for CSS sources to FR\\,I-type morphology remains unclear.  \\end{enumerate}"
    },
    "0512/astro-ph0512202_arXiv.txt": {
        "abstract": "We present spatially-resolved, near-diffraction-limited  10~$\\mu$m spectra of the nucleus of the Seyfert 2 galaxy NGC~1068, obtained with Michelle, the mid-IR imager and spectrometer on the 8.1 m Gemini North telescope.  The spectra cover the nucleus and the central 6.0\\arcsec\\ $\\times$ 0.4\\arcsec\\ of the ionization cones at a spatial resolution of approximately 0.4\\arcsec\\ ($\\approx$30 parsecs). The spectra extracted in 0.4\\arcsec\\ steps along the slit reveal striking variations in continuum slope, silicate feature profile and depth, and fine structure line fluxes on subarcsecond scales, illustrating in unprecedented detail the complexity of the circumnuclear regions of this galaxy at mid-IR wavelengths. A comparison of photometry in various apertures reveals two distinct components: a compact (radius $<$ 15 pc), bright source within the central 0.4\\arcsec\\ $\\times$ 0.4\\arcsec\\ and extended, lower brightness emission. We identify the compact source with the AGN obscuring torus, and the diffuse component with the AGN-heated dust in the ionization cones. While the torus emission dominates the flux observed in the near-IR, the mid-IR flux measured with apertures larger than about 1\\arcsec\\ is dominated instead by the dust emission from the ionization cones; in spite of its higher brightness, the torus contributes less than 30\\% of the 11.6~\\mic\\ flux contained in the central 1.2\\arcsec\\ region. Many previous attempts to determine the torus spectral energy distribution are thus likely to be significantly affected by contamination from the extended emission. The observed spectrum of the compact source is compared with clumpy torus models, the first detailed comparison of such models with observational data. The models require most of the mid-IR emitting clouds to be located within a few parsecs of the central engine, in good agreement with recent mid-IR interferometric observations. We also present a UKIRT/CGS4 5 $\\mu$m spectrum covering the R(0) -- R(4) lines of the fundamental vibration-rotation band of $^{12}$CO. None of these lines was detected, and we discuss these non-detections in terms of the filling factor and composition of the nuclear clouds. ",
        "introduction": "\\label{intro}  The advent of the unified model of active galactic nuclei \\citep[AGN; e.g.][]{Antonucci93}, with its torus of dusty material,  helped make sense of many of the observed properties of the numerous classes of active galaxy, but has brought with it its own set of challenges. An intense modeling effort has been under way with the aim of understanding and predicting the emission from dusty tori and explaining the observational data \\citep{PK92,PK93,Granato94,EHY95,Manske98,Nenkova,vanBem,Schartmann05}, and with some success in matching the gross infrared properties of AGN, but several observations have proven particularly difficult to reproduce. Especially problematic have been the width of the IR spectral energy distributions (SEDs), the similarity of the observed SEDs given the wide range of variation in X-ray column densities, and, most notoriously, the behavior of the 9.7 $\\mu$m silicate feature.  While the feature is commonly seen in absorption in type 2 objects \\citep{Roche91,Sieb04}, the corresponding emission feature expected for type 1 objects has to date only been detected in a handful of objects, mostly QSOs \\citep{Hao05,Siebenmorgen05,Sturm05}. Torus models must therefore be capable of producing a silicate feature in type 2 AGN while suppressing it in most type 1 objects.  Recently, progress has been made towards overcoming some of these problems. For instance, using a model in which the dust is contained in discrete clouds, thought to be necessary to ensure dust survival in the harsh AGN environment, \\cite{Nenkova} were able to construct SEDs broad enough to match typical observed SEDs. In these models, the silicate absorption feature was visible in an edge-on view of the torus but the emission could be greatly weakened in a face-on view. By biasing the grain size distribution to large sizes, \\citet{vanBem} were able to suppress the silicate emission feature, although their models tend to predict larger torus inner radii ($\\ge$10 pc) than recent observational limits (see below). \\citet{Schartmann05} were also able to suppress the silicate emission and fit the mean type 1 SED by introducing different sublimation radii for large and small grains. Given advances such as these, torus models will benefit from high-quality observations against which they can be thoroughly tested.  Recent high spatial resolution observations of several active galaxies at 10 and 20 $\\mu$m, where a torus re-emits a substantial fraction of the radiation incident upon it, indicate that the tori may be very small. For example, the outer diameter of the mid-IR emitting dust associated with the torus has been shown to be no greater than a few parsecs in NGC~1068 and the Circinus galaxy \\citep{Jaffe,Packham05}, and $\\le$35 pc in NGC~4151 \\citep{Radomski03}, although cooler dust at larger distances is not ruled out. At the same time, some AGN are also associated with extended mid-IR emission from dust on scales of a few hundred parsecs or larger \\citep{Maiolino95, Krabbe01, Radomski02, Gorjian04, Sieb04}. On moving from large to small apertures, the IR spectrum of an active galaxy will often change from being dominated by PAH emission bands to showing a featureless continuum or silicate absorption feature \\citep{LeFloch,Soifer02,Sieb04}. Where the AGN component can be spectrally isolated from large-aperture data, the expected clear dichotomy in hard X-ray/mid-IR flux ratios between types 1 and 2 AGN is not observed, perhaps because dust heated by the nuclear source but not associated with the torus itself is contributing to the mid-IR flux \\citep{Lutz04}. It is therefore clear that observations at the highest possible spatial resolution will be necessary to disentangle the torus and extended emission and provide a fair test for models of the dusty torus.  As one of the nearest, brightest and best-studied AGN, the archetypal Seyfert 2 galaxy, NGC~1068 (14.4 Mpc distant for $\\rm H_{0}=75~km \\ s^{-1} \\ Mpc^{-1}$; 1\\arcsec\\ = 72 pc), is an excellent candidate for mid-IR observations at high spatial resolution. The existence of a hidden broad line region in NGC~1068, as revealed by the broad emission lines detected in polarized light \\citep{AM85}, along with evidence for high extinction along our line of sight, strongly suggests the presence of obscuring dust in a toroidal geometry. Further evidence for a torus comes from the strong nuclear mid-IR source coupled with emission from ionized gas extended in a roughly conical structure oriented approximately northeast - southwest \\citep[e.g.][]{Macchetto94}. Modeling of 10~$\\mu$m interferometric observations of the nucleus suggests a hot, compact ($\\le$1 pc) structure, surrounded by a more extended (approx. 2$\\times$3 pc) warm component, interpreted as dust in the hot inner wall of the torus \\citep{Jaffe}. Larger-scale 10 and 20 $\\mu$m emission from NGC~1068 is observed to extend out from the central source to the northeast and southwest \\citep{Braatz93,Cameron93,Alloin,Galliano05}; further out, low surface brightness mid-IR emission is detected from a ring of star formation located at kiloparsec distances from the nucleus \\citep{Telesco88}.  The 9.7 $\\mu$m silicate feature is seen in absorption in NGC~1068 \\citep{Roche84,Jaffe,Sieb04,Rhee05}. Previous narrow-band imaging suggests that the depth of the absorption varies over the nuclear region \\citep{Bock2,Galliano}, but the spatial coverage and spectral resolution of those studies has been insufficient to allow detailed conclusions about the distribution and profile of the feature to be drawn.  We have therefore performed spatially resolved N band spectroscopy of the nucleus and ionization cones of NGC~1068 with Michelle, the mid-IR imager and spectrometer on the 8.1 m Gemini North telescope. The spatial resolution of the spectra, approximately 0.4\\arcsec, allows us to separate the nuclear and extended emission regions and reveals the distribution of the silicate feature across the nuclear region.  Having largely isolated the spectrum of the torus in the spectrum of the central 0.4$\\times$0.4\\arcsec\\ ($\\approx$ 30$\\times$30 pc) of NGC~1068, we compare it with the clumpy torus model of \\citet{Nenkova}. The data set also allows us to examine the detailed spatial distribution of fine structure emission lines in the 10~$\\mu$m region, which have been reported previously only in very large aperture ISO measurements \\citep{Lutz00} or in few-arcsecond aperture ground-based spectroscopy \\citep{Sieb04}.  In \\S\\ref{obs} we describe the observations and reduction of the data, while \\S\\ref{results} presents the results obtained for the acquisition images, nuclear spectra and model comparison, extended silicate feature and emission lines. We also report the results of a sensitive search at 5~$\\mu$m for CO lines towards the nucleus of NGC~1068, which serves as a probe of the nature and distribution of molecular gas near the center of this galaxy. The implications of this work are discussed in \\S\\ref{discuss} and summarized in \\S\\ref{conc}. ",
        "conclusions": "\\label{conc}  N-band near-diffraction-limited observations of the nuclear region of NGC~1068 at the Gemini North telescope show that the complexity of the nuclear region, previously seen in mid-IR imaging and observations at shorter wavelengths, persists in mid-IR spectra. The spectra vary widely on spatial scales of less than 1\\arcsec, in the slope of the continuum, the strength of the silicate absorption feature, and the fine structure line fluxes. The optical depth of the silicate feature is observed to vary between about 0.2 and 0.45, and the silicate absorption and all three of the [Ar~III]+[Mg~VII], [S~IV] and [Ne~II] fine structure features are asymmetrically distributed within 1\\arcsec\\ of the center. The relative strengths of the above three emission features imply a strong contribution by the high excitation [Mg~VII] line to the blend of that line and [Ar~III].  A high spectral resolution UKIRT/CGS4 spectrum of the central 1\\arcsec\\ of NGC~1068 shows no evidence for absorption by CO. This upper limit implies either a small filling factor for dense molecular clouds obscuring the nuclear continuum source(s), a predominance of obscuring diffuse cloud material lacking in CO (whose existence also is implied by the interstellar 3.4~$\\mu$m absorption feature seen towards the nucleus), or a combination of the two.  The mid-IR emission from NGC 1068 originates in two distinct components. Emission from the central region with diameter 0.4\\arcsec\\ is dominated by the IR - bright obscuring torus. Our observations place an upper limit of 15 pc on the torus radius, in agreement with the compact ($\\sim$3 pc) size determined by mid-IR interferometry. We have been able to obtain good fits with the clumpy torus models of \\citet{Nenkova} to this nuclear region. Mid-IR emission in apertures larger than $\\sim$ 1\\arcsec\\ is dominated by a second component, the lower brightness, diffuse emission from the AGN-heated dust in the ionization cones. If NGC 1068 is representative of other sources, the implication is that AGN SEDs are highly uncertain since the torus does not dominate large-aperture photometry at wavelengths longer than the near IR. While this has long been known to be the case for facilities such as ISO and IRAS, our results underline the need for future thermal-IR observations of AGN to be obtained at the highest possible spatial resolution even at mid-IR wavelengths."
    },
    "0512/astro-ph0512458_arXiv.txt": {
        "abstract": "We investigate the effects of an external magnetic helicity production on the evolution of the cosmic axion field. It is shown that a helicity larger than $(few \\times 10^{-15} \\G)^2 \\Mpc$, if produced at temperatures above a few $\\GeV$, is in contradiction with the existence of the axion, since it would produce too much of an axion relic abundance. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512172_arXiv.txt": {
        "abstract": "{This paper presents accurate spectral energy distributions (SEDs) of 16 M31 globular clusters (GCs) confirmed by spectroscopy and/or high spatial-resolution imaging, as well as 30 M31 globular cluster candidates detected by Mochejska et al. (1998). Most of these candidates have $m_V > 18$, deeper than previous searches, and these candidates have not yet been confirmed to be globular clusters. The SEDs of these clusters and candidates are obtained as part of the BATC Multicolor Survey of the Sky, in which the spectrophotometrically-calibrated CCD images of M31 in 13 intermediate-band filters from 4000 to 10000{\\AA} were observed. These filters are specifically designed to exclude most of the bright and variable night-sky emission lines including the OH forest. In comparison to the SEDs of true GCs, we find that some of the candidate objects are not GCs in M31. SED fits show that theoretical simple stellar population (SSP) models can fit the true GCs very well. We estimate the ages of these GCs by comparing with SSP models. We find that, the M31 clusters range in age from a few ten Myr to a few Gyr old, as well as old GCs, confirming the conclusion that has been found by Barmby et al. (2000), Williams \\& Hodge (2001), Beasley et al. (2004), Burstein et al. (2004) and Puzia et al. (2005) in their investigations of the SEDs of M31 globular clusters. ",
        "introduction": "Galactic globular clusters are among the oldest stellar objects in the universe, provide vitally important information regarding the minimum age of the universe and the early formation history of our Galaxy. Studying the integrated properties of extragalactic globular clusters can help us understand the evolutionary history of distant galaxies (cf. \\cite{bur04}). Globular clusters (GCs) are bright, easily identifiable stellar populations with homogeneous abundances and ages. M31 is an ideal target for studying GCs, since it comprises the largest and nearest sample of GCs, which is more than all GCs combined in the other Local Group members (\\cite{battis87}; \\cite{raci91}; \\cite{harris91}; \\cite{fusi1993}).  Many GC searches in M31 have been conducted. The first catalog of M31 GC was presented by Hubble (1932), who discovered 140 GCs with $m_{pg}\\leq 18$ mag. Then, a number of catalogs of GC candidates were published. Vete\\u{s}nik (1962a) compiled the first major catalog, containing about 300 GC candidates identified by the previous works (e.g., \\cite{hubble32}; \\cite{seynas45}; \\cite{hilt58}; \\cite{mayegg53}; \\cite{kronmay60}).  Later, several other major catalogs of M31 GC candidates were compiled by Sargent et al. (1977), Crampton et al. (1985) and by the Bologna Group (\\cite{battis80,battis87,battis93}). Although these catalogs may be fairly complete down to $V=18$ ($M_v \\sim -6.5$) (\\cite{fusi1993}), recent works have searched for fainter GCs in M31 (e.g., \\cite{moche98}; \\cite{bh01}).  These latter samples provide a good database for studies of M31 GCs, including their luminosity function (e.g., \\cite{auri92}; \\cite{moche98}; \\cite{barmby01}), reddening and intrinsic color (e.g., \\cite{vete62b}; \\cite{bage77}; \\cite{iyeri85}; \\cite{bh00}; \\cite{barmby02b}), metallicities (e.g., \\cite{vand69}; \\cite{ash93}; \\cite{bh00}; \\cite{perr02}), structure parameters (\\cite{barmby02a}) and comparisons with the Galactic GCs and M33 GCs (e.g., \\cite{hilt60}; \\cite{fpc80}; \\cite{rhh92}; \\cite{moche98}). Near infrared photometry of M31 GCs has been done as well (e.g., \\cite{fpc80}; \\cite{Sitko}; \\cite{Bonolia,Bonolib}; \\cite{Cohen}; \\cite{bh00}), with Galleti et al. (2004) using 2MASS data for 693 GC candidates.  In a previous paper from our group, Jiang et al. (2003) obtained CCD multicolor photometry for 172 M31 GCs and candidates in 13 intermediate-band filters spanning wavelengths from 4000 to 10000{\\AA}.  Barmby et al. (2000); Beasley et al. (2004); Burstein et al. (2004) and Puzia et al. (2005), using SEDs, have found a set of young to intermediate-age GCs in M31, as well as the ``usual'' complement of GCs as old as Galactic GCs. Jiang et al. (2003) estimated ages for their M31 GCs by comparing their SEDs to those of theoretical simple stellar population (SSP) models, finding at least 8 M31 GCs and candidates younger than 1 Gyr. Recently, Fusi Pecci et al. (2005) found a population of 67 massive blue clusters in the disk of M31, which they interpret as globular clusters, with ages less than $\\sim2~\\rm{Gyr}$.  M31 was observed as part of galaxy calibration program of the Beijing-Arizona-Taiwan-Connecticut (BATC) Multicolor Sky Survey (e.g., \\cite{fan96}; \\cite{zheng99}), which has a custom-designed set of 15 intermediate-band filters to do spectrophotometry for preselected 1 deg$^{2}$ regions of the northern sky. In this paper, we present the SEDs for 15 M31 GCs confirmed by spectroscopy and/or high spatial-resolution imaging and 30 M31 GC candidates detected by Mochejska et al. (1998) that lie within the BATC field of view. Details pertaining to the observations and data reduction are given in \\S~2. Analysis of the M31 GC SEDs is reported in \\S~3. A summary is given in \\S~4. ",
        "conclusions": "We have obtained SEDs of 45 faint M31 GCs and GC candidates detected by Mochejska et al. (1998) in 13 intermediate colors with the NAOC 60/90 cm Schmidt telescope. By analyzing their SEDs, 3 are likely background objects. SED fits show that theoretical simple stellar population (SSP) models can fit the true M31 GCs very well. We confirm 20 candidates as M31 GCs, and we can estimate their ages and metallicities by comparing with SSP models. We find that these faint M31 GCs range in age from a few ten Myr to a few Gyr old, as well as old GCs, confirming the conclusions that has been found by Barmby et al. (2000), Williams \\& Hodge (2001), Beasley et al. (2004), Burstein et al. (2004) and Puzia et al. (2005).  The fact that we can find faint GCs around M31 that have a similar wide range of ages as the brighter GCs that Beasley et al. (2004) and Burstein et al. (2004) have found is consistent with the formation scenario proposed in Burstein et al. (2004). In that scenario, these GCs come from dwarf galaxies that have been accreted by M31 in the past. Evidently, the history of accretion of its contingent of dwarf galaxies by M31 appears to be very different from that our Galaxy has experienced.  In Table 5, there is a GC, the estimated age of which is 20 Gyr. This age only means that this GC is as old as Galactic GCs."
    },
    "0512/astro-ph0512491_arXiv.txt": {
        "abstract": "We present new astrometry of Pluto's three satellites from images taken of the Pluto system during 2002-3 with the High Resolution Camera (HRC) mode of the Advanced Camera for Surveys (ACS) instrument on the Hubble Space Telescope.  The observations were designed to produce an albedo map of Pluto but they also contain images of Charon and the two recently discovered satellites, S/2005 P1 and S/2005 P2\\null. Orbits fitted to all three satellites are co-planar and, for Charon and P2, have eccentricities consistent with zero.  The orbit of the outermost satellite, P1, has a significant eccentricity of 0.0052 $\\pm$ 0.0011. Orbital periods of P1, P2, and Charon are 38.2065 $\\pm$ 0.0014, 24.8562 $\\pm$ 00013, and 6.3872304 $\\pm$ 0.0000011 days, respectively.  The total system mass based on Charon's orbit is $1.4570 \\pm 0.0009$~x~$10^{22}$~kg. We confirm previous results that orbital periods are close to the ratio of 6:4:1 (P1:P2:Charon) indiciative of mean-motion resonances, but our results formally preclude precise integer period ratios.  The orbits of P1 and P2, being about the barycenter rather than Pluto, enable us to measure the Charon/Pluto mass ratio as 0.1165$\\pm$0.0055.  This new mass ratio implies a density of 1.66 $\\pm$ 0.06 g~cm$^{-3}$ for Charon and 2.03 $\\pm$ 0.06 g~cm$^{-3}$ for Pluto thus adding confirmation that Charon is somewhat under-dense relative to Pluto.  Finally, by stacking all images, we can extract globally averaged photometry.  P1 has a mean opposition magnitude of $V=24.39 \\pm 0.02$ and color of $(B-V) = 0.644 \\pm 0.028$.  P2 has a mean opposition magnitude of $V=23.38 \\pm 0.02$ and color of $(B-V) = 0.907 \\pm 0.031$.  The colors indicate that P1 is spectrally neutral and P2 is slightly more red than Pluto.  The variation in surface color with radial distance from Pluto is quite striking (red, neutral, red, neutral) and begs further study. ",
        "introduction": "The study of Pluto was greatly facilitated in 1978 with the discovery of its first satellite, Charon \\citep{chr78}\\null.  That discovery made possible the accurate determination of its mass which had previously been largely a matter of conjecture.  Later, in the late 1980's, Pluto studies were transformed by the mutual events between Pluto and Charon \\citep[e.g.,][]{bui92, bin97, you01}\\null.  Charon remains an interesting object it is own right, but its role as a tool from which to understand the system should not be understated.  Two new moons were recently discovered in orbit around Pluto \\citep{wea05}\\null.  More precisely, they orbit the center of mass of the system, which is very close to the Pluto-Charon barycenter. As with Charon, these new objects will be studied in their own right and will also be useful as probes or test masses in the Pluto system. Given astrometry of sufficient precision and time-base, one can now easily deduce the precise Charon\\slash Pluto mass ratio.  One might also hope to determine the masses of the new satellites through their mutual perturbations.  However, their mutual gravitational force is more than 3 orders of magnitude weaker than the force exerted on them by Pluto and Charon, so the dynamics of their presumably resonant orbits may well completely mask any measurable effect P1 and P2 may have on each other.  The preliminary orbits computed by \\citet{wea05} were based on just two epochs of data separated by only three days, much less than a full orbit of either satellite.  Also, the data were derived from images where Pluto and Charon were both saturated.  These constraints led to a restricted solution for the orbit where it was assumed that the objects were in circular orbits in the same orbital plane as Charon.  As we shall show, this assumption turned out to be very close to the correct answer.  The data presented in this work are derived from pre-discovery HST observations that span multiple orbits of all satellites and do so with images where Pluto and Charon are not saturated.  This paper presents the first unrestricted fits to the orbits of the new satellites using pre-discovery HST observations. ",
        "conclusions": ""
    },
    "0512/astro-ph0512344_arXiv.txt": {
        "abstract": "Hypervelocity stars have been recently discovered in the outskirts of galaxies, such as the unbound star in the Milky Way halo, or the three anomalously fast intracluster planetary nebulae (ICPNe) in the Virgo Cluster. These may have been ejected by close 3-body interactions with a binary supermassive black hole (SMBBH), where a star which passes within the semimajor axis of the SMBBH can receive enough energy to eject it from the system. Stars ejected by SMBBHs may form a significant sub-population with very different kinematics and mean metallicity than the bulk of the intracluster stars. The number, kinematics, and orientation of the ejected stars may constrain the mass ratio, semimajor axis, and even the orbital plane of the SMBBH. We investigate the evolution of the ejected debris from a SMBBH within a clumpy and time-dependent cluster potential using a high resolution, self-consistent cosmological N-body simulation of a galaxy cluster. We show that the predicted number and kinematic signature of the fast Virgo ICPNe is consistent with 3-body scattering by a SMBBH with a mass ratio $10:1$ at the center of M87. ",
        "introduction": "A significant fraction of the stellar component of a galaxy cluster is not confined to any galaxy. This intracluster light (ICL) has been identified via planetary nebulae (PNe) (Feldmeier et al. 2004; Aguerri et al. 2005), Red Giant Branch (RGB) stars (Durrell et al. 2002), and ultra-deep surface photometry (e.g., Feldmeier et al. 2002, 2004; Mihos et al. 2005). It is commonly thought that intracluster stars are stripped from galaxies as the cluster assembles (Merritt 1984), via high speed galaxy encounters (Moore et al 1996), interactions with the cluster potential (Byrd \\& Valtonen 1990), or by tidal stripping in infalling groups (Mihos 2004). These processes will generate a debris field that is highly inhomogeneous, with distinctly non-Gaussian velocities, reflecting an unrelaxed intracluster population (Napolitano et al. 2003).   The continued search for intracluster PNe (ICPNe) has revealed several in the Virgo Cluster with radial velocities that are extremely rapid compared both to its nearest galaxy and to the background cluster. In one pointing $\\sim 65$ kpc from M87, Arnaboldi et al. (2004) found that of 15 ICPNe radial velocities, only 12 were consistent with M87's stellar velocity dispersion profile (Romanowsky \\& Kochanek 2001) and systemic velocity. Of the three remaining ICPNe, two had velocities which were offset $\\sim -900$ km/sec from the systemic velocity of M87 and one had a $\\sim +1500$ km/sec difference.  While these ICPNe could be associated with an unrelaxed tidally-stripped intracluster stellar population, there is an alternative explanation. In this paper, we explore the possibility that the high velocity ICPNe were ejected after interacting with a supermassive binary black hole (SMBBH) at the center of M87. Supermassive black holes (SMBH) are a standard component of elliptical galaxies and spiral bulges (e.g., Richstone et al. 1998). When galaxies evolve through hierarchical merging, each SMBH participates and at some point forms a hard binary. This implies that many galaxies host SMBBHs. Perhaps the best evidence for this is the double X-ray bright nucleii in NGC 6240 (Komossa et al. 2003).  During the hard binary phase, stars that pass between the two SMBHs siphon energy from the binary's orbit via 3-body scattering, and most are ejected with  hypervelocities (Hills 1988; Yu \\& Tremaine 2003). These SMBBH-driven ejecta may constitute a different ICL population than stars stripped by dynamical interactions between galaxies; they are more likely to be metal rich, and, as we will show, have distinct kinematics and spatial structure. This makes the two populations easy to separate.   We investigate the structure and kinematics of the debris ejected via 3-body scattering within a high resolution N-body simulation of a cluster potential. We model the SMBBH ejection velocities and study how well the ejected population remains kinematically distinct within a clumpy and still evolving cluster potential. We compare the expected mass fractions of the ICL generated by 3-body interactions with more traditional tidal interactions. According to our models, it is plausible that the fast ICPNe were ejected after interacting with a SMBBH in M87. We discuss how this theory can be tested with spectroscopic observations. ",
        "conclusions": ""
    },
    "0512/astro-ph0512399_arXiv.txt": {
        "abstract": "{ In our continuing optical spectroscopic campaign to identify the longer-wavelength counterparts of newly-discovered hard X--ray sources detected by {\\it INTEGRAL}, we observed the putative optical counterparts of seven southern sources at the South African Astronomical Observatory and at the European Southern Observatory. For two of these objects, optical photometry was also acquired. These observations firmly established the nature of four of them: we found that IGR J10404$-$4625 (=LEDA 93974), 4U 1344$-$60 and IGR J16482$-$3036 are Active Galactic Nuclei (AGNs) at redshifts $z$ = 0.0237, 0.013 and 0.0313, respectively, and that 2RXP J130159.6$-$635806 is a Galactic High-Mass X--ray Binary (HMXB). We also give possible optical identifications for three further objects, namely IGR J11215$-$5952, IGR J11305$-$6256 and IGR J16207$-$5129, which are consistent with being Galactic HMXBs. Physical parameters for these objects are also evaluated by collecting and discussing the available multiwavelength information. The detection of four definite or likely HMXBs out of seven objects in our sample further stresses {\\it INTEGRAL}'s crucial contribution in hunting this class of object. Also, the determination of the extragalactic nature of a substantial fraction of the {\\it INTEGRAL} survey sources underlines the importance of hard X--ray observations for the study of background AGNs located beyond the `Zone of Avoidance' of the Galactic Plane. ",
        "introduction": "One of the objectives of the {\\it INTEGRAL} mission (Winkler et al. 2003) is to survey the hard (20--100 keV) X--ray sky, with particular attention to the Galactic Plane. The capabilities of the IBIS instrument (Ubertini et al. 2003) allow {\\it INTEGRAL} to detect hard X--ray sources at the mCrab level with a typical localization accuracy of 2-3$'$ (Gros et al. 2003). This combination of sensitivity and positional accuracy is unprecedented for surveys in the hard X--ray band above 20 keV and made it possible, for the first time, to resolve crowded regions of the hard X--ray sky such as the Galactic Centre and the spiral arms. These capabilities also allowed the discovery of many new hard X--ray extragalactic objects beyond the Galactic Plane (the so-called `Zone of Avoidance'), where interstellar obscuration hampers observations in soft X--rays.  Since the launch of {\\it INTEGRAL} in October 2002, IBIS has already scanned large portions of the sky, and allowed the production of two deep (down to mCrab sensitivities) Galactic Plane Surveys (Bird et al. 2004, 2005; Bassani et al. 2004, 2005) as well as surveys of the Galactic Centre (Revnivtsev et al. 2004), of spiral arms of the Galaxy, such as the Sagittarius (Molkov et al. 2004) and Crux (Revnivtsev et al. 2005) arms, and of the Coma cluster of galaxies (Krivonos et al. 2005).  Within these surveys, the ISGRI detector of IBIS revealed more than 200 sources between 20 and 100 keV, the relative majority of them ($\\sim$50\\%) being Galactic X--ray binaries and with a smaller percentage ($\\sim$10\\%) of Active Galactic Nuclei (AGNs). However, about one fourth of the sample has no secure counterpart at longer wavelengths and therefore cannot yet be associated with any known class of objects. The majority of these unidentified sources are believed to be X--ray binaries, although some of them have been identified as AGNs (e.g., Masetti et al. 2004, 2005; hereafter Papers I and II).  In order to reduce the {\\it INTEGRAL} error circle, correlations with catalogues at longer wavelengths (soft X--ray, optical, near- and far-infrared, and/or radio) are employed. In particular, Stephen et al. (2005a,b) found a strong positional correlation between the ISGRI objects and the softer X--ray sources in the {\\it ROSAT} catalogue of bright sources (Voges et al. 1999), showing that a bright {\\it ROSAT} source, if present within an ISGRI error circle, is very likely the soft X--ray counterpart of the corresponding {\\it INTEGRAL} object. Similarly, Sazonov et al. (2005) accurately determined the positions of the soft X--ray counterparts of 6 {\\it INTEGRAL} sources by using {\\it Chandra} data. The use of the positional information coming from soft X--ray satellites therefore increases the positional accuracy to a few arcsecs, thus making the optical searches much easier.  Also, the presence of a radio object within the IBIS error box can again be seen as an indication of an association between the radio emitter and the {\\it INTEGRAL} source (e.g., Paper I; Paper II). However, whereas the cross-correlation with catalogues at other wavebands is fundamental in pinpointing the putative optical candidates, only optical spectroscopy can reveal the exact nature of the X--ray emitting object.  Additionally, broadband optical photometry can help to determine other characteristics, such as the overall spectral energy distribution and absolute magnitude.  Thus, in our continuing effort to identify the unknown {\\it INTEGRAL} sources, here we concentrate on a sample of 7 southern objects for which likely bright candidates could be pinpointed on the basis of their association with sources in other wavebands.  These are mainly in soft X--rays and radio, or for which a bright emission-line star could be detected within the ISGRI error circle.  Admittedly, the latter cases are the weaker ones among our selection of putative counterparts, due to the relatively large {\\it INTEGRAL} error boxes and the lack of more accurate catalogued localizations at shorter wavelengths (other than optical). However, in several instances a bright emission-line star in an {\\it INTEGRAL} error circle has been found to be the actual optical counterpart of the source responsible for the detected hard X--ray emission (e.g., Reig et al. 2005), so we consider these stars as important and viable candidates. We refer the reader to Reig et al. (2005) for a detailed discussion of the possibilities and caveats that these putative associations imply.  We moreover remark the following: using the spectral information of the stars in the Hipparcos catalogue (Perryman et al. 1997), together with (i) the known number densities of bright stars, (ii) the proportion of early-type stars showing emission lines and (iii) the percentage of stars in each spectral class (as in Allen 1973), we find that, in our cases, the chance probability of finding a bright blue emission-line star within the {\\it INTEGRAL} error box is less than 0.05 \\% along the Galactic Plane. This makes us confident that our choice of putative optical bright candidates of early spectral type and with emission lines is fully justified.  We present here the spectroscopic results on these 7 sources obtained at the South African Astrophysical Observatory (SAAO) and at the European Southern Observatory (ESO). For two of them, optical photometry was also obtained, and is reported here. In Sect. 2 we present the sample of objects selected for this observational campaign, whereas in Sect. 3 a description of the observations is given; Sect. 4 reports the results for each source and discusses their nature. Conclusions are drawn in Sect. 5.  In the following, when not explicitly stated otherwise, for our X--ray flux estimates we will assume a Crab-like spectrum. ",
        "conclusions": "We have presented the results of the third stage (the first one dealing with southern objects), this time accomplished at SAAO and at ESO, of our ongoing observational campaign aimed at the identification of newly-discovered {\\it INTEGRAL} objects of unknown nature. We spectroscopically observed the putative optical counterparts of seven southern {\\it INTEGRAL} sources. For two cases, optical photometry was also obtained. This approach, already very successful as demonstrated in Papers I and II, allowed us to firmly determine the nature of at least four of these objects.  We found that: (i) IGR J10404$-$4625 (=LEDA 93974) is a Seyfert 2 galaxy in the Compton-thin regime at $z$ = 0.0237; (ii) 2RXP J130159.6$-$635806 is a non-supergiant HMXB in the Crux Arm, at a distance $\\approx$7 kpc; (iii) 4U 1344$-$60 is a Seyfert 1 galaxy at $z$ = 0.013; and (iv) IGR J16482$-$3036 is a Seyfert 1 galaxy at $z$ = 0.0313 with a central black hole of mass $\\sim$1.4$\\times$10$^{8}$ $M_\\odot$.  We also give possible identifications for three further cases, namely: (i) HD 306414 is the likely counterpart of IGR J11215$-$5952 and thus probably a SFXT HMXB located at a distance $d \\sim$ 6.2 kpc; (ii) HD 100199 is possibly associated with the hard X--ray source IGR J11305$-$6256, which, if correct, would then resemble the persistent HMXB X Per, but at a distance $\\sim$3 kpc; (iii) HD 146803 is tentatively associated with IGR J16207$-$5129; which would identify it as a HMXB, possibly similar to 1H 0739$-$529, and located at a distance of $\\sim$1.9 kpc. However, more extensive multiwavelength studies of the error boxes of these sources, especially through pointed soft X--ray observations (obtainable with satellites such as {\\it Chandra}, {\\it XMM-Newton} or {\\it Swift}), will allow a conclusive test of these three possible associations.  Moreover, all of the above once again demonstrates, as already remarked in Papers I and II, that {\\it INTEGRAL} is making a fundamental contribution in the detection and study of a substantial fraction of persistent and transient HMXBs along the Galactic Plane, and of background AGNs located beyond the Zone of Avoidance of the Galactic Plane."
    },
    "0512/astro-ph0512485_arXiv.txt": {
        "abstract": "Extended, soft X-ray emission from the halo of a very large disk galaxy has been detected. The luminosity and surface brightness distribution is in excellent agreement with predictions by recent, cosmological galaxy formation models. Predicted Ly$\\alpha$ emission, associated with ``cold'' accretion of filamentary gas onto galaxies, is discussed in relation to Ly$\\alpha$ ``blobs''. Finally, the predicted evolution of the Tully-Fisher relation, going from $z$=0 to 1, is discussed in relation to recent observations. ",
        "introduction": "\\begin{figure}[ht] \\hfill\\includegraphics[angle=0,width=2.3in]{sommerlarsen.fig1_left.ps} \\hfill\\includegraphics[angle=0,width=2.3in]{sommerlarsen.fig1_right.ps}\\hspace*{\\fill} \\caption{Left: Predicted and observed 0.2-2 keV luminosities of X-ray haloes as a function of disc circular velocity. All X-ray luminosities have been calculated within the same physical aperture as used for NGC 5746. The filled circles are from the observations of NGC 5746 and NGC 5170 (1-$\\sigma$ upper limit of $L_X<2.9~10^{39}$ erg/s, 0.3-2 keV), while other symbols are the predictions from simulations with a range of different resolutions and circular velocities: Triangles are for simulations with low-metallicity chemical composition while squares are for simulations with self-consistent chemical evolution run at 8, 64, and 512 times the original resolution (Sommer-Larsen, Gotz and Portinari 2003) corresponding to gas particle masses of 7.5~10$^5$, 9.4~10$^4$ and 1.2~10$^4~h^{-1}M_{\\odot}$, respectively ($h$=0.65). Results from simulations re-run at higher resolution are connected with lines. Open squares, except the simulated galaxy with a circular velocity of $\\sim$225 km/s, are for simulations run with a universal baryon fraction of 0.15. All other simulations were run with a baryon fraction of 0.1. Right:  Predicted and observed surface brightness profile of X-ray haloes as function of the distance to the disc mid-plane. Filled circles are NGC 5746 data while other symbols mark simulations with different resolutions and circular velocities. The vertical dashed line indicates D25 of NGC 5746.} \\end{figure} Disk galaxies continue forming to the present day, as evidenced by, e.g., infall of high-velocity clouds and satellite galaxies. Self-consistent models of disk galaxy formation predict that, at present, part of the inflowing gas originates as hot and dilute, low-metallicity halo gas, slowly cooling out and accreting onto the disk (e.g., Abadi et al. 2003, Sommer-Larsen et al. 2003, Governato et al. 2004, Robertson et al. 2004). The X-ray luminosity of the hot halo is predicted to increase strongly with galaxy mass, and the haloes of the most massive galaxies should be detectable (Toft et al. 2002, Rasmussen et al. 2004). Searches for this hot halo gas have so far failed to detect such soft X-rays, challenging galaxy formation theory. Moreover, it has been suggested that for galaxies of total mass less than a few times 10$^{11} M_{\\odot}$, most gas stays cold during the accretion onto the galaxy (apart from a very transient radiative shock phase, e.g., Birnboim \\& Dekel 2003; Dekel \\& Birnboim 2005; next section). This may result in reduced halo X-ray emission. On the other hand, the recent detection of a warm-hot phase of the intergalactic medium shows the presence of a reservoir of hot and dilute gas at galactic distances (Nicastro et al. 2005). Furthermore, absorption of the OVI line in quasar spectra (Wakker et al. 2004), and the head-tail structure of some high-velocity clouds in the halo of the Milky Way (Bruns et al. 2000) provide circumstantial evidence  that the Milky Way is surrounded by an extended hot halo.  As a test of current disk galaxy formation models we used the Chandra X-ray telescope to study the most promising candidate spiral galaxy for detecting halo X-ray emission, NGC 5746. We also studied a similar, but less massive galaxy, NGC 5170, as a test of our procedure. The galaxies are massive and nearby (NGC 5746 is an SBb galaxy at $d$=29.4 Mpc and has $V_c$=307$\\pm$5 km/s; NGC 5170 is an Sc galaxy at $d$=24.0 Mpc and has $V_c$=250$\\pm$5 km/s). Both galaxies are quiescent, showing no signs of either starburst activity, interaction with other galaxies, or an active galactic nucleus. The disks of the galaxies are viewed almost perfectly edge-on.  Diffuse, soft X-ray emission extending more than 20 kpc from the stellar disc was detected around NGC 5746. A total of about 200 net counts in the 0.3-2 keV band were detected from the halo of  NGC 5746, corresponding to a 4.0-$\\sigma$ detection. The same observing technique and data analysis revealed no diffuse emission around the less massive galaxy NGC 5170. Moreover, from Fig.1 it follows that the agreement between observations and models is very good. Full discussions of the results are presented in Pedersen et al. (2005), and Rasmussen et al. (2005). ",
        "conclusions": ""
    },
    "0512/astro-ph0512166_arXiv.txt": {
        "abstract": "We present the small-scale ($0.01<r<8\\,h^{-1}~\\Mpc$) projected correlation function $\\wpp(\\rp)$ and real space correlation function $\\xi(r)$ of $24520$ luminous early-type galaxies from the Sloan Digital Sky Survey Luminous Red Galaxy (LRG) sample ( $0.16<z<0.36$). ``Fiber collision'' incompleteness of the SDSS spectroscopic sample at scales smaller than $55~\\arcs$ prevents measurements of the correlation function for LRGs on scales smaller than $\\sim 0.3~\\Mpc$ by the usual methods. In this work, we cross-correlate the spectroscopic sample with the imaging sample, with a weighting scheme to account for the collisions, extensively tested against mock catalogs. We correct for photometric biases in the SDSS imaging of close galaxy pairs. We find that the correlation function $\\xi(r)$ is surprisingly close to a $r^{-2}$ power law over more than 4 orders of magnitude in separation $r$. This result is too steep at small scales to be explained in current versions of the halo model for galaxy clustering. We infer an LRG-LRG merger rate of $\\lesssim 0.6~\\Gyr^{-1}\\Gpc^{-3}$ for this sample. This result suggests that the LRG-LRG mergers are not the main mode of mass growth for LRGs at $z<0.36$. ",
        "introduction": "Galaxy clustering is a tool for the study of a diverse set of phenomena at different scales. On large scales ($\\sim 100~\\Mpc$) the density perturbations are small enough to be described in a linear framework, allowing analysis of cosmological models in detail and constraints on cosmological parameters \\citep{tegmark03b}. In addition, recent studies have used extremely large-scale galaxy clustering to test fundamental cosmological hypotheses, such as homogeneity of the Universe \\citep{hogg05a} and flatness of the universe \\citep{eisenstein05b}.  On intermediate scales ($0.3$ to $30 ~\\Mpc$), galaxy clustering probes the relation of galaxies to dark matter through the biased clustering of dark matter halos. On these scales the two-point correlation function is found to be very close to a power law. Recent work has shown that the observed small deviations of the correlation function from a power-law on these scales can be interpreted in the framework of the Halo Occupation Distribution \\citep[HOD;][]{peacock00a, seljak00a, scoccimarro01a, berlind02a} as a natural transition between galaxy pairs within a single virialized halo and galaxy pairs in separate halos \\citep[\\eg, ][]{zehavi04a,zehavi05b}.  On sufficiently small scales within a halo ($< 1~\\Mpc$), galaxy clustering will probe more than just the properties of the dark matter. On these scales one expects a variety of more complex processes to modify the galaxy clustering and thus presumably give rise to features in the correlation function $\\xi(r)$. These processes include dynamical friction, tidal interactions, stellar feedback and other dissipative processes. These processes will force close galaxy pairs to merge over time-scales of order of the dynamical time ($\\sim 10^8~\\mathrm{yr}$) \\citep{toomre77a, larson78a, vandokkum05a, bell05a}. Thus measurements of the correlation function $\\xi(r)$ on small scales can be translated to a number density of near-future merger events. Some understanding of the merger process can be used to turn this density into a galaxy merger rate or at least a constraint thereon, providing an empirical measure of the importance of mergers as a fundamental mode of stellar mass addition to galaxies.  In addition, numerical simulations within the cold dark matter paradigm find that self-bound substructures within dark matter halos (subhalos) merge to form larger systems. It is tempting to associate these subhalos with galaxies in groups and clusters \\citep{klypin99a, kravtsov99a, colin99a, moore98a}. Comparison between galaxy clustering and subhalo clustering on very small scales can lead to a better understanding of the connection between galaxies and dark matter subhalos \\citep{kravtsov04b,zentner05a}.  There are no measurements of the real-space correlation function on scales smaller than $ 100~\\kpc$. This is mainly due to the fact that on small angular separations the selection function of typical surveys becomes too complex and the pair counts are limited by shot noise. \\cite{gott79a} and \\cite{maller05a} measured the angular correlation function for angular separations that, for the median redshift of their sample translates to $10<r<100~ \\kpc$. They both find that the correlation function $\\xi(r)$ power law extends down to these small scales. The drawback of these studies is that they measure only angular, not proper correlations, and they are difficult to interpret without precise knowledge of the radial selection function and density field.  The Sloan Digital Sky Survey \\citep[SDSS;][]{york00a} has provided the largest spectroscopic sample ever of these massive galaxies; its special Luminous Red Galaxy (LRG) target selection \\cite{eisenstein01a} uses color to pre-select luminous early-type galaxies in a volume much larger than that of the SDSS Main sample. Because of their high luminosities, their association with massive halos, and their spectral uniformity, LRGs are extremely useful probes of large-scale structure.  In addition, however, their morphological simplicity makes it straightforward to study their clustering at very small scales; Their lack of cold gas (and, indeed, the lack of equally bright blue galaxies) means that close passages and mergers do not trigger significant star formation, change the colors significantly, or change the morphology significantly, and therefore do not ``remove'' systems of interest at small scales from the parent sample. In this paper we capitalize on these properties of the LRGs---and the enormous volume of the SDSS LRG sample---to measure the clustering to extremely small scales ($\\sim 10~\\kpc$) and to constrain the LRG--LRG merger rate.  Throughout this paper all distances are comoving, calculated for a cosmological world model with $(\\Omega_\\mathrm{m},\\Omega_\\Lambda)=(0.3,0.7)$ and Hubble constant parameterized by $H_0\\equiv 100~h\\,\\km\\,\\s^{-1}\\,\\Mpc^{-1}$. At the mean redshift of the sample ($z\\sim 0.3$), the fiber collision scale, $55~\\arcs$, corresponds to $\\sim 200~\\kpc$ and the SDSS median PSF width, $1.4~\\arcs$, translates to $\\sim 6~ \\kpc$. ",
        "conclusions": "We have combined the data for $24520$ luminous red galaxies from the SDSS spectroscopic sample with LRG targets in the SDSS imaging sample to measure the strength of clustering for LRGs on small scales ($0.01$ to $8\\,h^{-1}~\\Mpc$). We deal with the fiber collision incompleteness of the SDSS spectroscopic sample on scales smaller than $55~\\arcs$ by cross-correlating the spectroscopic and imaging samples and statistically removing the interloping galaxies. This method is extremely powerful on small scales at which the correlation function becomes much larger than unity.  We find that the correlation function on these scales is very close to an extrapolation of the correlation function power-law found on larger scales $\\xi(r)\\propto r^{-2}$ \\citep{zehavi05a}. This is surprising, as one might expect the direct interactions between galaxies (\\eg, dynamical friction, galaxy merger, tidal impulses, etc.) to create features in the correlation function.  This result cannot be simply explained with the current best-fit of the of Halo Occupation Distribution (HOD) models for galaxy clustering \\citep[\\eg,][]{zehavi05a,tinker05a}. This inconsistency arises from the fact that in current HOD models the galaxy-galaxy correlation function on small scales simply follows the convolution of the dark matter halo profile with itself (1-halo term). For an NFW profile \\citep{navarro97a} this relation is proportional to $r^{-1}$ toward the core of the halo and similarly for the Moore profile \\citep{moore98b} produces a $r^{-1.5}$ relation; neither is sufficiently steep to fit the results of this paper. Figure ~\\ref{fig:xi} shows the 1-halo term of the correlation function calculated for the HOD parameters given in \\cite{zehavi05b}.  If HOD models are modified to have a galaxy distribution much more concentrated than the dark matter or to have dark matter halos with density profiles much steeper than NFW toward the core of the halo, this result could in principle be accommodated within HOD models.  We convert the correlation function on these scales to statistics of galaxies in close dynamical pairs, and we infer an LRG-LRG merger rate of $\\lesssim 1/160~\\Gyr^{-1}$ or a comoving volume merger rate, $\\dot{\\phi}_M \\sim 0.6 \\times 10^4~\\Gyr^{-1}\\,\\mathrm{Gpc}^{-3}$. The fact that in large clusters, galaxies can closely pass each other without merging, turns any merger rate inferred from close dynamical pair statistics into an strict upper limit. This upper limit on the LRG merger rate could in principle be violated if the merger process triggers such strong star-formation activities in these galaxies that they become too blue to make it to the LRG sample selection cuts. However, this proposal is not tenable, as LRGs contain little apparent gas or dust and essentially there are no blue galaxies with luminosities comparable to the LRGs \\citep{blanton03d,eisenstein01a}. Therefore the LRG merger rate is in fact very close to merger rate for all galaxies with LRG luminosities regardless of their color.  In previous work, there are two main methods for measuring merger rates, the most popular one is to convert the close pair statistics to a merger rate \\citep[\\eg,][]{zepf89a, burkey94a, patton97a, patton00a, carlberg00clust,bell05a}. The other method uses the asymmetry parameter ($A$) to determine the fraction of galaxies undergoing mergers \\citep{conselice03a, conselice03b, vandokkum05a}. It is difficult to compare these previous studies of the merger rate with that found here because our LRG sample is more luminous and has a significantly smaller number density than the samples uses in those studies. The merger rate has been found to be higher for higher redshifts and higher for fainter samples \\citep{conselice03a}. For this reason, an SDSS-size sample was required to measure the merger rate for LRG type galaxies at low redshifts. Nevertheless our result is, in principal, consistent with the extrapolation of the best fits to samples at higher redshifts or fainter samples \\citep{conselice05a}.  One question to address by merger-rate measurements is the importance of major mergers in the mass build-up of massive galaxies. This could be answered by measuring the cross-correlation of galaxies of different mass with LRGs on small scales to obtain a mass spectrum for objects merging into LRGs. \\cite{murali02a} created a simulation for $L^{*}$ galaxies to answer this question. They find that smooth accretion plays the dominant role in mass build-up. Unfortunately the resolution of the simulation ($L\\sim L^{*}/4$) was too close to the galaxies in question to be able to distinguish true smooth accretion from mergers with objects below the mass resolution. LRGs are massive enough that the merger of objects far below their mass could be studied in hydro-dynamical simulations to give a detailed mass spectrum for merger events."
    },
    "0512/astro-ph0512216_arXiv.txt": {
        "abstract": "We study eigenmodes of acoustic oscillations of high multipolarity $l \\sim 100$ -- $1000$ and high frequency ($\\sim 100$ kHz), localized in neutron star envelopes. We show that the oscillation problem is self-similar. Once the oscillation spectrum is calculated for a given equation of state (EOS) in the envelope and given stellar mass $M$ and radius $R$, it can be rescaled to a star with any $M$ and $R$ (but the same EOS in the envelope). For $l\\gtrsim 300$ the modes can be subdivided into the outer and inner ones. The outer modes are mainly localized in the outer envelope. The inner modes are mostly localized near the neutron drip point, being associated with the softening of the EOS after the neutron drip. We calculate oscillation spectra for the EOSs of cold-catalyzed and accreted matter and show that the spectra of the inner modes are essentially different. A detection and identification of high-frequency pressure modes would allow one to infer $M$ and $R$ and determine also the EOS in the envelope (accreted or ground-state) providing thus a new and simple method to explore the main parameters and internal structure of neutron stars. ",
        "introduction": "Neutron stars can be considered as resonators where various oscillation modes can be excited. These oscillations are attracting much attention because, in principle, they can be used to study the internal structure of neutron stars. Some of them (for instance, r-modes) can be accompanied by gravitational radiation. Because neutron stars are relativistic objects, their oscillations must be studied in the framework of General Relativity. The relativistic theory of oscillations was developed in a series of papers by Thorne and coauthors \\citep{ThornPaperI,ThornPaperII,ThornPaperIII,ThornPaperIV,ThornPaperV,ThornPaperVI}. In particular, the rapid ($\\sim 1$ s) damping of p-modes with multipolarity $l = 2$ by gravitational radiation was demonstrated by \\citet{ThornPaperIII}. An exact treatment of general-relativistic effects is complicated, but in many cases it is possible to use the relativistic Cowling approximation \\citep{McDermott1983}. An analysis of various oscillation modes and mechanisms for their dissipation was carried out by \\citet{McDermott1988}. Let us also note the review paper by \\citet{sterg03}, which contains an extensive bibliography. As a rule, one considers neutron star oscillations with low values of $l$.  Although neutron stars are objects at the final stage of stellar evolution, they can be seismically active for many reasons. Possible mechanisms for the generation of oscillations have been widely discussed in the literature (see, e.g., \\citet{McDermott1988,sterg03} and references therein). Recently, much attention has been paid to r-modes -- vortex oscillations that can be generated in rapidly rotating neutron stars and accompanied by powerful gravitational radiation. In addition, oscillations can be excited in neutron stars, for example, during X-ray bursts (nuclear explosions in outermost layers of accreting neutron stars), bursting activity of magnetars (anomalous X-ray pulsars and soft gamma-ray repeaters; see, e.g., \\citet{kaspi04}), and glitches (sudden changes of spin periods) of ordinary pulsars.   In this paper we focus of high-frequency ($\\sim 100$ kHz) pressure oscillations (p-modes) with high multipolarity $l \\gtrsim 100$ localized in neutron star envelopes (crusts). In our previous paper \\citep{Chug2005} we have studied these oscillations for $l \\gtrsim 500$. In that case p-modes are localized in the outer envelope (before the neutron drip point, at densities $\\rho \\lesssim 4 \\times 10^{11}$ g~cm$^{-3}$), where the equation of state (EOS) of stelar matter is relatively smooth. Accordingly, the oscillation spectrum is simple and well established.  In the present paper we extend our analysis to p-modes with lower $l$. These oscillations penetrate into the inner envelope of the star, where the EOS undergoes considerable softening due to neutronization and becomes more complicated (essentially different for ground-state and accreted matter). We show that the neutron drip affects strongly the oscillation spectrum. If detected, this spectrum would give valuable information on the EOS in neutron star envelopes and also on global parameters of neutron stars (their masses and radii). ",
        "conclusions": "We have studied high-frequency pressure oscillations which are localized in the envelopes of neutron stars composed of the accreted or ground-state matter.  Our main conclusions are as follows.  (1) The oscillations are almost insensitive to various modifications of the EOS for the ground-state matter (section \\ref{subsec_EigenFrec}). All EOSs we have used (HP, RHS, SCM) give the same oscillation spectrum.  (2) The neutron drip and associated softening of the EOS in the inner envelope do not affect strongly the spectrum of the well known (outer) oscillation modes which are localized predominantly in the outer envelope (section \\ref{subsec_EigenFrec}).  (3) However, the neutron drip leads to the appearance of inner oscillation modes localized mostly near the neutron drip point (section \\ref{SecRes}). The spectrum of these modes is sensitive to the EOS in the envelope (accreted or ground-state).  (4) The p-mode oscillation problem is self-similar (in the plane-parallel approximation). Once the problem is solved for one stellar model, it can easily be rescaled to neutron star models with any mass and radius (but the same EOS in the envelope; see section \\ref{SubSecModeScale}).  (5) A detection and identification of one outer mode and several inner modes would enable one to discriminate between the ground-state and accreted envelope and determine neutron star mass and radius (section \\ref{sect:infer}). For example, a detection of the fundamental mode with $l=900$ at the frequency 74 kHz and of two inner modes with $l=300$ and $l=900$ at 56 kHz and 140 kHz, respectively, would indicate a canonical neutron star with the ground-state envelope.  Therefore, high-frequency pressure modes are excellent tools to explore the physics of matter in neutron star envelopes and to determine masses and radii of neutron stars. The oscillation frequencies could be detected by radio-astronomical methods very precisely. A detailed analysis of pulse shapes of some radio pulsars reveals that oscillations with large multipolarity % are possibly excited there \\citep{Clemens2004} but their frequencies are $\\sim$30 Hz, so that they are not high-frequency p-modes we discuss here.  A search for high-frequency p-modes could be useful. High-multipolarity p-modes do not damp very quickly because they do not produce any powerful gravitational or electromagnetic emission \\citep[see, e.g.,][]{Chug2005}. They are robust because they are relatively independent of the thermal state of the star, and they should not be strongly affected by neutron star magnetic fields. The inner p-modes, localized in the inner envelope, could be easily triggered by pulsar glitches, which are thought to occur just in inner envelopes of pulsars. \\citet{Chug2005} studied the dissipation of p-modes localized in the outer envelope; this dissipation is mainly produced by the shear viscosity. It may be enhanced  by thin viscous layers near numerous nuclear phase transitions. The viscosity in these layers can be diffusive or turbulent. Note, that fundamental modes do not produce viscosity layers because they pass phase transitions without velocity discontinuities (see Sec.\\ \\ref{SecFormal}). Their dissipation is not enhanced by phase transitions.   Finally, p-modes in neutron star envelopes are relatively insensitive to the EOS and composition of neutron star cores. However, these modes can be useful to discriminate between ordinary neutron stars and strange stars with crust. The latter stars are thought to contain extended cores composed of strange quark matter. Nevertheless, a core is assumed to be surrounded by an envelope of normal matter \\citep[see, e.g.,][]{Zdunik2002}, so that a strange star with the crust may look like an ordinary neutron star from outside. The density of the normal matter in a strange star does not exceed the neutron drip density. The pressure modes in the envelopes of such stars should easily `feel' underlying dense quark matter, and the oscillation spectrum would reflect the presence of the quark core. We intend to consider this effect in a future publication."
    },
    "0512/astro-ph0512020_arXiv.txt": {
        "abstract": "We studied the WMAP temperature anisotropy data using two different methods. The derived signal gradient maps show regions with low mean gradients in structures near the ecliptic poles and higher gradient values in the wide ecliptic equatorial zone, being the result of non-uniform observational time sky coverage. We show that the distinct observational time pattern present in the raw (cleaned) data leaves also its imprints on the composite CMB maps. Next, studying distribution of the signal dispersion we show that the north-south asymmetry of the WMAP signal diminishes with galactic altitude, confirming the earlier conclusions that it possibly reveals galactic foreground effects. As based on these results, one can suspect that the  instrumental noise sky distribution and non-removed foregrounds can have affected some of the analyses of the CMB signal. We show that actually the different characteristic axes of the CMB sky distribution derived by numerous authors are preferentially oriented towards some distinguished regions on the sky, defined by the observational time pattern and the galactic plane orientation. ",
        "introduction": "The first-year data from the Wilkinson Microwave Anisotropy Probe (WMAP) (Bennett at al. 2003a,b) became a basis for testing cosmological scenarios, substantial improving of evaluation accuracy of basic cosmological parameters related to the geometry of our Universe, dynamics of its expansion, matter and energy contents. Therefore any studies of reliability of the provided data are of great interest and importance. In particular, application of non-standard methods provide additional independent tests of data quality and consistency, as well as the correctness of involved interpretations.  All the sky CMB anisotropy maps give possibility to test with high accuracy the fundamental principles of current cosmology, such as the isotropy of our Universe and the Gaussianity of primordial cosmological perturbations. The first analysis of WMAP data for Gaussianity was performed by WMAP team \\cite{komatsu2003}. They found that the WMAP data are consistent with the Gaussian primordial fluctuations and provided limits to the amplitude of non-Gaussian ones. Other authors, however,  applying advanced or sophisticated statistical methods to analyse the foreground-subtracted WMAP maps pointed out  non-Gaussian signatures or north-south asymmetry of galactic hemispheres (Chiang et al. 2003, Park 2004, Eriksen et al. 2004a, Hansen et al. 2004a,b,c, Eriksen et al. 2005a, Cayon et al. 2005). Some authors suppose that it may be caused by residual foreground contamination from unknown galactic or extra-galactic  sources and find their possible locations, angular scales, and amplitudes (Park 2004, Eriksen et al. 2004a, Hansen et al. 2004b, Chiang \\& Naselsky 2004, Patanchon et al. 2004, Naselsky et al. 2005, Tojeiro et al. 2005, Land \\& Magueijo 2005). For example, Tojeiro et al. (2005) argued that the evidence for non-Gaussianity on large scales is associated with cold spots of unsubtracted  foregrounds. Some part of non-Gaussian  statistics of $\\Delta T/T$ fluctuations can be due to the ring of pixels of radius of 5 degrees centered at ($l\\approx 209^o, b\\approx -57^o$) \\cite{cayon2005}, close to the spot detected in Vielva et al. (2004) and Cruz et al. (2005). Wibig \\& Wolfendale (2005) performed the correlation analysis of CMB temperature and gamma-ray whole-sky maps and detected positive correlation, which may suggest an actual contamination of the WMAP data by, if not cosmic rays directly, some component accompanying them in interstellar space with a sufficiently flat emission spectrum.  All these studies indicate that WMAP Q-, V-, and W-band maps contain information on unidentified extended low-amplitude foreground structures that contaminate the relic CMB signal. The effective extraction of these structures from the WMAP maps would improve the current and future CMB analyses, as well as reveal possible new foreground objects or effects (cf. Davies et al. 2005, Eriksen et al. 2005b).  Our principal goal in the present paper is to understand the physical meaning of structures within WMAP measurements in order to detect or limit possible non-cosmological factors affecting the data used for determining cosmological parameters. We apply two simple analytical methods to study possible foreground or observational effects in the WMAP first-year data, not recognized and/or p removed by the cleaning methods used by the WMAP team (Bennett et al. 2003b). In the paper, first we shortly provide the basic information about the data analysed, and next describe our method of generating gradient maps from the foreground-cleaned WMAP maps and interpret the large-scale gradient structures by the observational strategy of the WMAP satellite. We demonstrate that the scanning effects present in the raw signal data can leave clear imprints on the gradients maps from all frequency channels, retained even in coadded, smoothed, and ILC maps. Then, we investigate distribution of signal dispersion in the WMAP maps to find a significant asymmetry between the northern and southern galactic hemispheres. Following some previous authors, we argue that a natural explanation of this asymmetry is residual galactic contamination. Finally, the inevitable influence of the above-mentioned effects on the analysis of the large-scale CMB structure is shortly considered.  \\begin{figure*}                                                 % \\includegraphics[width=55mm,height=84mm,angle=270]{fig1a.ps} \\includegraphics[width=55mm,height=84mm,angle=270]{fig1b.ps} \\includegraphics[width=55mm,height=84mm,angle=270]{fig1c.ps} \\includegraphics[width=55mm,height=84mm,angle=270]{fig1d.ps} \\caption{ The maps in the Mollveide projection of the galactic coordinate frame present: a) the signal $\\Delta T$ in the WMAP V2 channel; b) the V2 gradient map derived from (a), blue color denotes low gradient values and red color - higher ones; c) the V2 gradient map smoothed out using the $1\\degr\\,$ Gaussian beam; d) the effective number of observation map for the channel V2; blue color denotes small numbers and red color - larger ones. All maps are presented in the histogram-scale format to enhance the apparent features. The structures in the galactic plane have roughly about ten times higher gradients than those outside it. } \\label{fig1} \\end{figure*} ",
        "conclusions": "The gradient analysis is a very efficient method in recognition of weak signal non-uniformities in the sky maps. By analysing the gradient maps we demonstrate the WMAP noise pattern resulting from non-uniformity of the observational time distribution. The pattern involves regions of low noise (small gradients) near the ecliptic poles and along the circles at the angular distance of $\\approx 23^\\circ$ from the ecliptic poles, and non-uniform extended regions of larger noise amplitudes (larger gradients) along the ecliptic equator. The pattern introduced by local variations in the observational time is much alike at all available frequency channels and must inevitably affect the resulting CMB maps, independently of the applied approach to constructing such maps from the original data (e.g. Bennett at al. 2003b, Eriksen et al. 2004b, Tegmark et al. 2003). The apparent gradient structure can \"propagate'' through the noise filtering procedures used in attempt to reconstruct the CMB signal maps, as it is proved in section 3 by showing its imprints in the frequently presented and discussed ILC map. It suggests that analyses of the ILC maps, or similar maps, must be affected - besides statistical flaws due to the signal generation procedure and non-removed foreground effects - by the WMAP observational time pattern (cf. in this context e.g. Fig.~1 in Hansen et al. 2004c, Fig.~4 in Jaffe et al. 2005a and Jaffe et al. 2005b). One may also note that for the CMB study based on the WMAP data (and the future Planck mission) the best regions to study are those with the lowest gradients around the ecliptic poles, and the lowest instrumental noise.  \\begin{figure}                                                  % \\includegraphics[width=8.5cm]{fig5.ps} \\caption{ The V2 gradient map from Fig.~1c (but in a linear scale) with plotted characteristic directions described in the text. For orientation, the lines are added for the ecliptic equator and the ecliptic meridian perpendicular to the galactic plane. } \\label{fig5} \\end{figure}  Independently, the possible galactic foreground effects were studied by simple methods of analysing distributions of signal fluctuations versus the observational time at the WMAP V2 band. We confirm existence of the south-north asymmetry between hemispheres, with the larger fluctuation amplitude in the south hemisphere (cf. Eriksen et al. 2004a, Hansen et al. 2004a,b). We show that the asymmetry extends throughout a large part of the sky and diminishes at higher galactic altitudes. It suggests that the foreground effects may be due to some galactic component concentrating along the galactic plane. Furthermore, higher amplitude of signal fluctuations in the southern hemisphere can be, in our opinion, explained by the known fact that our Sun is situated approximately 20 pc north off this plane. Let us also note that regularity of observed effects with growing galactic altitude, also in the separate ranges of $N$ at Fig.~4, does not support the hypothesis suggesting an extragalactic origin of the asymmetry.  The findings suggest that the analysis of large-scale patterns of the CMB signal based on the WMAP data must be affected by observational time non-uniformities, which are not possible to be quite removed. Thus we expect that analysis results of the global CMB sky distribution can be also influenced by this fact, introducing the ecliptic symmetry to the data, and by the foreground effects, leading to some symmetry/asymmetry with respect to the galactic plane. Thus let us look at examples of characteristic orientations/axes of the universal CMB distribution discovered by the other authors (e.g. in discussion of non-trivial universe topology by de Oliveira-Costa et al. 2004, Weeks et al. 2004, Roukema et al. 2004\\footnote{but Cornish et al. (2004) do not confirm non-trivial cosmological effects}). Such 10 directions imposed onto the gradient map in Fig.~5 include: 1.)  ($l=207^\\circ$, $b=-59^\\circ$) Vielva et al. (2004); 2) (222$^\\circ$, -62$^\\circ$) Jaffe et al. (2005a); 3) (-100$^\\circ$, 60$^\\circ$) Land et al. (2005); 4) (-110$^\\circ$, 60$^\\circ$) de Oliveira et al. (2004); 5) (252$^\\circ$ , 65$^\\circ$); 6) (51$^\\circ$, 51$^\\circ$); 7) (144$^\\circ$, 38$^\\circ$); 8) (271$^\\circ$, 3$^\\circ$); 9) (207$^\\circ$, 10$^\\circ$); 10) (332$^\\circ$, 25$^\\circ$) Roukema et al. (2005). The presented directions show more or less preferential orientation with respect to four circles: the ecliptic equator and its meridian perpendicular to the galactic plane, plus two small circles in the distance of $23^\\circ$ around the ecliptic poles, the sites corresponding to the local maxima of the observational time distribution. Also the discussed unexplained quadrupole and octopole orientations in respect to the ecliptic frame (cf. de Oliveira-Costa et al. 2004, Land \\& Magueijo 2005, Jaffe et al. 2005) can be possibly explained in a natural way, if their detections are affected by the sky pattern due to the observational time, retained in the cleaned maps.  In the present work we were not able to verify the supposition that several detected structures in the CMB maps result from the discussed instrumental or foreground effects, but the distinct grouping of the observed structure orientations within some regions of the ecliptic and galactic reference frames suggests that such possible effects should be properly considered.  \\thanks{This work made use of the {\\it WMAP} data archive and the HEALPIX software package. We are grateful to Reinhard Schlickeiser and Andrzej Woszczyna for stimulating discussions. The work was performed within the research program of the Jagiellonian Center of Astrophysics.}"
    },
    "0512/astro-ph0512093_arXiv.txt": {
        "abstract": "We demonstrate that artificial bipolar structure can be detected using spectro-astrometry when the point spread function (PSF) of a point source suffers distortion in a relatively wide slit. Spectro-Astrometry is a technique which allows us to probe the spatial structure of astronomical sources on milliarcseond (mas) scales making it possible to detect close binaries and to study the geometry and kinematics of outflowing gas on scales much smaller than the seeing or the diffraction limit of the telescope. It is demonstrated that distortion of the PSF, caused by tracking errors of the telescope or unstable active optics during an exposure can induce artificial signals which may be misinterpreted as a real spectro-astrometric signal.  Using simulations we show that these may be minimised by using a narrow slit relative to the seeing. Spectra should be obtained at anti-parallel slit position angles (e.g., $0^{\\circ}$ and $180^{\\circ}$) for comparison in order to allow artificial signatures to be identified. ",
        "introduction": "Spatial resolutions of optical-IR facilities have dramatically improved over the last decade.  Indeed, the \\emph{Hubble Space Telescope} and adaptive optics on ground based 8-10m telescopes have provided spatial resolutions of 0.05 arcseconds, leading to a better understanding of the nature of a variety of astronomical objects. Despite this, even higher spatial resolutions are desired to study the geometry of extrasolar planetary systems, populations and formation of close binaries, mechanisms of mass ejection and accretion in young and evolved stars and the nature of active galactic nuclei. Optical-IR interferometry has begun to provide resolutions on milliarcsecond (mas) scales, however its applicability is limited due to the complexity of the technique.  Spectro-astrometry is an alternative approach to study milliarcsecond structures at optical-IR wavelengths.  The basic concept is to measure the relative position of the source as a function of wavelength using either an intensity weighted centroiding algorithm or profile fitting (see Bailey 1998a for a more detailed description). The resultant ``position spectrum'' will show structure on any spectral feature that is displaced from the centroid of the continuum source: e.g., an emission line arising in a binary companion, outflowing jets or any other structure that is not perfectly symmetric about the centroidal source, or whose displacement from the source varies with wavelength. The smallest spatial scale observed is limited by the seeing or the diffraction limit of the telescope, however this technique indeed provides astronomically useful information at spatial scales much smaller than these.  The concept was introduced in the 1980's as ``Differential Speckle Interferometry'' or ``Chromatic Position Difference'' (e.g.,Beckers 1982), and demonstrated to detect close binaries (e.g., Sorokin \\& Tokovinin 1985).  This broad band technique was limited by the necessity to correct for atmospheric dispersion which can be much larger than the observed structure.  This difficulty is removed by measuring the displacement of a spectral line at high resolution. Aime et al. (1988) were able to measure spectro-astrometric structure using special instrumentation. See Bailey (1998a) for review of these works.  More recently Bailey (1998a) adapted this technique using a standard long-slit CCD spectrograph and coined the term ``spectro-astrometry''.  The method is similar to that used by Solf \\& B\\\"ohm (1993), who studied a jet from a young stellar object (YSO) on subarcsecond/arcsecond scales. Bailey (1998a) revised this method and has shown that with good spatial uniformity of the CCD coupled with a pixel scale that allows excellent sampling of the seing profile positional accuracies as small as 1mas are achieved (see Takami 2003). Bailey has also shown that the technique is a powerful tool for discovering binary companions toward young and evolved stars, separating their spectra, and studying the kinematics of YSO jets/winds and narrow line regions of active galactic nuclei (AGN). Since this method doesn't require any special instrumentation, its popularity has steadily increased in recent years. Studies made for young binaries and YSO jets include Bailey(1998b), Takami et.al. (2001,2002,2003), Davis et al. (2001,2003), Whelan, Ray \\& Davis(2004) and Baines et al. (2004). Measurements using an integral field unit (IFU) were first conducted by Garcia, Thi\\'ebaut \\& Bacon (1999).  A similar approach using VLTI-AMBER is investigated by Petrov (2003), Marconi et al. (2003) and Dougadous et al. (2003). Since the positional accuracy achieved depends on the PSF size as well as the photon noise, using an interferometer for this approach has the potential to achieve information on spatial structures at \\emph{micro}arcsecond scales (Bailey 1998a).  Instrumental effects which may compromise spectro-astrometric observations have not been fully investigated. These effects may indeed create false detections of close binaries or bipolar outflows. Possible instrumental effects include; misalignment of the spectrum with CCD columns, any departure of the CCD pixels from a regular grid, imperfect flatfielding or charge transfer deficiencies in the CCD (Bailey 1998b; Takami et al. 2001). To eliminate instrumental effects, Bailey (1998a,b) proposed to obtain position spectra at anti-parallel slit position angles (e.g.,0$^{\\circ}$ and $180^{\\circ}$) via rotation of the instrument. On comparison of these spectra, any real signal from the object will change sign, while the instrumental effects remain constant.  We can thus eliminate instrumental effects by subtracting one position spectrum from the other.  This method is followed by Takami et al. (2001,2002,2003) and Baines (2004). Garcia et al. (1999), Davis et al. (2001,2003) and Whelan et al. (2004) however eliminate instrumental effects by polynomial fitting of the continuum position.  Bailey (1998a) reports a systematic effect upon observation of sharp unresolved lines.  This is caused by either telescope tracking errors or unstable active optics during the exposure, and appears in the position spectrum as a P Cygni type profile. Bailey warns that the target spectral lines should be well resolved to observe true positional displacement.  This instrumental effect could be eliminated in some cases by Bailey's method, but this cannot be done by polynomial fitting of the continuum position.  In this paper we show that this instrumental effect can appear even if the line profile is fully resolved as was the case in position spectra of RU Lupi, obtained using the 3.5m New Technology Telescope (NTT). Takami et al. had previously detected the presence of bipolar outflow in this object using the 3.9m Anglo-Australian Telescope (AAT). Indeed the observed positional displacement in the NTT looks similar to that seen in the AAT data.  We show that the signal detected in the NTT data is not real, motivating us to investigate the effect of uneven illumination on spectro-astrometry, how to avoid it and whether previously detected bipolar outflows are in fact real.  In \\S2 we describe the spectro-astrometric observations obtained using EMMI on the 3.6m NTT and the results in which artificial signatures were detected. These artefacts are identified by qualitatively inspecting position spectra obtained using anti-parallel slit position angles. In \\S3 we (1) show that the observed position spectra are explained by distortion/motion of the stellar image in a relatively wide slit, (2) show that the extent of the artefacts, in a given set of observing conditions, can be greatly reduced by using a slit width narrow in comparison to the seeing, and (3) show an example of the effect these artefacts have on real spectro-astrometric data. In \\S4 we give some recommendations on how to monitor the data for the false signatures of bipolar structure. In Appendix A, we show that the bipolar outflows reported by Takami et al. (2001,2003) are real and not due to the instrumental effect described here.  \\section[]{Observations and Results} \\begin{figure*} \\centering \\includegraphics[angle=-90,width=1.0\\textwidth]{Fig1_ref.eps} \\caption{Intensity and position spectra of RU Lup obtained using EMMI-NTT.  From left to right;  Top:  H$\\alpha$ $\\lambda6563$\\AA{}, HeI $\\lambda 6678$\\AA{} \\& [SII] $\\lambda 6731$\\AA{} emission and LiI $\\lambda 6708$\\AA{} absorption lines. The intensities shown are arbitrary; Lower four panels:  spectro-astrometric signal (mas) for each corresponding spectral feature detected at 4 position angles $0^{\\circ},180^{\\circ},90^{\\circ}\\&270^{\\circ}$ in the spectrum of RU Lup. The bottom two panels show position spectra for anti-parallel slit angles $90^{\\circ}$ \\& $270^{\\circ}$ in each of the spectral features, similarly in the third and fourth panels from the bottom position spectra for anti-parallel slit angles $0^{\\circ}$ \\& $180^{\\circ}$ in each of the spectral features is shown. In the blue wing of [SII] $\\lambda 6731$\\AA{} emission feature the angular displacement is seen to change sign in corresponding anti-parallel slits, whereas displacement near zero velocity of the [SII] $\\lambda 6731$\\AA{} line (marked with a dot-dashed line), in the permitted emission lines and the LiI $\\lambda 6708$\\AA{} absorption line does not change sign in anti-parallel slits.} \\label{Fig1}% \\end{figure*} Medium resolution spectroscopy of RU Lup was obtained with the 3.5m NTT in the early morning of April 8 2004, when the seeing was measured to be 0.8\\arcsec. The observations were made using the REad Medium Dispersion (REMD) mode of the ESO Multi-Mode Instrument (EMMI) with the \\#6 grating and an arcsecond slit. The wavelength coverage spanned 6100 to 6800\\AA{} with resolving \\mbox{power $=5000$}.  The spectra were obtained at 4 position angles: $0^{\\circ}, 90^{\\circ}, 180^{\\circ}, 270^{\\circ}$ by rotating the instrument. HeAr lamp spectra were used to calibrate the wavelength scale and a pixel scale of 0.165\\arcsec{} with the MIT detector provided good sampling of the seeing profile.  Flat fields were obtained by combining many exposures of the spectrograph illuminated by a halogen lamp.  The data were reduced using the FIGARO package provided by Starlink, following Bailey (1998a,b) and Takami (2001,2003). After subtracting the bias level and flat-fielding, ``position spectra'' were determined by fitting the seeing profile at each wavelength with a gaussian function.  Global instrumental effects in the position spectra are removed by polynomial fitting. This allows us to see the positional displacement of any emission/absorption features at each slit angle relative to the continuum.  In addition to the position spectra, the intensity spectra were obtained by subtracting the bias, flat-fielding, subtracting the adjacent sky and extracting bright columns on the CCD.  The spectrum of RU Lup shows a wealth of emission lines as observed by Lago \\& Penston (1982) and Takami (2001).  The intensity and position spectra of the H$\\alpha$ $\\lambda 6563$\\AA{}, HeI $\\lambda 6678$\\AA{} \\& [SII] $\\lambda 6731$\\AA{} emission lines and LiI $\\lambda 6708$\\AA{} absorption line, obtained using EMMI-NTT, are shown in Figure 1.  The H$\\alpha$ emission line shows a large offset in the position spectrum, $\\sim$25mas from the continuum position and the blue and red shifted wings are displaced in different directions. Similar displacement in H$\\alpha$ is reported in Takami (2001). Such structure usually indicates the presence of a bipolar outflow (see Appendix A). However a qualitative inspection of position spectra, obtained using anti-parallel slits (i.e. $0^{\\circ} \\& 180^{\\circ}$) in this data, reveals that the displacement seen in anti-parallel position spectra have equal sign indicating that the signature must be systematic. In addition, the displacement observed with slit position angle $0^{\\circ}$ is less than that observed with position angle $180^{\\circ}$. In fact two position spectra observed at a position angle of $0^{\\circ}$ exhibit dissimilar profiles indicating that the effect responsible for the spurious displacement signal is time variable and as such cannot be removed even by obtaining position spectra at anti-parallel slit angles. \\begin{figure} \\includegraphics[width=0.5\\textwidth]{Fig2.eps}\\\\ \\caption{Image capture in slit viewer 7/4/2004 using 3.5m NTT.  Pixel scale: 0.4\\arcsec. Left and Centre: distorted image of the Classical T Tauri Star (CTTS) SZ68. Right: distorted image of the CTTS SZ69}\\label{fig2} \\end{figure}  In addition to H$\\alpha$ $\\lambda 6563$\\AA{} and HeI $\\lambda 6678$\\AA{}, evidence for bipolar structure is seen in other bright permitted lines including \\mbox{FeII $\\lambda 6238/6248/6456/6516$\\AA{}}.  The LiI $\\lambda 6708$\\AA{} absorption line exhibits displacement similar to that of H$\\alpha$ and the bright permitted emission lines, although in any given spectrum the direction of displacement in Li I absorption is of opposite sign to that of the emission lines.  For example, at a position angle of $180^{\\circ}$ the blue and redshifted components of the bright permitted emission lines show positive and negative displacement respectively, whereas the Li I absorption line shows negative and positive displacement in the blue and red wings respectively.  Conversely, the direction of displacement due to blueshifted forbidden line emission reverses sign at anti-parallel slit positions ($0^{\\circ}$ vs. $180^{\\circ}$ ; $90^{\\circ}$ vs. $270^{\\circ}$). This is best seen in the [SII] $\\lambda 6731$\\AA{} emission line in Figure 1. However a large, sharp displacement is observed close to zero velocity in the position spectrum at both [SII] $\\lambda 6731$\\AA{} and [OI]\\protect\\footnote{not shown as the effect is in this line is the same as that in the [SII] line, but the data at the wavelength of [OI] is of lower quality} $\\lambda 6300$\\AA{} emission lines. This displacement is observed in all the position spectra for [SII] $\\lambda 6731$\\AA{} (marked in Figure 1 with a dot-dashed line) and [OI] $\\lambda 6300$\\AA{}. This sharp displacement does not reverse sign in anti-parallel slit position angles and therefore this is likely to be due to instrumental effects and not to actual signal from the object. ",
        "conclusions": "Artificial spectro-astrometric signal can be produced in a position spectrum as a result of a distorted stellar image in a relatively wide slit, as was the case in data obtained using the 3.6m NTT. The observed artificial displacement looks similar to that produced when observing an object exhibiting bipolar structure, the only difference being that a real signal from bipolar structure will change sign when observed using anti-parallel slit position angles whereas the artificial signal will not.  Time variable distortion of the stellar image is a significant problem, the time variation seen is presumably due to variation of the PSF, and as such the resulting magnitude of the effect is time-averaged. As a result the artefacts cannot be removed by subtracting position spectra obtained using long slit spectra with anti-parallel slit position angles, as is the case for other instrumental effects. Spectro-astrometric data must therefore be carefully monitored during observations to determine when these artefacts arise and the observations adjusted to compensate.  Simulations show that the use of a slit width narrow in comparison to the seeing greatly reduces the extent of this instrumental effect.  However, an image of high quality and good tracking of the telescope are mandatory in order to obtain position spectra with high positional accuracy.  Based on observations and simulations we recommend that in order to minimise these artefacts and obtain position spectra with high positional accuracy the following are taken into account: \\begin{enumerate} \\item When observing use a slit width narrow in comparison to the seeing. This systematic effect was not observed in observations where the seeing is typically 1.5\\arcsec{}     - 2.0\\arcsec{} and the slit width used was 1.0\\arcsec. \\item The use of an integral field unit (IFU) may be advantageous as they do not experience uneven illumination (e.g. IFU's with micro lenses or fibres), however although previous work by Garcia et al. (1999) appear to have been successful, instrumental effects arising in IFU's may well be more complicated than those from long slit spectrographs.  These should be investigated. \\item Obtain spectra with a wide spectral range.  The systematic effect will create the same artificial signal in a number of emission and absorption features.  Therefore if a wide spectral range is observed to include a number of absorption and emission lines with a range of excitations it will be possible to deduce if the bipolar signal is real. \\item Spectra should be obtained at anti-parallel slit position angles.  Not all instrumental effects are well understood or anticipated, as such comparison of position spectra obtained using anti-parallel slit position angles is the best method to determine whether or not the signal detected is due to real structure in the source. \\end{enumerate}"
    },
    "0512/astro-ph0512600_arXiv.txt": {
        "abstract": "We present spectral results, from {\\it Chandra} and {\\it XMM-Newton} observations, of a sample of 22 low-redshift ($z < 0.1$) radio galaxies, and consider whether the core emission originates from the base of a relativistic jet, an accretion flow, or contains contributions from both. We find correlations between the unabsorbed X-ray, radio, and optical fluxes and luminosities of FRI-type radio-galaxy cores, implying a common origin in the form of a jet. On the other hand, we find that the X-ray spectra of FRII-type radio-galaxy cores is dominated by absorbed emission, with $N_{\\rm H} \\gappeq 10^{23}$ atoms cm$^{-2}$, that is likely to originate in an accretion flow. We discuss several models which may account for the different nuclear properties of FRI- and FRII-type cores, and also demonstrate that both heavily obscured, accretion-related, and unobscured, jet-related components may be present in all radio-galaxy nuclei. Any absorbed, accretion-related, components in FRI-type galaxies have low radiative efficiencies. ",
        "introduction": "The physical origin of X-ray emission in the nuclei of radio galaxies is a topic of considerable debate. It is unclear whether the emission primarily originates in an accretion flow, or instead has an origin associated with a parsec-scale radio jet. The evidence in favor of each interpretation provides motivation to this paper, which uses spectroscopy of the X-ray cores to distinguish the two possibilities. On the one hand, it has been suggested that at least a fraction of the nuclear X-ray emission of radio galaxies has an origin at the unresolved base of a parsec-scale radio jet (e.g.,~\\citealt{fab84}). The best pieces of evidence in favor of this hypothesis are the observed correlations between the {\\it ROSAT} X-ray and VLA radio core fluxes and luminosities measured in the B2 (\\citealt{can99}) and 3CRR (\\citealt{har99}) samples, supporting a nuclear jet-related origin for at least the soft X-ray emission. \\cite{har99} demonstrated that, since the jet-generated radio emission is believed to be strongly affected by relativistic beaming, the flux--flux and luminosity--luminosity correlations would have considerable intrinsic scatter if the X-ray emission were instead to originate in an isotropic accretion flow.  Further constraints on the physical processes present in radio-galaxy nuclei come from {\\it HST} observations (\\citealt{chi99,har00}). Correlations are observed between the radio and optical luminosities of 3CR FRI-type nuclei, again supporting a common origin at the base of a parsec-scale jet. However, FRII-type nuclei show a range of behavior at optical wavelengths: those sources with weak or narrow optical line-emission lie on the (presumably jet-related) radio--optical luminosity--luminosity correlation established for the FRI-type sources (\\citealt{chi00b}). However, the optical luminosities of broad-line FRII-type sources lie significantly above this trendline, and \\cite{chi00b} interpret these sources as being oriented to the observer such that significant optical emission from an accretion disk is observed, consistent with predictions from AGN-unification schemes (e.g.,~\\citealt{ant93,up95}). Additional evidence in favor of jet-related X-ray emission in the nuclei of radio galaxies is found when considering the multiwavelength spectral energy distributions (SED) of these sources. As demonstrated by \\cite{cap02}, \\cite{chi03}, and \\cite{pel03}, the nuclear SED of certain FRI-type radio galaxies may be modeled by synchrotron and synchrotron self-Compton emission from the base of a relativistic jet.  There is, however, an alternative interpretation of the origin of nuclear X-ray emission: that it is emission from an accretion flow. Observations in favor of this model include the detection of broadened Fe K$\\alpha$ lines and short-timescale ($\\sim$ks) variability, both of which were used to argue for a physical origin in the inner regions of an accretion flow (\\citealt{gli03,gli04}). In a recent study of a sample of FRI-type radio-galaxy nuclei, \\cite{don04} showed that several sources with strong optical jet emission are not accompanied by strong X-ray emission, which may suggest a different physical origin for the two. Further, \\cite{mer03} argued that the observed correlations between the luminosities of radio and X-ray emission in radio-galaxy cores may be part of a `fundamental plane', linking the radio and X-ray emission with black hole mass in both radio-loud and radio-quiet AGN. They suggest that the X-ray emission in all these sources has an origin in the form of an accretion flow, although they cannot rule out a significant contribution from the jet to the X-ray emission of the radio-loud sources.  In addition to the debate surrounding the origin of the X-ray emission, there has been some controversy over whether physical differences exist in the accretion-flow modes of FRI- and FRII-type sources, and over the presence and r\\^{o}le of obscuring tori in these sources. It is unclear whether the \\frd observed on large (kpc--Mpc) scales (\\citealt{fr74}) is due to the primarily {\\it extrinsic} impact of the hot-gas environment on the jet propagating through it (e.g.,~\\citealt{bic94,bic95,gkw00}), or rather ensues from {\\it intrinsic} differences in the structure of the accretion flow and/or torus (see discussion by \\citealt{gkw00}). There is growing evidence (e.g.,~\\citealt{rey96,don04}) to suggest that FRI-type radio-galaxy nuclei possess radiatively inefficient, possibly advection-dominated (e.g.,~\\citealt{nar95}) accretion flows, rather than standard, geometrically thin disks. It has also been claimed from the low intrinsic absorption measured at X-ray wavelengths (\\citealt{don04}), together with the high optical-core detection rate (\\citealt{chi99}), that an obscuring torus, required by AGN-unification schemes, is {\\it absent} in FRI-type radio galaxies. However, as noted by \\cite{har02b},~\\cite{wor03b}, and \\cite{cao04}, if the optical and X-ray emission is dominated by a jet and occurs on scales larger that of a torus, then one cannot comment directly on the presence or absence of the torus using the X-ray absorption and optical reddening properties alone.  As the orientation-dependent effects of relativistic beaming and the putative obscuring torus are expected to play a large part in determining the observed properties of a radio-galaxy nucleus, it is important to select sources based on their low-frequency (and hence isotropic) emission characteristics, such as in the 3C and 3CRR samples. In this paper, we present the results of a \\Ch and \\XMM spectral analysis of a sample of the nuclei of 22 radio galaxies at $z < 0.1$, 19 of which are from the 3CRR catalogue, with the remaining sources, 3C~403, 3C~405 (Cygnus~A), and Centaurus~A, included due to their high quality spectra. The high spatial resolution of \\Ch means that it is possible to disentangle confusing kpc-scale jet emission from that of the core, while the large collecting area of \\XMM allows tight constraints to be placed on spectral parameters. Each observatory covers a sufficiently large energy range that both soft unobscured X-ray emission and hard obscured emission, possibly viewed through a torus, may be observed, providing direct tests of AGN-unification models.  This paper is organized as follows. Section 2 contains a description of the data and a summary of their analysis. The results from fitting the spectra of the nuclei of each source are presented in Section 3. In Section 4, we discuss the X-ray, radio, and optical flux and luminosity correlations. In Section 5, we discuss the presence that an obscuring torus is present in FRI-type radio galaxies, while in Section 6 we interpret the observed differences in the X-ray nuclei of FRI- and FRII-type sources. We end with our conclusions in Section 7. All results presented in this paper use a cosmology in which $\\Omega_{\\rm m, 0}$ = 0.3, $\\Omega_{\\rm \\Lambda, 0}$ = 0.7, and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. When distinguishing between different model fits to the data, we present $F$-statistic results, although we note that this method may be unreliable in such circumstances (\\citealt{pro02}). We adopt thresholds of 95 and 99.9 per cent for marginally and highly significant improvements in the fit, respectively. Errors quoted in this paper are 90 per cent confidence for one parameter of interest (i.e., $\\chi^2_{\\rm min}$ + 2.7), unless otherwise stated. ",
        "conclusions": "We have presented results from a \\Ch and \\XMM spectral analysis of the nuclei of a sample of the nuclei of 22 low-redshift ($z < 0.1$) radio galaxies.  We find that:   \\begin{enumerate}  \\item The nuclear X-ray spectra of FRI-type sources are unabsorbed, or absorbed simply by gas related to the known kpc-scale dusty disks in the host galaxy. The strong observed correlations between the X-ray, radio, and optical fluxes and luminosities imply that the emission has a common origin at the base of a relativistically beamed parsec-scale jet.  \\item The nuclear X-ray emission of narrow-line FRII-type sources is dominated by heavily absorbed components of emission with $N_{\\rm H} > 10^{23}$ atoms cm$^{-2}$, and is accompanied by emission from neutral fluorescent Fe K$\\alpha$ lines. We argue that this absorbed emission is likely to originate in an accretion flow and be surrounded by a structure such as the putative torus. We also find that the nuclear X-ray spectrum of every FRII galaxy has a corresponding component of soft X-ray emission, which is consistent with having a jet-related origin.   \\item If the (jet-dominated) X-ray emission of FRI-type sources occurs on scales larger than the torus, it is impossible to test for the presence of a torus using the X-ray data, but important constraints can still be made. We estimate the maximum level of a `hidden', accretion-related component of emission that could be obscured by an adopted column of $10^{23}$ atoms cm$^{-2}$ to be in the range $10^{39}$--$10^{41}$ ergs s$^{-1}$. The X-ray data do not exclude the presence of a torus, but the luminosity of the accretion flow it obscures is significantly less than in FRII-type sources unless there is more obscuring matter in the FRI-type sources, which seems unlikely based on infrared constraints.  \\item Any `hidden' accretion flows in jet-dominated FRI-type sources are likely to be significantly sub-Eddington in nature. This implies that their accretion flows are mass-starved, and/or radiate at a low efficiency.  \\item The accretion-flow luminosities of FRII-type sources are typically several orders of magnitude higher than those of FRI-type sources. The ratio of X-ray to Eddington luminosities, $\\eta_{\\rm X,Edd}$, is $\\sim$$10^{-3}$--$10^{-2}$, while the ratio of bolometric to Eddington luminosities is still higher. This implies that the accretion flows of FRII-type sources tend to be fed at a high rate, and/or possess significant contributions from high-radiative-efficiency flows, plausibly in the form of a standard, geometrically thin, optically thick disk.  \\item Two models may account for the observed differences in the nuclear properties of FRI- and FRII-type sources, although neither is without problems. One model, in which the relative contribution of the accretion-related and jet-related emission varies smoothly as a function of total AGN power, can successfully account for the observed X-ray emission characteristics of these sources. However, it is then difficult to understand why the transition between a source being jet-dominated and accretion-dominated occurs at the FRI/FRII boundary. Alternatively, there is a real dichotomy in the accretion-flow modes of FRI- and FRII-type sources. We note that the manner in which the accretion-flow mode might then affect the large-scale (FRI versus FRII) characteristics of radio galaxies remains poorly understood.  \\end{enumerate}"
    },
    "0512/astro-ph0512436_arXiv.txt": {
        "abstract": "{We review four mechanisms for forming brown dwarfs: (i) turbulent fragmentation (producing very low-mass prestellar cores); (ii) gravitational instabilities in discs; (iii) dynamical ejection of stellar embryos from their placental cores; and (iv) photo-erosion of pre-existing cores in HII regions. We argue (a) that these are simply the mechanisms of {\\it low-mass star formation}, and (b) that they are not mutually exclusive. If, as seems possible, all four mechanisms operate in nature, their relative importance may eventually be constrained by their ability to reproduce the binary statistics of brown dwarfs, but this will require fully 3-D radiative magneto-hydrodynamic simulations. ",
        "introduction": "The existence of brown dwarfs was first proposed on theoretical grounds by Kumar (1963) and by Hayashi \\& Nakano (1963). However, more than three decades then passed before brown dwarfs were observed unambiguously (Rebolo et al., 1995; Nakajima et al., 1995; Oppenheimer et al. 1995). Brown dwarfs are now observed routinely, and are estimated to be comparable in number with hydrogen-burning stars. It is therefore appropriate to ask how brown dwarfs form, and in particular to ascertain (a) whether brown dwarfs form in the same way as hydrogen-burning stars, and (b) whether there is a clear distinction between the mechanisms that produce brown dwarfs and those that produce planets.  In Section \\ref{SEC:STAR} we argue that brown dwarfs do form in the same way as stars, on the grounds that their statistical properties (mass function, binary statistics, clustering properties, etc.) appear to form a smooth continuum with those of low-mass hydrogen-burning stars. We also suggest that understanding how brown dwarfs form is the key to answering a fundamental anthropic question, namely, what determines the lower mass limit for star formation, and thereby the likelihood of long-lived stars with habitable zones. In Section \\ref{SEC:TURB} we consider the formation of brown dwarfs by turbulent fragmentation, as suggested by Padoan \\& Nordlund (2002), and we address the question of whether an isolated core of brown-dwarf mass formed in this way can cool sufficiently fast to condense out. In Section \\ref{SEC:DISC} we consider the formation of brown dwarfs by gravitational instabilities in discs. We stress that only in massive discs, and at large radii, can fragments of a disc contract and cool sufficiently fast to condense out; closer in they are likely to bounce and be shredded. We also point out that, in a dense proto-cluster, impulsive interactions between discs, or between a disc and a naked star, should be common, and may be necessary to ensure disc fragmentation. In Section \\ref{SEC:EJEC} we consider the formation of brown dwarfs by the ejection mechanism, as suggested by Reipurth \\& Clarke (2001). We point out that the requirements for this mechanism to operate are very general, and therefore it is likely to occur in nature, although it is probably not the only mechanism forming brown dwarfs, given the difficulty it has producing close BD-BD binaries. In Section \\ref{SEC:EROS} we consider the formation of brown dwarfs by photo-erosion of pre-existing cores which are overrun by HII regions, as suggested by Hester et al. (1996). We stress that this a very robust mechanism, in the sense that it does not require very fine tuning of the parameters; but it is also a very inefficient mechanism, in the sense that it requires a very massive inital core to form a brown-dwarf, and it clearly cannot deliver the brown dwarfs in regions like Taurus. In Section \\ref{SEC:CONC} we summarise our review. ",
        "conclusions": "\\label{SEC:CONC}  We have discussed four possible mechanisms for forming brown dwarfs: turbulent fragmentation, disc fragmentation, dynamical ejection and photo-erosion. None of these is mutually exclusive, and in fact the first three may occur consecutively. We emphasise that none of these mechanisms has been modelled properly with a fully radiative 3-D magneto-hydrodynamical code. Therefore, neither the thermal effects which presumably determine the minimum mass for star formation, nor the angular momentum transport processes which presumably determine the binary statistics, nor the $N$-body dynamics which presumably determine the clustering properties of brown dwarfs, has yet been properly evaluated."
    },
    "0512/hep-th0512001_arXiv.txt": {
        "abstract": "We study the possibility of having a static, asymptotically AdS black hole localized on a braneworld with matter fields, within the framework of the Randall and Sundrum scenario. We attempt to look for such a brane black hole configuration by slicing a given bulk spacetime and taking ${\\Bbb Z}_2$ symmetry about the slices. We find that such configurations are possible, and as an explicit example, we provide a family of asymptotically AdS brane black hole solutions for which both the bulk and brane metrics are {\\em regular} on and outside the black hole horizon and brane matter fields are {\\em realistic} in the sense that the dominant energy condition is satisfied. We also find that our braneworld models exhibit signature change inside the black hole horizon. ",
        "introduction": "In \\cite{RS}, using the ``squashing\" effect a negative bulk cosmological constant has on a four dimensional hypersurface, a brane, Randall and Sundrum (RS) uncovered a  mechanism of dynamical localization of gravity. All dimensions, including the one(s) we don't see, could now be infinitely large and one would still recover, on the brane, the General Relativistic and Newtonian limits of gravity at low energies \\cite{GT,Shiromizu}. It has also been shown that Standard Model matter fields can be constrained on a brane \\cite{gonzalo} and observations do not rule out braneworlds as cosmological models (see \\cite{maartcosmo} for a review). From the astrophysical point of view, both numerical and analytical models of stars have been found \\cite{toby,maarger}.   \\par On the other hand, black holes are not well understood in the RS braneworld scenario. The first attempt to find a static black hole solution on the brane was developed in \\cite{Chamblin} where the 4D Schwarzschild black hole metric on the brane was embedded in a 5D bulk containing an extended singularity: a black string. Were our universe to be a brane in a higher dimensional bulk, such a state of affair is not satisfactory: one might indeed expect astrophysical black holes formed by collapsing matter {\\it on the brane} to be localized on (or at least very close to) the brane (see however \\cite{Maa}.) Study of a simple gravitational collapse model \\cite{BGM} on a RS braneworld indicates the difficulty of finding a static vacuum black hole solution localized on a RS braneworld. The difficulties in understanding black hole solutions arise from the fact that in general, brane dynamics generate bulk Weyl curvatures, which then backreact on the brane dynamics. One is then left with the very hard task of solving equations of motion for the coupled system of bulk and braneworld with given suitable initial data. Such a program has not been answered yet, and even numerical approaches are still rather approximative \\cite{Tanaka1}. To simplify the problem, the authors of \\cite{Dadhich} looked for analytic solutions to the projected Einstein equations on the brane only and found an exact (Reissner-Nordstrom looking) black hole solution. Other similar solutions were subsequently found \\cite{Casadio1}. These bulks' geometries are not known.  \\par % Under such circumstances, it is interesting to ask whether a brane on which a 4D black hole is localized can be found by looking for a slice that intersects a bulk black hole. However, generalizing the work of \\cite{Chamblin}, Kodama showed in \\cite{Kodama} that brane solutions with a black hole geometry cannot be found as a slicing of a bulk with $G(D-2,k)$ symmetry, if the brane is {\\em vacuum} and {\\em not} totally geodesic \\footnote{$D$ is the dimensionality of the bulk spacetime and $k$ is the $D-2$ sectional curvature $k=\\pm 1,\\,0$; a brane is called totally geodesic if its extrinsic curvature is vanishing. }. In other words if, for simplicity, one wants to keep studying slices of bulks with $G(D-2,k)$ symmetry to find localized black holes in the RS braneworld scenario---in which the brane is not totally geodesic---one has to look for a {\\it non-vacuum} brane.   \\par %  Recently, an attempt to find a localized static but {\\em non-vacuum} brane black hole solution as a slice of a $G(D-2,k)$ bulk was made by Seahra~\\cite{Seahra}. There, the bulk chosen was the Schwarzschild and Schwarzschild-AdS black hole with spherical three-dimensional geometry, i.e., $G(3,k=1)$, and branes were taken as a planar, asymptotically flat slice of these bulks. Unfortunately these slicing turned out to produce naked singularities with respect to the induced geometry, except when corresponding to the equatorial plane of a bulk black hole, a special case of a totally geodesic brane.  The aim of this paper is to find a regular RS braneworld on which a static, spherically symmetric black hole surrounded by realistic matter is localized, by slicing a fixed 5D black hole bulk spacetime. The choice of slicing we will use is motivated by the AdS/CFT-inspired ``classical braneworld black hole\" vs ``quantum black hole\" duality of \\cite{Emparan} which states: ``The black hole solutions localized on the brane in the $AdS_{d+1}$ braneworld which are found by solving the classical bulk equations in $AdS_{d+1}$ with the brane boundary conditions, correspond to quantum-corrected black holes in $d$-dimensions, rather than classical ones\" (see also \\cite{Tanaka2}). Since due to Hawking radiation, black holes in asymptotically flat spacetimes are semi-classically unstable, such a duality would explain the impossibility \\cite{BGM} of finding a static exterior to the Oppenheimer-Snyder collapse of a star in asymptotically flat RS braneworlds (see also \\cite{Casadio2}). However, asymptotically AdS spacetimes allow (big enough) black holes to be in semi-classical equilibrium \\cite{HawkingPage} with their Hawking radiation. That is the main motivation for turning our attention to the specific slices we study, which are non-vacuum and asymptotically AdS (in a weak sense, see \\ref{Sect:IV}). Encouraging results were already obtained in \\cite{Emparan2}. In this paper we will indeed show with explicit examples that it is possible to construct a localized braneworld black hole surrounded by matter that satisfies the dominant energy condition, when the braneworld is asymptotically AdS (the case of an asymptotically AdS brane black hole as a slice of a bulk black string was studied in \\cite{char}).  \\par % The plan of the paper is as follows. In the next section, we show that regular slices that cross a bulk black hole horizon can be constructed and we point out why the planar slices of \\cite{Seahra} exhibit a curvature singularity there. In Sect.~\\ref{Understanding:Slicing}, we fix our notations and define a bulk slice which is an asymptotically AdS braneworld. We then consider a simple one-parameter family of slices which correspond to asymptotically vacuum and asymptotically AdS braneworlds with black hole horizon, filled with matter satisfying the dominant energy condition. Under our slicing ansatz, we find such braneworlds are possible only for the bulk with three-dimensional spatial geometries corresponding to $k=0$ and $k=-1$. For the bulk spacetimes with spherical three-dimensional geometry ($k=1$), our slicing define braneworld with matter violating the dominant energy condition. We also find that some of our braneworlds exhibit an intriguing property: signature change inside brane black holes. In the conclusion, we summarize and discuss our results. ",
        "conclusions": "\\label{conclusion}  One of the most important unsolved problems in braneworld scenario is probably the missing localized black hole solution on a braneworld. Until now, only negative results of (vacuum) localized braneworld black holes had been put forward. Motivated by the conjecture that {\\it localized black hole on a Randall-Sundrum brane} holographically corresponds to a {\\it semiclassical four dimensional black hole} \\cite{Emparan, Tanaka2}, we tried to find a non-vacuum, asymptotically AdS black hole solution on such a (non totally geodesic) brane.  To achieve this goal, we hunted for possible slices of the Schwarzschild-AdS bulk which cut the bulk black hole horizon producing a smooth horizon on the brane. We showed that this is possible if one introduces suitable matter on the brane, solely determined by the junction conditions on the brane itself. Requiring such matter to be realistic, we looked for slicing corresponding to a brane filled by matter satisfying the dominant energy conditions.  More explicitly, we studied the simplest one-parameter family of slices which obeyed these constrains. Although our parameter turned out to be constrained for our energy conditions to be satisfied (outside the horizon), we found a whole range of values for which our slices correspond to a (regular) localized ``black hole\" on a brane. In some particular cases, corresponding to a generalization of the ``equatorial slice\" of the spherical Schwarzschild AdS bulk to hyperbolic and flat three-dimensional geometries, we found explicit analytical solutions with a horizon hiding a single point-like singularity at the center. We also noticed that for non-zero large bulk black hole mass, the energy conditions for the effective matter are satisfied outside a spherical region surrounding the black hole horizon. This is reminiscent of semi-classical results obtained in asymptotically flat black hole spacetimes (see \\cite{Page,Visser} and references therein). For zero mass bulk black hole case, both the real and effective energy conditions are satisfied everywhere outside the horizon.  We also showed that it is possible for the part of our braneworlds that is hidden by the Killing horizon to undergo a signature change. From the braneworld view point, the appearance of the Euclidean signature region might be interpreted in the quantum theoretical context, such as Euclidean quantum gravity on the braneworld. On the other hand, the bulk spacetime is everywhere Lorentzian, hence one may expect that the braneworld signature change could entirely be understood in terms of bulk classical theory. Such an expectation is in accord with the spirit of a holographic idea in the sense that quantum phenomena on braneworld have some correspondence to bulk classical phenomena. However, whether such a signature changing braneworld can be realized as a solution of a well-posed initial value problem (e.g., 5D Einstein equations with suitable initial data) is a non-trivial question.  By construction, our results fit in a Randall-Sundrum braneworld scenario. We nevertheless believe that our solutions can be of phenomenological importance beyond that framework in understanding real astrophysical or microscopic black holes if extra dimensions are part of an ultimate theory of gravity.  Finally, from the perspective of the conjecture of \\cite{Emparan,Tanaka2}, we would like to conclude by pointing out that the brane black hole solutions found in this paper are, to our knowledge, neither known semi-classical solutions nor possibly obtained by perturbations thereof (see \\cite{Emparan2} for further similar comments). We found that black holes satisfying energy conditions are possible only in the case $\\Lambda_4\\simeq \\Lambda_5$. In this regime, four dimensional gravity is not localized at least at spatial infinity \\cite{karch}. Therefore, although our results apparently contradict the conjecture of \\cite{Emparan} and \\cite{Tanaka2}, it is not clear whether our black hole solutions can be used for this holographic conjecture \\footnote{We thank Takahiro Tanaka for pointing this out to us.} as, at least at spatial infinity, we would expect our black hole solutions to correspond to a deformed conformal field theory without gravity \\cite{deformed}. Clarification of this problem is beyond the scope of the present work.  \\bigskip  \\begin{center} {\\bf Acknowledgments} \\end{center} The authors are indebted with M.~Sasaki for enlightening discussions. We also wish to thank S.~Hartnoll, S.~Hawking, R.~Maartens and G.~Procopio for useful comments and discussions. AI was supported by NSF grant PHY 00-90138 to the University of Chicago. CG was supported in part by JSPS fellowship (JSPS/FF1/412) and in part by PPARC research grant (PPA/P/S/2002/00208). \\bigskip"
    },
    "0512/astro-ph0512147_arXiv.txt": {
        "abstract": "We present the rest-frame Js-band and Ks-band luminosity function (LF) of a sample of about 300 galaxies selected in the  HDF-S at Ks$\\le23$ (Vega). We use calibrated photometric redshift together with spectroscopic redshift for 25\\% of the sample. { The accuracy reached in the photometric redshift estimate is 0.06 (rms) and the fraction of outliers is 1\\%.} We find that the rest-frame Js-band luminosities obtained by extrapolating the observed Js-band photometry are consistent with those obtained by extrapolating the photometry in the redder H and Ks bands closer to the rest-frame Js, at least up to $z\\sim2$. Moreover, we find no significant differences among the luminosities obtained with different spectral libraries. Thus, our LF  estimate is not dependent either on the extrapolation made on the best-fitting template or on the library of models used to fit the photometry. The selected sample has allowed to probe the evolution of the LF in the three redshift bins [0;0.8), [0.8;1.9) and [1.9;4) centered at the median redshift $z_m\\simeq[0.6,1.2,3]$ and to probe the LF at $z_m\\simeq0.6$ down to the unprecedented faint luminosities M$_{Js}\\simeq-13$ and M$_{Ks}\\simeq-14$. { We find hints of a  raise of the faint end (M$_{Js}>-17$ and M$_{Ks}>-18$)  near-IR LF at $z_m\\sim0.6$, raise which cannot be probed at higher redshift with our sample}. The values of $\\alpha$ we estimate are consistent with the local value and do not show any trend with redshift. We do not see evidence of evolution from $z=0$ to $z_m\\sim0.6$ suggesting that the population of local bright galaxies was already formed at $z<0.8$. On the contrary, we clearly detect an evolution of the LF to $z_m\\sim1.2$ characterized by a brightening of M$^*$  and by a decline of $\\phi^*$. To $z_m\\sim1.2$ M$^*$ brightens by about 0.4-0.6 mag  and  $\\phi^*$ decreases by a factor 2-3. This trend  persists, even if at a less extent, down to $z_m\\sim3$ both in the Js-band and in the Ks-band LF. The decline of the number density of bright galaxies seen at $z>0.8$ suggests that a significant fraction of them increases their stellar mass at $1<z<2-3$ and that they underwent a strong evolution in this redshift range. On the other hand, this implies also that a significant fraction of local bright/massive galaxies was already in place at $z>3$. Thus, our results suggest that  the assembly of massive galaxies is spread over a large redshift range and that the increase of their stellar mass has been very efficient also at very high redshift at least for a fraction of them. ",
        "introduction": "The luminosity function (LF) of galaxies is a fundamental statistical tool to study the populations of galaxies. Its dependence on  morphological type, on wavelength and on look-back time provides constraints on the evolution of the properties of the whole population of galaxies, of the populations of various morphological types and on their contribution to the luminosity density at different wavelength. The parameters derived by the best-fitting of the LF, the characteristic luminosity M$^*$, the slope $\\alpha$ and the normalization $\\phi^*$, provide strong constraints to the models of galaxy formation.  The recent estimates of the LF based on local wide surveys in the optical, such as the Two Degree Field (2dF; Folkes et al. 1999; Madgwick et al. 2002; Norberg et al. 2002) and the Sloan Digital Sky Survey (SDSS; Blanton et al. 2003) and in the near-IR, such as the Two Micron All Sky Survey (2MASS; Kochanek et al. 2001), have provided a comprehensive view of the local LF for different morphological types and wavelengths. It is now well established that the LF depends on morphological type, that the faint-end is increasingly dominated by galaxies with late-type morphology and spectra and that their number density increases going to lower luminosities. On the contrary, the bright-end is dominated by early-type galaxies whose fraction increases with luminosities (e.g. Marzke et al. 1994, 1998; Folkes et al. 1999; Kochanek et al. 2001; Bell et al. 2003).  The studies of the LF based on optically selected redshift surveys at $z<1$ have provided the first evidences of a differential evolution of galaxies. Lilly et al. (1995a), using the Canada France Redshift Survey (CFRS; Lilly et al. 1995b) show that the rest-frame B-band LF of the blue population of galaxies brightens by about 1 mag to $z\\sim1$ contrary to the red population which shows a little evolution over the redshift range probed. They used the observed I-band photometry to derive the rest-frame B-band luminosities. Subsequent studies confirmed the different behaviors followed by the LF of the different populations of galaxies at various redshifts (e.g. Lin et al. 1997; Liu et al. 1998) and the differential evolution they underwent (e.g. Bell et al. 2004; Wolf et al. 2003). At higher redshift ($z>1$) the studies of the LF in the optical rest-frame confirm the presence of the bi-modality and provide evidences of luminosity and of density evolution down to $z\\sim2-3$ (e.g. Poli et al. 2003; Gabasch et al.2004; Giallongo et al. 2005).  Contrary to the light at UV and at optical wavelengths, the near-IR light is less affected by dust extinction and by ongoing star formation e.g. Rix \\& Riecke 1993; Kauffmann \\& Charlot 1998). Moreover, since the near-IR light of a galaxy is dominated by the evolved population of stars, the near-IR light is weakly dependent on galaxy type and it is more related to the stellar mass of the galaxy. Therefore, the evolution of the LF of galaxies in the near-IR rest-frame can provide important clues on the history of the stellar mass assembly in the Universe rather than on the evolution  of the star formation.  The studies of the K-band luminosity function conducted so far have shown no or little evolution of the population of galaxies at $z<0.4-0.5$ (e.g. Glazebrook et al. 1995; Feulner et al. 2003; Pozzetti et al. 2003) with respect to the local population (Kochaneck et al. 2001; Cole et al. 2001; Glazebrook et al. 1995; Gardner et al. 1997). On the contrary, evidences of evolution emerge at $z>0.5$, even if  some discrepancies are present among the results obtained by the various authors (see e.g. Cowie et al. 1996; Pozzetti et al. 2003; Drory et al. 2003; Caputi et al. 2004, 2005; Dahlen et al. 2005). One of the possible reasons of these discrepancies is related to the rest-frame K-band luminosities, usually extrapolated from the observed K-band photometry making use of the best-fitting template which should reproduce the unknown SED of the galaxy. To minimize the uncertainties, photometry at wavelengths long-wards the rest-frame K-band would be required. Given the obvious difficulties in getting deep mid-IR observations, some authors have used the observed K-band magnitudes to construct the rest-frame J-band LF (e.g. Pozzetti et al. 2003; Bolzonella et al. 2002; Dahlen et al. 2005) and to constrain the evolution of the LF in the near-IR. Another reason could be that deep K-band observations  allowing a good sampling of the LF at magnitudes fainter than M$^*$ over a wide redshift range are not so common. Consequently  the LF faint-wards of the knee is progressively less constrained with increasing redshift affecting the LF estimate.  In this work, we aim at measuring the rest-frame Js-band and Ks-band LF of galaxies and its evolution to $z\\simeq3$ through a sample of about 300 galaxies selected at Ks$\\le23$ on the Hubble Deep Field-South (HDF-S). The extremely deep near-IR observations collected on this field coupled with the optical HST data have allowed us to probe the LF over a wide redshift range and to sample the faint end of the LF down to unprecedented faint luminosities. We use photometric redshift spectroscopically calibrated with a sample of more than 230 spectroscopic redshift, $\\sim80$ of which in the HDF-S. We pay particular attention to the calculation of the rest-frame near-IR luminosities by comparing different methods and libraries of models. In \\S 2 we describe the photometric catalogue and the criteria adopted to construct the final sample of galaxies used to derive the LF. In \\S 3 we describe the procedure used to obtain calibrated photometric redshifts and we present the redshift and color distribution of galaxies. In \\S 4 we discuss the derivation of the rest-frame Js-band and Ks-band absolute magnitudes and assess the influence of the different estimates on the LF. The resulting LF at various redshifts is derived in \\S 5. In \\S 6 our LFs are compared with the local estimates and with those previously obtained by other authors at comparable redshift to depict the evolution of the LF of galaxies up to $z\\sim3$. In \\S 7 we summarize and discuss the results.  Throughout this paper, magnitudes are in the Vega system unless explicitly stated otherwise. We adopt an $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ cosmology with H$_0=70$ Km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have probed the evolution of the rest-frame Js-band and Ks-band LFs of field galaxies with a complete sample of about 300 galaxies selected in the HDF-S at Ks$\\le23$ (Vega). Photometric redshifts have been obtained from template SED fitting to U$_{300}$, B$_{450}$, V$_{606}$, I$_{814}$, Js, H and Ks photometry after having calibrated and optimized the procedure and the set of templates with a control sample of 232 spectroscopic redshifts, 151 from the HDF-N and 81 from the HDF-S. The accuracy in the redshift estimate we obtained is 0.06 (rms).  We investigate the reliability of the rest-frame near-IR absolute magnitudes obtained using the conventional method based on the extrapolation of the observed photometry on the best-fitting template and  using the photometry approaching the rest-frame near-IR. We find that the rest-frame Js-band absolute magnitudes obtained through the photometry in the redder bands (H and Ks according to the redshift of the galaxy) are consistent with those obtained from the Js-band photometry at least down to $z\\sim2$. This shows that the LF is not dependent either on the extrapolation made on the best-fitting template or on the library of models used.  We derived the Js-band and the Ks-band LF in the three redshift bins [0,0.8), [0.8,1.9) and [1.9,4) centered at the median redshift $z_m\\sim0.6$, $z_m\\sim1.2$ and $z_m\\sim3$. Each bin contains 101, 100 and 84 galaxies of which 52 (50\\%), 12 (12\\%) and 10 (12\\%) with spectroscopic redshift respectively. The analysis of the observed LF at different redshifts and the comparison with those previously found by other authors at comparable redshift and at $z=0$ provided the following results:  \\noindent 1. We find hints of a raise of the faint end (M$_{Js}>-17$ and M$_{Ks}>-18$) of the near-IR LF at $z_m\\sim0.6$. The raise is defined by 9 galaxies at $z<0.3$ with irregular morphology as found for the local LF at optical wavelength (e.g. Marzke et al. 1994; Marzke et al. 1998; Zucca et al. 1997; Folkes et al. 1999). They account for a comoving number density $\\bar n=0.4\\pm0.23$ Mpc$^{-3}$ at $<z>\\simeq0.2$, almost one order of magnitude higher than that of brighter galaxies. { However, given the low statistics, such raise of the faint end  needs a larger sample to be established;}   \\noindent 2. We find no evidence of a steepening with redshift of the near-IR LFs of galaxies in the HDF-S. The value of $\\alpha$ we find is consistent with the local value. Given the size of our sample, we estimate that if $\\alpha$ changes with redshift its evolution is constrained within $\\pm0.3$ from $z\\sim0$ to $z\\sim3$.  \\noindent 3. We do not find evidence of strong evolution of the LF up to $z_m\\simeq0.6$, where 50\\% of our galaxies has spectroscopic redshift. The comparison with the local estimates of the near-IR LF (Cole et al. 2001; Kochanek et al. 2001) shows that M$^*$ has evolved not more than 0.2-0.3 magnitude in this redshift range, while the number density of galaxies ($\\phi^*$) has not evolved significantly in this redshift range, in agreement with the previous studies (e.g. Drory et al. 2003; Pozzetti et al. 2003; Feulner et al. 2003; Glazebrook et al. 1995; Caputi et al. 2004, 2005).  \\noindent 4. We clearly detect  an evolution of the LF at  $z_m>0.6$ characterized by a brightening of M$^*$ of about 0.6 and a decrease of $\\phi^*$ by a factor 2-3 to $z_m\\sim1.2$ both in the Js-band and in the Ks-band. { By fixing $\\alpha=1.0$, this evolution  is characterized by $\\Delta M^*\\simeq-0.6$ (at 1$\\sigma$) coupled with $\\Delta\\phi^*/\\phi^*\\simeq-0.65$ (at $>3\\sigma$). The brightening  persists (at 1$\\sigma$) up to $z_m\\sim3$ together with the decline of $\\phi^*$ (at $\\sim 2\\sigma$) even if at a lower extent.}  Our results agrees, at least qualitatively, with most of the analysis previously done (Drory et al. 2003; Pozzetti et al. 2003; Feulner et al. 2003; Caputi et al. 2004, 2005) while deviates from the recent estimate of Dahlen et al. (2005). Taking into account that the near-IR light is rather tightly connected to the stellar mass, our results can  give insights on the evolution of the stellar mass of the galaxies besides their luminosity. Up to $z\\simeq0.8$, clearly little luminosity evolution and no density evolution is observed. This suggests that the population of local bright galaxies was already formed at $z\\sim0.8$ and that their stellar mass growth was already completed at this redshift. This agrees with the results on the evolution of the stellar mass density obtained by many authors: Fontana et al. (2003) and Rudnick et al. (2003) in  the HDF-S, Fontana et al. (2004) on the K20 sample (Cimatti et al. 2002), Yamada et al. (2005) in the Subaru Deep Survey Field, Bundy et al. (2005) and Drory et al. (2005) in the GOODS and FDF fields and  Feulner et al. (2005a) in the MUNICS fields. Indeed, all of them  find very little or even no evolution in the stellar mass function of galaxies and in the specific star formation rate (SSFR) at $z<1$ for high-mass galaxies.  At $z>0.8$ the evolution is stronger and the larger brightening observed is accompanied by a decrease of the number density of bright/massive galaxies. In the redshift bin [0.8,1.9) the number density of bright galaxies is $\\sim30-50$\\% of the local value and reaches $\\sim20-30$\\% at $1.9<z<4$. This decline in the number density suggests that up to $70$\\%  of the local massive/bright galaxies has grown at $1<z<2-3$ through star formation, merging or both. Thus, we should expect to observe massive star-forming galaxies with disturbed/irregular morphology and merging systems in this high redshift range, as indeed found in same cases (e.g. Daddi et al. 2004). On the other hand, the observed evolution implies also that at least 30\\% of the local bright/massive galaxies was already in place at $1<z<2$  as indeed found by some authors (Fontana et al. 2004; Glazebrook et al. 2004; F$\\ddot{\\rm o}$rster Schreiber et al. 2004). Pozzetti et al. (2003) and Caputi et al. (2005a) show that the bright end of the LF is dominated by early-type galaxies and that their contribution to the bright end does not decline with redshift. Thus, that 30\\% should be dominated by early-types implying that we should observe  massive evolved galaxies fully assembled at $1<z<2$  and even at $z>2$. This is indeed observed by many authors (Cimatti et al. 2004; McCarthy et al. 2004; Saracco et al. 2004, 2005; Longhetti et al. 2005; Daddi et al. 2005; Labb\\'e et al. 2005). It is worth noting that this picture could be accounted for, at least qualitatively, in the hierarchical picture of galaxy formation as suggested by the recent results obtained by Nagamine et al. (2005).  From these considerations we can gather that the growth of massive galaxies does not follow a unique way but displays different behaviours. A significant fraction (50-70\\%) of the brighetst/most massive galaxies increases their stellar mass over a large redshift range  at $z>1$. The remaining fraction reaches their final mass in a narrower redshift range at $z>3$ since it is already in place by this redshift. This suggests that, for the former, the stellar mass growth has been less efficient and has proceeded more slowly than for the latter. On the contrary, in the latter, the stellar mass has to be grown rapidly in a short interval, surely much shorter than 1 Gyr given their redshift. This is  supported by the recent results derived by mid-IR massive galaxies observed at high-z (e.g. Caputi et al. 2005b) and suggests a high efficiency in the accretion of the stellar mass in massive haloes in the early Universe, possibly through a very efficient star formation. The recent results obtained on the evolution of the SSFR of galaxies seems to  point toward this direction. The strong increase in the SSFR of the most massive galaxies with redshift (Feulner et al. 2005b) favors indeed an efficient SFR in the brightest galaxies at $z>3-4$ constraining their growth in a very short interval."
    },
    "0512/astro-ph0512371_arXiv.txt": {
        "abstract": "We report  spectra obtained with the {\\it Spitzer Space Telescope} in the ${\\lambda}$ = 5 -- 35 ${\\mu}$m range of HD 233517, an evolved K2 III giant with circumstellar dust.  For ${\\lambda}$ $>$ 13 ${\\mu}$m, the flux is a smooth continuum that varies approximately as ${\\nu}^{-5/3}$. For ${\\lambda}$ $<$ 13 ${\\mu}$m, although the star is oxygen-rich,   PAH features produced by carbon-rich species at 6.3 ${\\mu}$m, 8.2 ${\\mu}$m,  11.3  ${\\mu}$m  and 12.7 ${\\mu}$m are detected along with likely broad silicate emission near 20 ${\\mu}$m. These results can be explained if there is  a passive, flared disk orbiting HD 233517.   Our data support  the hypothesis that organic molecules in orbiting disks  may  be synthesized {\\it in situ} as well as being incorporated from the interstellar medium. ",
        "introduction": "We  study  dusty circumstellar disks  in order to understand better the origin and  evolution of planets and stars.  Analogous to young stellar objects, disks around evolved stars such as the Red Rectangle are important in the angular momentum budget of the system and also may be sites where planets form, as occurred around the pulsar PSR 1215+12 (Wolszczan \\& Frail 1992).  Here, we report data obtained with the Infrared Spectrograph (Houck et al. 2004) on the {\\it Spitzer Space Telescope} (Werner et al. 2004) that provide  supporting evidence for the hypothesis (Jura 2003) that the evolved red giant  HD 233517 (K2 III) possesses an orbiting disk.    About 10$^{5}$ red giants  were  detected with IRAS. While Asymptotic Giant Branch  (AGB) stars with luminosities larger than 10$^{3}$ L$_{\\odot}$ usually have large enough mass-loss rates to produce substantial infrared excesses (see, for example, Sopka et al. 1985), the less powerful first-ascent red giants with luminosities near 100 L$_{\\odot}$ typically do not have measurable infrared excesses. Even when present, the typical fractional excess is only between  10$^{-4}$ and 10$^{-3}$ of the star's luminosity (Zuckerman et al. 1995), and at least in some cases actually may be produced by interstellar cirrus rather than mass loss (Kim et al. 2001). Only a handful of  class III red giants with larger  excesses  are known; a well-studied example  is  HD 233517 with a luminosity of 90 L$_{\\odot}$ and a fractional excess of  ${\\sim}$ 0.06 (Sylvester et al. 2001). Identified as anomalous in the IRAS data base (Walker \\& Wolstencroft 1988), initially, it was thought to be a main-sequence  star (Skinner et al. 1995).  Although its parallax is not measured, high-resolution optical spectroscopy shows that it is a  red giant (Fekel et al. 1996, Balachandran et al. 2000, Zuckerman 2001). Furthermore, HD 233517 is likely to be post-main-sequence rather than pre-main-sequence  because it is spatially isolated from any known region of star formation, does not fall on any theoretical tracks in the H-R diagram for pre-main-sequence stars and shows an abundance of lithium of [Li]/[H] = 1.7 ${\\times}$ 10$^{-8}$ about an order of magnitude greater than the standard interstellar value of 2 ${\\times}$ 10$^{-9}$ (Anders \\& Grevesse 1989, Howarth et al. 2002).  This large abundance of lithium can be explained by models of surface mixing of processed material in some models of post-main-sequence evolution (see, for example, Denissenkov \\& Weiss 2000,  Denissenkov \\& Herwig 2004), but not by current models for pre-main-sequence stars.  HD 233517 has  $v\\,\\sin\\,i$ = 17.6 km s$^{-1}$ (Balachandran et al. 2000) which is substantially greater than the typical rotational speed of a class III red giant of ${\\leq}$5 km  s$^{-1}$ (de Medeiros \\& Mayor 1995, Gray \\& Pallavicini 1989, Schrijver \\& Pols 1993), but smaller than the maximum possible of ${\\sim}$35 km s$^{-1}$.  There are a few other first-ascent red giants which also have marked infrared excesses, distinctively high lithium abundances and unusually rapid rotation (see, for example, Drake et al. 2002, Reddy \\& Lambert 2005).   Any model to describe HD 233517 also may  pertain to these  stars.  Previously, Jura (2003) has argued that the  infrared excess around HD 233517 is unlikely to be produced by a recent outflow in a stellar wind.   For most  red giants that are currently losing mass,  F$_{\\nu}$  typically varies as ${\\nu}^{1.5}$ for  ${\\lambda}$ $>$ 10 ${\\mu}$m (see, for example, Sopka et al. 1985).  However, for HD 233517, the IRAS fluxes between  ${\\lambda}$ = 13 ${\\mu}$m and ${\\lambda}$ = 60 ${\\mu}$m vary  approximately as ${\\nu}^{-5/3}$, a spectral energy distribution that can be naturally modeled with  a passive, flared, orbiting disk.  Also, the CO radio emission from the system is probably undetected (Jura \\& Kahane 1999, Dent et al. 2005).  In contrast, the winds from  mass-losing red giants typically are easily detectable CO sources (see, for example, Olofsson et al. 1993). Here, we report additional evidence in support of the view that the infrared excess of HD 233517 is produced by orbiting dust.  Our study of HD 233217 can help us understand other circumstellar disks. For example, there is controversy about the origin of organic molecules in the early solar system.  The usual view is that most of the carbonaceous material was incorporated as carbon-rich compounds from the interstellar medium and then further processed (Ehrenfreund \\& Charnley 2000, Kerridge 1999).  However, a contrary view is that even in the oxygen-rich solar nebula, Fischer-Tropsch catalysis on the surface of metal grains at $T$ $>$ 400 K was an important route for the synthesis of carbon-rich molecules (Kress \\& Tielens 2001, Zolotov \\& Shock 2001).  Alternatively, gas-phase synthesis of such molecules may be important (Morgan et al. 1991).  Since the disk around HD 233517 is free of interstellar contamination, it can serve as an indirect test of these models. ",
        "conclusions": "The infrared spectrum of  HD 233517 shows emission from PAHs and has a continuum that can be  reproduced approximately with a passive, flared, orbiting disk.  The data can be explained if the material resides in a long-lived orbiting disk which may have been created when  HD 233517  engulfed a companion. The data lend support to the hypothesis that organic molecules in disks may  be synthesized {\\it in situ} as well as being incorporated from the interstellar medium."
    },
    "0512/astro-ph0512001_arXiv.txt": {
        "abstract": "We present a color-magnitude diagram (CMD) for a field in the giant tidal stream of the Andromeda galaxy (M31).  These observations, taken with the Advanced Camera for Surveys on the {\\it Hubble Space Telescope}, are 50\\% complete at $V\\approx 30$, reaching 1~mag below the oldest main-sequence turnoff.  Striking similarities between the stream and a previous spheroid CMD imply they have very similar age and metallicity distributions, but present something of an enigma; we speculate on possible interpretations of this result, but note that none are without problems. Distinct multiple turnoffs, as might be expected from pulses of star formation caused by interaction with Andromeda,  are not apparent in the stream CMD. Star formation in both fields lasted about 6 billion years, building up to relatively high metallicities and being largely complete 6 billion years ago. The close similarity of the spheroid and stream suggests that both may have derived from the same event; it would be worth exploring to what extent stars in these structures are the remnants of a disk galaxy that interacted with M31, or even were disrupted from the M31 disk itself by the interaction. ",
        "introduction": "According to hierarchical models of galaxy formation, spheroids form in a series of mergers between galaxies and proto-galaxies.  These models generally predict more dwarfs than observed, leading to suggestions that most of the proto-dwarfs in the early universe have since dissolved into the spheroids of giant galaxies (e.g., Bullock et al.\\ 2000).  Until relatively recently, traces of such activity were not obvious in the two giant galaxies of the Local Group.  However, the discovery of the Sgr dwarf (Ibata et al.\\ 1994), cannibalized by the Milky Way, sparked renewed interest in the formation of spheroids via the accretion of dwarfs.  Subsequently, a spectacular tidal stream was discovered in Andromeda (Ibata et al.\\ 2001), along with a variety of substructures in the spheroid and outer disk (Ferguson et al.\\ 2002), suggesting an active merger history.  Further insight into the formation of galaxies can be found by studying their star formation histories.  Photometry extending below the oldest main-sequence turnoff in a population can yield its complete formation history, but until the advent of the Advanced Camera for Surveys (ACS; Ford et al.\\ 1998) on the {\\it Hubble Space Telescope (HST)}, it was not feasible to obtain such data for populations much beyond our own Milky Way and its satellites.  However, Brown et al.\\ (2003) used the ACS to obtain extremely deep photometry of a minor-axis field in the \\linebreak  {\\small \\noindent $^1$Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at STScI, and associated with proposal 10265.\\\\ $^2$STScI, 3700 San Martin Dr., Baltimore, MD 21218;  tbrown@stsci.edu, edsmith@stsci.edu, ferguson@stsci.edu.\\\\ $^3$UCO / Lick Observatory, 271 ISB, 1156 High Street, Santa Cruz, CA 95064; raja@ucolick.org.\\\\ $^4$Div. of Astronomy, Dpt.\\ of Physics \\& Astronomy, UCLA, Los Angeles, CA 90095; rmr@astro.ucla.edu.\\\\ $^5$Osservatorio Astronomico, Vicolo Dell'Osservatorio 5, I-35122 Padova, Italy; arenzini@pd.astro.it.\\\\ $^6$Code 667, NASA/GSFC, Greenbelt, MD 20771; \\\\ allen.v.sweigart@nasa.gov, randy.a.kimble@nasa.gov.}  \\noindent Andromeda (NGC~224; M31) spheroid; the resulting CMD shows a dominant intermediate-age population of 6--11~Gyr along with a significant, old, metal-poor population.  The age distribution, high metallicity (Mould \\& Kristian 1986; Durrell, Harris, \\& Pritchet 2001), and substructure in the Andromeda spheroid all point to a violent merger history. Given these disturbances, our use of the term ``spheroid'' in this Letter does not imply a smooth and relaxed structure, but simply refers to the extraplanar stars of M31.  Andromeda's tidal stream remains the most prominent merger remnant in the Local Group, and several years of intense study have yielded important constraints on its origin.  Ibata et al.\\ (2004) found that the stream is kinematically cold, with a velocity dispersion ($\\sigma_v$) of only 11 km s$^{-1}$, and that the stream is increasingly blue-shifted as it approaches Andromeda, implying that the stream is falling into Andromeda from behind the galaxy. More recently, Kalirai et al.\\ (2005) found two distinct kinematic stream components -- both cold ($\\sigma_v \\approx 15$ km s$^{-1}$) and blue-shifted with respect to the spheroid (Figure 1$a$).  In all of these studies, the stream's age distribution has remained largely unconstrained.  Ferguson et al.\\ (2005) published a stream CMD reaching a few magnitudes below the horizontal branch (HB); they found no evidence for very young populations, but speculated that the presence of a strong RGB bump might indicate a narrow age dispersion.  In late 2004, we obtained deeper stream photometry, extending 1~mag below the oldest main-sequence turnoff, thus revealing the entire star formation history.  In this Letter, we report on our initial analysis, which shows a remarkable similarity between the stream population and the spheroid population of Brown et al.\\ (2003), in both the age and metallicity distributions. ",
        "conclusions": "The stream CMD presents a challenge for the hypothesis that the progenitor is a disrupted dwarf galaxy on an early passage close to M31.  The dwarf apparently stopped forming stars at least 6 Gyr ago ($\\lesssim$5\\% of the stars in the stream field are less than 4~Gyr old). Orbit models (e.g., Font et al.\\ 2005) imply that the progenitor of the stream only recently ($\\lesssim 1$~Gyr ago) approached within 100 kpc of M31.  Those models do not explain why star formation stopped long before the interaction started.  The relatively high metallicity of the stream implies that the progenitor had a stellar mass of $\\sim 10^9$ $M_\\odot$ if it followed the scaling relation for Local Group dwarfs found by Dekel \\& Woo (2003). However, according to their correlations, a dwarf galaxy with $\\sigma_v \\approx 15$~km s$^{-1}$ (observed for the stream) would have a stellar mass of only $\\sim 10^6$ $M_\\odot$.  Font et al.\\ (2005) claim that $\\sigma_v$ can vary significantly over the orbit, and could be much lower than that in the progenitor; it remains to be seen if this explanation can account for the large mass discrepancy.  We note in this context that the only known systems that are kinematically cold and metal rich are disks. Thus a disk galaxy may be a more attractive candidate for the progenitor of the stream than a pressure-supported dwarf.  Although the stream and spheroid populations have very different kinematic and spatial distributions, their CMDs are strikingly similar, implying nearly identical age and metallicity distributions. A subtle difference between the CMDs is a slightly brighter ($\\approx$~0.04~mag) SGB in the stream, suggesting a slightly younger mean age ($\\sim$300 Myr).  Could the similarity of the CMDs be explained by the stream passing through the spheroid field, as suggested by some orbital models (e.g., Ibata et al.\\ 2004)?  The stream {\\it dominates} by a 3:1 ratio in our stream CMD (Figure 1$a$; Kalirai et al.\\ 2005).  Thus the stream would have to similarly dominate in the spheroid field, yet show no clear morphological signature in the Ferguson et al.\\ (2002) maps, and exhibit the kinematics shown in Figure 1$b$ (a broad distribution at the M31 systemic velocity; Rich et al.\\ 2005, in prep.); this seems implausible given the current orbital models.  Thus, if there is ``contamination'' from the stream in the spheroid field, it is likely to be much more complex and extended over a much wider region than that suggested by models where the infalling object has only completed one or two orbits.  Perhaps the stars in the spheroid came from a different dwarf galaxy (or galaxies) that merged with M31 earlier. This hypothesis can help explain why other locations in the spheroid (Ferguson et al.\\ 2005) show RGB and HB morphologies similar to our spheroid field (Brown et al.\\ 2003).  However, it then becomes mere coincidence that the star-formation histories in the stream and the spheroid look nearly identical.  Perhaps the spheroid is a disrupted disk -- either the M31 disk itself (Brown et al.\\ 2003), or the remnants of a disk galaxy that merged with M31.  The star formation history in the spheroid might be naturally explained if the disruption occurred 6--8 Gyr ago.  We are still left invoking coincidence to explain the similar CMDs -- unless the stream is also part of the disrupted disk.  It would be interesting to explore a wider range of dynamical simulations than those considered previously, to see whether the stream might {\\it not} be a dwarf galaxy that has encountered M31 in the last 1 Gyr, but instead a remnant of a merged disk galaxy or a plume of stars disrupted from the M31 disk.  It is unclear if this interpretation can be reconciled with the stream's structure and kinematics.  Compared to a dwarf, a disk population might better explain the extended star formation history and high metallicity present in the stream and spheroid.  Measurements of the formation history at multiple locations in the stream and distorted spheroid of M31 are undoubtedly going to be critical for sorting out what happened. Guhathakurta et al.\\ (2005b) and Irwin et al.\\ (2005) find that the minor-axis surface-brightness profile changes from a de Vaucouleurs $r^{1/4}$ law to an $r^{-2.3}$ power law beyond $\\sim$30 kpc; exploration of this region might yield important constraints on the ``primordial'' halo unaffected by Andromeda's violent merger history."
    },
    "0512/astro-ph0512237_arXiv.txt": {
        "abstract": "{Protoplanetary nebulae typically present non-spherical envelopes. The origin of such geometry is still controversial. There are indications that it may be carried over from an earlier phase of stellar evolution, such as the AGB phase. But how early in the star's evolution does the non-spherical envelope appear? } {Li-rich giants show dusty circumstellar envelopes that can help answer that question. We study a sample of fourteen Li-rich giants using optical polarimetry in order to detect non-spherical envelopes around them. } {We used the IAGPOL imaging polarimeter to obtain optical linear polarization measurements in ${\\it V}$ band. Foreground polarization was estimated using the field stars in each CCD frame.} {After foreground polarization was removed, seven objects presented low intrinsic polarization (0.19 $-$ 0.34)\\% and two (\\object{V859 Aql} and \\object{GCSS 557}) showed high intrinsic polarization values (0.87 $-$ 1.16)\\%. This intrinsic polarization suggests that Li-rich giants present a non-spherical distribution of circumstellar dust. The intrinsic polarization level is probably related to the viewing angle of the envelope, with higher levels indicating objects viewed closer to edge-on. The correlation of the observed polarization with optical color excess gives additional support to the circumstellar origin of the intrinsic polarization in Li-rich giants. The intrinsic polarization correlates even better with the IRAS  25 $ \\mu $m far infrared emission. Analysis of spectral energy distributions for the sample show dust temperatures for the envelopes tend to be between 190 and 260 K. We suggest that dust scattering is indeed responsible for the optical intrinsic polarization in Li-rich giants.} {Our findings indicate that non-spherical envelopes may appear as early as the red giant phase of stellar evolution. } ",
        "introduction": "Mass loss plays a central role in the late stages of stellar evolution. In particular, asymptotic giant branch (AGB) stars may be responsible for up to 60\\% of the (all stars) interstellar dust input into the interstellar medium (Gehrz \\cite{gehr89}).  Our current understanding of stellar evolution places the AGB stars as precursors of protoplanetary nebulae (PPN). On the other hand, we know from direct imaging that PPN typically present non-spherically symmetric envelopes. While extrinsic (to the PPN) effects such as binarity have been said to cause the non-sphericities (for a review see Balick \\& Frank \\cite{bali02}), could it not be instead that the asymmetry is carried over from the earlier, AGB phase of stellar evolution?  One piece of evidence showing that departures from spherical symmetry exist around AGB stars is that light from these objects may show some degree of linear polarization (Coyne \\& Magalh\\~aes \\cite{coyn77,coyn79}; Magalh\\~aes et al. \\cite{maga86a,maga86b}; Kahane et al. \\cite{kaha97}; Magalh\\~aes \\& Nordsieck \\cite{maga00}). Further evidence of asphericity in the envelopes of late-type giant and supergiant star envelopes comes from the details of OH maser emission profiles (Collison \\& Fix \\cite{coll92}), KI resonant scattering (Plez \\& Lambert \\cite{plez94}), rings of SiO maser emission (Diamond et al. \\cite{diam94}, Greenhil et al. \\cite{gree95}), and OH radio images (Chapman et al. \\cite{chap94}).  Aspherical symmetries, such as those found by Trammell et al. (1994) in post-AGB stars from spectropolarimetry, may then be understood naturally, since such symmetries are already present in earlier evolutionary stages. The observed and more obvious non-spherical symmetries in PPN  (Balick \\& Franck \\cite{bali02}) and the inferred assymetries in other evolved objects obtained from polarimetry (Johnson \\& Jones \\cite{john91}, Parthasarathy \\& Jain \\cite{part93}, Parthasarathy et al. \\cite{part05}) are also consistent with the origin of the asymmetries early in the AGB phase.  Aside from the fact that polarimetry {\\it per se} allows the study of otherwise unresolved objects, the question arises as to how early in stellar evolution the non-spherical symmetry appears. In this paper we use the fact that Li-rich red giants (RG) show dusty circumstellar envelopes to explore their environments and to help answer that question.  Over the past two decades around 40 red giants were found to have Li abundances that are 100 times larger than the mean values observed in red giants (Brown et al. \\cite{brow89}, de la Reza \\& Drake \\cite{dela95}). This number indicates that about 2$\\%$ of the Population I red giants show significantly larger lithium abundances than expected by dilution due to mixing by classical convection. In some of the giants, the lithium abundance reaches values that are similar to (and even larger than) the Pop I value (meteoritic, open clusters, etc.), around logN(Li) = 3.3. Castilho et al. (\\cite{cast99a}) obtained very low Be abundances for two Li-rich giants (LRG) providing evidence that the original Li in these stars must have been almost completely destroyed and that the high Li abundances in the Li-rich red giants are due to Li production in these stars.  It seems to be clear that LRG are quite normal stars, except for their high Li abundance and large infrared excess (Castilho et al. \\cite{cast95}, de la Reza et al. \\cite{reza96}, Castilho et al. \\cite{cast00}). The similarities to normal giants, namely mass, chemical composition, temperature, and metallicity, combined with the far-infrared emission, indicate that LRG do not form a unique class of objects but are ordinary low-mass stars observed during a short phase of their evolution, when Li is created.  If indeed all low mass red giants go through a phase of Li production in the RGB (that could be cyclic), together with an increase in mass loss seen in the IRAS color diagram (Gregorio-Hetem et al. \\cite{greg93}, de la Reza et al. \\cite{reza96}), and since some of them reach a larger lithium abundance than the Pop I abundance, they could be an important source of Li enrichment in the Galaxy. Some information needed to quantify the LRG contribution for the interstellar medium Li enrichment remains unknown, such as: the maximum Li abundance reached for each star and its relation with stellar mass and/or metallicity, the duration of the Li production phase(s), the mass loss process, and rate, and  the simultaneity of the envelope ejection with the Li production.  Up to now only one envelope of LRG has been studied in detail. The $\\sim$ 3x4 arcmin envelope of \\object{HD 65750} (Castilho et al. \\cite{cast98}) has a butterfly geometry and the present mass loss rate of the star is not enough to form the observed envelope. Witt \\& Rogers (\\cite{witt91}) proposed that a past and more efficient mass-loss event about 32,000 years ago was responsible for the observed structure.  The aim of this work is to measure the polarization of LRG stars so as to study the spatial distribution of the circumstellar dust around these objects. Scattering of the stellar radiation by the circumstellar dust can produce polarization that is measurable in objects with non-resolved envelopes. The amount of polarization will depend upon the density and nature of the scatterers, as well as on the aspect angle and exact geometry of the envelope (Magalh\\~aes \\cite{maga92}).  If intrinsic polarization is detected, we can infer that an asymmetric spatial distribution of circumstellar dust is present in these objects.  Here we present the results of optical linear polarization measurements of fourteen LRG and two normal giants (NG). In Sect. 2 we describe the observations and data reduction along with the calculations of foreground and intrinsic polarizations. In Sect. 3 we show the stellar parameters of our sample and a general discussion is presented including correlations between the polarimetric data and optical and near infrared excess color. Correlation with IRAS colors also are explored and spectral energy distributions used to estimate the dust temperature associated to the circumstellar envelopes. The conclusions are drawn in Sect. 4.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{4270f01.eps}} \\caption{Histograms of (a) the observed and (b) intrinsic polarization of LRG. The NG in the sample are indicated in black.} \\label{histo} \\end{figure}   \\begin{table*} \\caption{Log of observations.\\label{log}} \\begin{tabular}{lccccccccc} \\hline \\hline Object     & Log $ \\varepsilon $(Li)$^{a}$     & {\\it l} & {\\it b} &  {\\it V}$^{b}$  &  Wav. pos. &  IT &  Aper. & Date & ID \\\\ & (dex) & &   & (mag) &  & (sec.) & (\\arcsec)  & \\\\ (1) & (2)  & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) \\\\ \\hline \\multicolumn{10}{c}{Li-rich giants}\\\\ \\hline \\object{HD 90082}\t&\t0.20\t(0.15)\t&\t286.27\t&\t-3.91\t&\t7.50\t&\t16\t&\t8\t&\t5.5\t&\t06/20/98\t&\ta\t\\\\ \\object{HD 95799}\t&\t3.10\t(0.15)\t&\t289.19\t&\t1.17\t&\t8.14\t&\t16\t&\t12\t&\t5.5\t&\t06/20/98\t&\tb\t\\\\ \\object{HD 96195}\t&\t0.40\t(0.20)\t&\t291.12\t&\t-2.57\t&\t7.94\t&\t16\t&\t13\t&\t6.7\t&\t06/20/98\t&\tc\t\\\\ \\object{HD 120602}\t&\t1.90\t(0.15)\t&\t338.54\t&\t64.21\t&\t6.01\t&\t16\t&\t1\t&\t3.7\t&\t06/20/98\t&\td\t\\\\ \\object{PDS 68}\t&\t2.00\t(0.40)\t&\t315.72\t&\t19.10\t&\t12.80\t&\t4\t&\t300\t&\t4.3\t&\t06/20/98\t&\te\t\\\\ \\object{HD 146850}\t&\t1.60\t(0.15)\t&\t359.59\t&\t24.44\t&\t6.10\t&\t6$^{c}$\t&\t1\t&\t3.7\t&\t06/20/98\t&\tf\t\\\\ \\object{GCSS 557} (\\object{V385 Sct})\t&\t0.70\t(0.15)\t&\t17.10\t&\t-1.37\t&\t13.26\t&\t4\t&\t300\t&\t3.7\t&\t06/20/98\t&\tg\t\\\\ \\object{HD 176588}\t&\t1.60\t(0.15)\t&\t30.08\t&\t-4.18\t&\t6.89\t&\t16\t&\t2\t&\t4.9\t&\t06/20/98\t&\th\t\\\\ \\object{IRAS 19012-0747}\t&\t2.60\t(0.15)\t&\t27.46\t&\t-6.28\t&\t11.17\t&\t16\t&\t60\t&\t3.1\t&\t06/20/98\t&\ti\t\\\\ \\object{IRAS 19038-0026}\t&\t0.60\t(0.25)\t&\t34.32\t&\t-3.49\t&\t14.48\t&\t8\t&\t300\t&\t3.7\t&\t06/20/98\t&\tj\t\\\\ \\object{HD 178168}\t&\t0.90\t(0.15)\t&\t32.54\t&\t-4.72\t&\t9.06\t&\t16\t&\t20\t&\t5.5\t&\t06/20/98\t&\tk\t\\\\ \\object{HD 112127}\t&\t2.80\t(0.15)\t&\t0.96\t&\t89.34\t&\t6.87\t&\t16\t&\t1\t&\t4.9\t&\t06/21/98\t&\tl\t\\\\ \\object{V859 Aql} (\\object{PDS 100})\t&\t2.50\t(0.15)\t&\t42.29\t&\t-6.28\t&\t10.44\t&\t16\t&\t40\t&\t4.9\t&\t06/21/98\t&\tm\t\\\\ \\object{HD 203251}\t&\t1.40\t(0.15)\t&\t35.57\t&\t-39.97\t&\t8.00\t&\t16\t&\t8\t&\t7.3\t&\t06/21/98\t&\tn\t\\\\ \\hline \\multicolumn{10}{c}{normal giants}\\\\ \\hline \\object{HD 190664}\t&\t$<$0.00 (0.30)\t&\t37.92\t&\t-18.48\t&\t6.47\t&\t16\t&\t1\t&\t3.7\t&\t06/21/98\t&\to\t\\\\ \\object{HD 124649}\t&\t$<$0.00 (0.30)\t&\t315.53\t&\t7.45\t&\t7.86\t&\t16\t&\t6\t&\t6.1\t&\t06/21/98\t&\tp\t\\\\ \\hline \\end{tabular} \\\\ Notes: (a) log $ \\varepsilon $(Li) were obtained from Castilho et al. \\cite{cast00}, except for: \\object{HD 120602} and \\object{HD 112127} (from Brown et al. \\cite{brow89}), \\object{HD 146850} (from Castilho et al. \\cite{cast95}), \\object{V959 Aql} (from Reddy et al. \\cite{redd02}), and \\object{HD203251} (from Fekel \\& Balachandran \\cite{feke93}). The value for \\object{PDS 68} is an estimation based on the equivalent width of the Li I (6707.8\\AA) line (from Gregorio-Hetem et al. \\cite{greg92}), and the values for \\object{HD 190664} and \\object{HD 124649} are higher limits (Castilho et al. \\cite{cast99a}). The errors are indicated in parenthesis; (b) {\\it V} were obtained from Castilho (\\cite{cast99b}), except for: \\object{HD 120602}, \\object{HD 112127}, \\object{HD 203251}, and \\object{HD 190664} (from GCPC, Mermilliod et al. \\cite{merm97}), \\object{PDS 68} and \\object{V859 Aql} (from Gregorio-Hetem et al. \\cite{greg92}), \\object{HD 124649} from CDS, and \\object{IRAS 19038-0026} (from the GSC 2.2 catalogue, using magnitude in the visual phographic band as {\\it V}); (c) 8 waveplate positions were initially observed but two of them were saturated. \\label{log} \\end{table*}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{4270f02.eps}} \\caption{Correlation between observed polarization and galactic latitude. The black dots are the LRG, and the white dots are the NG. The letters indicate the ID for each object as it appears in column 10 of Table \\ref{log}.} \\label{lat} \\end{figure} ",
        "conclusions": "Optical polarimetry for a sample of fourteen LRG stars was obtained. For nine of them intrinsic polarization was estimated using field stars. Seven LRG (five of them with {\\it P}/$\\sigma$$_{{\\it P}}$ $>$ 5) have lower but non-negligible intrinsic polarization (0.19 $-$ 0.34)\\%, and in two cases (\\object{V859 Aql} and \\object{GCSS 557}) intrinsic polarizations higher than 0.5\\% are found. These results indicate that an asymmetric spatial distribution of circumstellar dust is present in LRG (but probably also in normal giants). An excess in observed polarization when it is correlated with the optical excess color gives additional support to the circumstellar origin of the intrinsic polarization in LRG. The optical intrinsic polarization in LRG is not correlated with the near infrared excess but is correlated with far infrared emission. This would suggest that grains emitting in $\\sim 25$ $\\mu$m are responsible for the optical intrinsic polarization, and the higher intrinsic polarization levels would indicate a favorable (edge-on) viewing angle for the envelopes, as in \\object{V859 Aql} and \\object{GCSS 557}. Analysis of spectral energy distributions for the sample provides an estimate of the dust temperature for the envelopes, which are mainly between 190 K and 260 K. Our findings indicate that non-spherical symmetries may appear as early as the RG phases of stellar evolution."
    },
    "0512/astro-ph0512598_arXiv.txt": {
        "abstract": "We discuss certain aspects of the production of $\\rm ^{60}Fe$ in massive stars in the range between 11 and 120 $\\rm M_\\odot$, both in the hydrostatic and explosive stages. We also compare the $\\rm ^{60}Fe/^{26}Al$ $\\gamma$-ray line flux ratio obtained in the present calculations to the detected value reported by INTEGRAL/SPI. ",
        "introduction": "\\label{intro} The nucleus $\\rm ^{60}Fe$ is a long lived ($\\tau\\sim$ 2 Myr) radioactive isotope that should be present in an appreciable amount in our galaxy. Historically, $\\rm ^{60}Fe$ has been considered as a key isotope to understand whether or not massive stars are the main contributors to the diffuse $\\rm ^{26}Al$ present in the Galaxy, another long lived radioactive isotope ($\\tau\\sim$ 1 Myr) traced by the detected 1.809 MeV $\\gamma$-ray emission line at a level of $\\rm \\sim 4\\times 10^{-4}$ $\\rm cm^{-2}$ $\\rm s^{-1}$. Indeed, already in the early 80's \\cite{clayton82} pointed out that SNII are the only candidate sources for $\\rm ^{26}Al$ to produce also $\\rm ^{60}Fe$ and hence that the detection of this isotope in the Milky Way could constitute a strong argument in favour of SNII as the main contributors to the galactic $\\rm ^{26}Al$. The first detection of $\\rm ^{60}Fe$ in the Galaxy was obtained with RHESSI and reported by \\cite{smith03}. The line flux detected implies a $\\rm ^{60}Fe/^{26}Al$ $\\gamma$-ray line flux ratio of $\\sim 0.16$ (for each $\\rm ^{60}Fe$ line). More recently \\cite{harrisetal05} reported the first detection of $\\rm ^{60}Fe$ decay lines at 1.173 MeV and 1.333 MeV in spectra taken by the SPI spectrometer on board INTEGRAL during its first year, yielding a $\\gamma$-ray line flux of $\\rm 3.7\\pm 1.1~\\times 10^{-5}~\\gamma~cm^{-2}~s^{-1}$. The same analysis applied to the 1.809 MeV line of $\\rm ^{26}Al$ yielded a $\\rm ^{60}Fe/^{26}Al$ $\\gamma$-ray line flux ratio of $\\sim 0.11\\pm0.03$. From a theoretical side, many groups have performed calculations of nucleosynthesis in massive stars, estimating the amounts of either $\\rm ^{26}Al$ or both $\\rm ^{26}Al$ and $\\rm ^{60}Fe$ ejected in the interstellar medium. However, no set of models covers an extended grid of stellar masses. Indeed, several research groups interested in the presupernova evolution and explosion of massive stars provided yields of both $\\rm ^{26}Al$ and $\\rm ^{60}Fe$ for stars up to a mass of $\\rm 40~M_\\odot$ \\citep{CL04,RHHW02,WW95,TNH96} - among these works only \\cite{RHHW02} included mass loss in the computations. On the other hand, groups mainly interested in the evolution of massive stars including mass loss computed the evolution up to the end of central He burning and hence provide only the hydrostatic yield of $\\rm ^{26}Al$ \\citep{MAPP97,PMVKSCM05,LBF95}.  In this paper we will discuss to some extent the production of $\\rm ^{60}Fe$ in massive stars in the range between 11 and 120 $\\rm M_\\odot$, both in the hydrostatic and explosive stages. We will also provide theoretical predictions for the $\\rm ^{60}Fe/^{26}Al$ $\\gamma$-ray line flux ratio of such a generation of massive stars and we will compare them with the observations. ",
        "conclusions": ""
    },
    "0512/astro-ph0512551_arXiv.txt": {
        "abstract": "We present an analysis of the HD 82943 planetary system based on a radial velocity data set that combines new measurements obtained with the Keck telescope and the CORALIE measurements published in graphical form. We examine simultaneously the goodness of fit and the dynamical properties of the best-fit double-Keplerian model as a function of the poorly constrained eccentricity and argument of periapse of the outer planet's orbit. The fit with the minimum $\\chi_{\\nu}^2$ is dynamically unstable if the orbits are assumed to be coplanar. However, the minimum is relatively shallow, and there is a wide range of fits outside the minimum with reasonable $\\chi_{\\nu}^2$. For an assumed coplanar inclination $i = 30^{\\circ}$ ($\\sin i = 0.5$), only good fits with both of the lowest order, eccentricity-type mean-motion resonance variables at the 2:1 commensurability, $\\theta_1$ and $\\theta_2$, librating about $0^{\\circ}$ are stable. For $\\sin i = 1$, there are also some good fits with only $\\theta_1$ (involving the inner planet's periapse longitude) librating that are stable for at least $10^8$ years. The libration semiamplitudes are about $6^{\\circ}$ for $\\theta_1$ and $10^{\\circ}$ for $\\theta_2$ for the stable good fit with the smallest libration amplitudes of both $\\theta_1$ and $\\theta_2$. We do not find any good fits that are non-resonant and stable. Thus the two planets in the HD 82943 system are almost certainly in 2:1 mean-motion resonance, with at least $\\theta_1$ librating, and the observations may even be consistent with small-amplitude librations of both $\\theta_1$ and $\\theta_2$. ",
        "introduction": "The first pair of extrasolar planets suspected to be in mean-motion resonance was discovered around the star GJ 876, with the orbital periods nearly in the ratio 2:1 \\citep{mar01}. A dynamical fit to the radial velocity data of GJ 876 that accounts for the mutual gravitational interaction of the planets is essential because of the short orbital periods ($\\approx 30$ and 60 days) and large planetary masses [minimum combined planetary mass relative to the stellar mass $(m_1 + m_2)/m_0 \\approx 0.0074$]. It is now well established that this pair of planets is deep in 2:1 orbital resonance, with both of the lowest order, eccentricity-type mean-motion resonance variables, \\begin{equation} \\theta_1 = \\lambda_1 - 2\\lambda_2 + \\varpi_1 \\label{theta1} \\end{equation} and \\begin{equation} \\theta_2 = \\lambda_1 - 2\\lambda_2 + \\varpi_2 , \\label{theta2} \\end{equation} librating about $0^\\circ$ with small amplitudes \\citep{lau01,riv01,lau05}. Here $\\lambda_{1,2}$ are the mean longitudes of the inner and outer planets, respectively, and $\\varpi_{1,2}$ are the longitudes of periapse. The simultaneous librations of $\\theta_1$ and $\\theta_2$ about $0^\\circ$ mean that the secular apsidal resonance variable, \\begin{equation} \\theta_{\\rm SAR} = \\varpi_1 - \\varpi_2 = \\theta_1 - \\theta_2 , \\label{thetaSAR} \\end{equation} also librates about $0^\\circ$ and that the periapses are on average aligned. The basic dynamical properties of this resonant pair of planets are not affected by the recent discovery of a third, low-mass planet on a 1.9-day orbit in the GJ 876 system \\citep{riv05}.  It is important to confirm other suspected resonant planetary systems, as the GJ 876 system has shown that resonant planetary systems are interesting in terms of both their dynamics and their constraints on processes during planet formation. The geometry of the 2:1 resonances in the GJ 876 system is different from that of the 2:1 resonances between the Jovian satellites Io and Europa, where $\\theta_1$ librates about $0^\\circ$ but $\\theta_2$ and $\\theta_{\\rm SAR}$ librate about $180^\\circ$. For small orbital eccentricities, the Io-Europa configuration is the only stable 2:1 resonance configuration with both $\\theta_1$ and $\\theta_2$ librating. For moderate to large eccentricities, the Io-Europa configuration is not stable, but there is a wide variety of other stable 2:1 resonance configurations, including the GJ 876 configuration \\citep{lee02,lee03a,bfm03,fbm03,lee04}. The resonances in the GJ 876 system were most likely established by converging differential migration of the planets leading to capture into resonances, with the migration due to interaction with the protoplanetary disk. While it is easy to establish the observed resonance geometry of the GJ 876 system by convergent migration, the observational upper limits on the eccentricities ($e_1 \\la 0.31$ and $e_2 \\la 0.05$) require either significant eccentricity damping from planet-disk interaction or resonance capture occurring just before disk dispersal, because continued migration after resonance capture could lead to rapid growth of the eccentricities \\citep{lee02,kle05}. Hydrodynamic simulations of the assembly of the GJ 876 resonances performed to date do not show significant eccentricity damping from planet-disk interaction and produce eccentricities that exceed the observational upper limits, unless the disk is dispersed shortly after resonance capture \\citep{pap03,kle04,kle05}.  HD 82943 was the second star discovered to host a pair of planets with orbital periods nearly in the ratio 2:1. The discovery of the first planet was announced in ESO Press Release 13/00\\footnote{ See \\anchor{http://www.eso.org/outreach/press-rel/pr-2000/pr-13-00.html} {http://www.eso.org/outreach/press-rel/pr-2000/pr-13-00.html} } and the discovery of the second, inner planet was announced in ESO Press Release 07/01\\footnote{ See \\anchor{http://www.eso.org/outreach/press-rel/pr-2001/pr-07-01.html} {http://www.eso.org/outreach/press-rel/pr-2001/pr-07-01.html} }. \\citet{may04} have recently published the radial velocity measurements obtained with the CORALIE spectrograph on the 1.2-m Euler Swiss telescope at the ESO La Silla Observatory in graphical form only. Unlike GJ 876, a double-Keplerian fit is likely adequate for HD 82943, because the orbital periods are much longer ($\\approx 220$ and 440 days) and the planetary masses are smaller [minimum $(m_1 + m_2)/m_0 \\approx 0.003$], and the mutual gravitational interaction of the planets is not expected to affect the radial velocity significantly over the few-year time span of the available observations. \\citet{may04} have found a best-fit double-Keplerian solution, and its orbital parameters are reproduced in Table \\ref{tab:mayor}.  \\citet{fmb05} have reported simulations of the \\citet{may04} best-fit solution that are unstable, but they assumed that the orbital parameters are in astrocentric coordinates and they did not state the ranges of orbital inclinations and starting epochs examined. The fact that the orbital parameters obtained from multiple-Keplerian fits should be interpreted as in Jacobi coordinates was first pointed out by \\citet{lis01} and was derived and demonstrated by \\citet{lee03b}. We have performed direct numerical orbit integrations of the \\citet{may04} best-fit solution, assuming that the orbital parameters are in Jacobi coordinates and that the orbits are coplanar with the same inclination $i$ from the plane of the sky. When we assumed that the orbital parameters correspond to the osculating parameters at the epoch JD 2451185.1 of the first CORALIE measurement, the system becomes unstable after a time ranging from $\\sim 500$ to $4 \\times 10^4\\yr$ for $\\sin i = 1$, $0.9$, $\\ldots$, $0.5$. Since the long-term evolution can be sensitive to the epoch that the orbital parameters are assumed to correspond to, we have repeated the direct integrations by starting at three other epochs equally spaced between the first (JD 2451185.1) and last (JD 2452777.7) CORALIE measurements. All of the integrations become unstable, with the vast majority in less than $10^6\\yr$. These results and those in \\citet[][who also examined mutually inclined orbits and the effects of the uncertainty in the stellar mass]{fmb05} show that the best-fit solution found by \\citet{may04} is unstable.  It is, however, essential that one examines not only the best-fit solution that minimizes the reduced chi-square statistic $\\chi_\\nu^2$, but also fits with $\\chi_\\nu^2$ not significantly above the minimum, especially if the $\\chi_\\nu^2$ minimum is shallow and $\\chi_\\nu^2$ changes slowly with variations in one or more of the parameters. \\citet{fmb05} have analyzed the CORALIE data published in \\citet{may04} by generating a large number of orbital parameter sets and selecting those that fit the radial velocity data with the RMS of the residuals (instead of $\\chi_\\nu^2$) close to the minimum value. The stable coplanar fits that they have found have arguments of periapse $\\omega_1 \\approx 120^\\circ$ and $\\omega_2 \\ga 200^\\circ$, and both $\\theta_1$ and $\\theta_2$ librate about $0^\\circ$ with large amplitudes.\\footnote{ For coplanar orbits, the longitudes of periapse $\\varpi_j$ in the resonance variables $\\theta_1$, $\\theta_2$, and $\\theta_{\\rm SAR}$ (eqs. [\\ref{theta1}]--[\\ref{thetaSAR}]) can be measured from any reference direction in the orbital plane. If we choose the ascending node referenced to the plane of the sky as the reference direction, $\\varpi_j$ are the same as the arguments of periapse $\\omega_j$ obtained from the radial velocity fit. } On the other hand, \\citet{may04} have noted that it is possible to find an aligned configuration with $\\omega_1 \\approx \\omega_2$ that fits the data with nearly the same RMS as their best-fit solution, although it is unclear whether such a fit would be stable and/or in 2:1 orbital resonance. Thus, it has not been firmly established that the pair of planets around HD 82943 are in 2:1 orbital resonance, even though it is likely that a pair of planets of order Jupiter mass so close to the 2:1 mean-motion commensurability would be dynamically unstable unless they are in 2:1 orbital resonance.  In this paper we present an analysis of the HD 82943 planetary system based on a radial velocity data set that combines new measurements obtained with the Keck telescope and the CORALIE measurements published in graphical form. The stellar characteristics and the radial velocity measurements are described in \\S 2. In \\S 3 we present the best-fit double-Keplerian models on a grid of the poorly constrained eccentricity, $e_2$, and argument of periapse, $\\omega_2$, of the outer planet's orbit. In \\S 4 we use dynamical stability to narrow the range of reasonable fits and examine the dynamical properties of the stable fits. We show that the two planets in the HD 82943 system are almost certainly in 2:1 orbital resonance, with at least $\\theta_1$ librating, and may even be consistent with small-amplitude librations of both $\\theta_1$ and $\\theta_2$. Our conclusions are summarized and discussed in \\S 5. ",
        "conclusions": "We have analyzed the HD 82943 planetary system by examining the best-fit double-Keplerian model to the radial velocity data as a function of the poorly constrained eccentricity and argument of periapse of the outer planet's orbit. We have not found any good fits that are non-resonant and dynamically stable (if the orbits are assumed to be coplanar), and the two planets in the HD 82943 system are almost certainly in 2:1 mean-motion resonance. The fit I with the minimum $\\chi_{\\nu}^2$ is unstable, but there is a wide range of fits outside the minimum with $\\chi_{\\nu}^2$ only slightly higher than the minimum. If the unknown $\\sin i \\approx 1$, there are stable good fits with both of the mean-motion resonance variables, $\\theta_1 = \\lambda_1 - 2\\lambda_2 + \\varpi_1$ and $\\theta_2 = \\lambda_1 - 2\\lambda_2 + \\varpi_2$, librating about $0^\\circ$ (e.g., fits II and III), as well as stable good fits with only $\\theta_1$ librating about $0^\\circ$ (e.g., fit IV). If $\\sin i \\approx 0.5$, only good fits with both $\\theta_1$ and $\\theta_2$ librating about $0^{\\circ}$ (e.g., fits II and III) are stable. Fit II is the fit with the smallest libration semiamplitudes of both $\\theta_1$ and $\\theta_2$, with $\\Delta\\theta_1 \\approx 6^\\circ$ and $\\Delta\\theta_2 \\approx 10^\\circ$.  Our analysis differs from that of \\citet{fmb05} in having the additional Keck data, in using $\\chi_\\nu^2$ instead of RMS as the primary measure of the goodness of fit, and in examining systematically the dynamical properties of all of the fits found. Although \\citet{fmb05} showed $e_j$ and $\\omega_j$ of many fits with RMS within $\\approx 0.2\\mps$ of the minimum RMS, they discussed in detail the dynamical properties of only two fits --- their solutions A and B --- with RMS within $0.06\\mps$ of the minimum. Their solutions A and B have $\\omega_1 \\approx 120^\\circ$ and $\\omega_2 \\ga 200^\\circ$ and hence large amplitude librations of both $\\theta_1$ and $\\theta_2$. Because the fit to the CORALIE data with the minimum RMS (which is near the minimum $\\chi_\\nu^2$ fit of \\citealt{may04} at $e_2 = 0.18$ and $\\omega_2 = 237^\\circ$; see Table \\ref{tab:mayor}) is close to the stability boundary in the $e_2$-$\\omega_2$ plane, \\citet{fmb05} were able to find stable fits with RMS within $0.06\\mps$ of the minimum. For the combined data set, the fit with the minimum RMS at $e_2 = 0.22$ and $\\omega_2 = 275^\\circ$ is far from the stability boundary, and none of our fits with RMS within $0.06\\mps$ of the minimum are stable (see Figs. \\ref{fig:param}{\\it b}, \\ref{fig:1.0}{\\it a}, and \\ref{fig:0.5}{\\it a}). However, by examining systematically the $\\chi_\\nu^2$, RMS, and dynamical properties of all of our fits, we have found fits like fit IV (which has stable, large amplitude libration of only $\\theta_1$ if $\\sin i \\approx 1.0$), fit III (which has large amplitude librations of both $\\theta_1$ and $\\theta_2$) and, in particular, fit II (which has small amplitude librations of both $\\theta_1$ and $\\theta_2$), all with RMS within $0.13\\mps$ of the minimum and $\\chi_\\nu^2$ within $0.08$ of the minimum. The stable and unstable regions in the $e_2 \\cos \\omega_2$-$e_2 \\sin \\omega_2$ plane shown in Figures 4 and 8 of \\citet{fmb05} are qualitatively similar to those in our Figures \\ref{fig:1.0}{\\it a} and \\ref{fig:0.5}{\\it a} in the $e_2$-$\\omega_2$ plane, but it should be noted that their figures show simulations of their solution B (for the CORALIE data) with $e_2$ and $\\omega_2$ changed, while our figures show simulations of the fits (to the combined data set) that minimize $\\chi_\\nu^2$ for the given $e_2$ and $\\omega_2$.  The relatively large values of $\\chi_\\nu^2$ ($\\ge 1.84$) and RMS ($\\ge 7.87\\mps$) of the fits presented in this paper suggest that the double-Keplerian model may not fully explain the radial velocity data of HD 82943. Also hinting at the same possibility is the increase in the RMS of the best fit from $6.99\\mps$ (for the extracted CORALIE data alone) to $7.88\\mps$ with the inclusion of the Keck data, which fill in some gaps in the CORALIE data and increase the time span of observations. However, the RMS and $\\chi_\\nu^2$ values are by themselves not very strong evidence, because our estimate ($4.2\\mps$) for the radial velocity jitter has an uncertainty of $\\sim 50\\%$ and the jitter can be $\\sim 6\\mps$. On the other hand, there appears to be systematic deviations of the radial velocity data from the double-Keplerian fits. Figure \\ref{fig:res} shows that the data within about $\\pm 200$ days of JD 2452600 are higher than the fits and that the data near JD 2453181 are lower than the fits. The fact that $\\sim 40\\%$ of the Keck measurements fall in these two regions result in the aforementioned increase in the RMS of the best fit from $6.99\\mps$ for the CORALIE data alone to $7.88\\mps$ for the combined data set. But it is important to note that the Keck and CORALIE data are consistent with each other where they overlap and that they are both higher than the fits in the region around JD 2452600. One possible explanation for the deviations is that they are due to the mutual gravitational interaction of the planets. However, unlike the GJ 876 system where the resonance-induced apsidal precession rate $d\\varpi_j/dt$ is $-41^\\circ \\yr^{-1}$ on average and the precession has been seen for more than one full period \\citep{lau05}, the average apsidal precession rates of, e.g., fits II and III are only $-0\\fdg51\\,(\\sin i)^{-1} \\yr^{-1}$ and $-0\\fdg71\\,(\\sin i)^{-1} \\yr^{-1}$, respectively, and the orbits have precessed a negligible $3$--$4\\,(\\sin i)^{-1}$ degrees over the $6.1\\yr$ time span of the available observations. On the other hand, there are small variations in the orbital elements on shorter timescales due to the planetary interaction that could produce deviations from the double-Keplerian model. Alternatively, the deviations could be due to the presence of additional planet(s). Continued observations of HD 82943 combined with dynamical fits should allow us to distinguish these possibilities."
    },
    "0512/astro-ph0512284_arXiv.txt": {
        "abstract": "We analyze a light curve of the symbiotic star BF Cyg, covering 114 years of its photometric history. The star had a major outburst around the year 1894. Since then the mean optical brightness of the system is in steady decline, reaching only in the last few years its pre-outburst value. Superposed on this general decline are some 6 less intense outbursts of 1-2 magnitude and duration of 2000-5000 days. We find a cycle of ~6376 days, or possibly twice this period, in the occurrence of these outbursts. We suggest that the origin of the system outbursts is in some magnetic cycle in the outer layers of the giant star of the system, akin to the less intense ~8000 days magnetic cycle of our Sun. We further find, that in addition to its well known binary period of 757.3 days, BF Cyg possesses also another photometric period of 798.8 days. This could be the rotation period of the giant star of the system. If it is, the beat period of these two periodicities, 14580 days, is the rotation period of a tidal wave on the surface of the giant. A 4th period of 4436 days, the beat period of the 14580 and the 6376 cycles is possibly also present in the LC. We predict that BF Cyg will be at the peak of its next outburst around the month of May in the year 2007. The newly discovered 798.8 days period explains the disappearance of the orbital modulation at some epochs in the light curve. The 757.3 oscillations will be damped again around the year 2013. ",
        "introduction": "Symbiotic stars (SS) are a class of variable stars consisting of a cool giant, a hotter object, either a hot subdwarf or a compact object, and an emission nebula.   The optical variability of symbiotics may take different forms and time scales. One form is of cyclic variations, due to the varying aspects of the revolving binary system, with or without an apparent eclipse in the light curve (LC). Binary periods of SS are of the order of 1 to a few years. Another type of variability has an explosive character, in the form of a single outburst, as for symbiotic novae, or multiple events. The time scales of these variations are quite long: the  decay times of the outburst of symbiotic novae range between a few months to more than a century. The cool giant in some symbiotic systems shows also intrinsic variability such as radial pulsations of Mira-type.  Variability of SS light on short time scale of minutes and hours has not been reported much in the literature. Recently, however, such variations on this time scale, reminiscing the flickering phenomenon in cataclysmic variables, have been discovered in symbiotics (Sokoloski et al. 2001).   We have analysed anew the historical light curve of BF Cyg. This analysis resulted in the discovery of new elements in the long-term light curve of the star which may give us new clues on the nature of this system.  In this paper  we present this analysis of the long-term  light curve of BF Cyg, covering 114 years of observations. The characteristics  of the symbiotic system BF Cyg are described in Section 2, and the data sets used on our analysis are reviewed in Section 3.  In Section 4 we describe the time series analysis and the periodicities detected.  In Section 5 we discuss the  physical interpretation of these periodicities.   \\section {Brief description of BF Cyg}  BF Cyg is a very bright object (V $\\simeq 9 $). This makes it a popular target for observations on a wide wavelength range. This is also probably the reason why the record of measurements of its magnitude goes back in history for over 115 years (see below). Similar to many SS stars, and in particular to the prototype Z And, the long-term light curve of BF Cyg shows two kinds of optical variability. One is a regular periodic modulation, with large changes in its amplitude. The other type of the long range variability of this star is of explosive character, taking the form of repeating outbursts (see Figure 1).  The cool component is a fairly normal M5 giant and the system is classified as an S-type symbiotic with near-IR colors consistent with those of normal cool giants (Kenyon \\& Fernandez-Castro, 1987; Munari et al. 1992). The IUE  ultraviolet continuum suggests the presence of a hot subdwarf with a temperature $>60000 K$ (Gonz\\'{a}les-Riestra, Cassatella \\& Fernandez-Castro 1990).   Photometric variability with an apparent  modulation period of 754 days was first noted by Jacchia (1941). The system has been extensively studied in the optical and the ultraviolet wavelength ranges (Mikolajewska et al. 1989; Fernandez-Castro et al. 1990; Gonz\\'{a}les-Riestra et al. 1990; Skopal et al. 1997). Its optical and ultraviolet emission lines  vary in phase or/and in anti phase with the photometric minima. The orbital nature  of this variation is confirmed by infrared radial velocities data  and  the spectroscopic orbital period is 757.2, nearly identical to the photometric one (Fekel et al. 2001).   \\section {The long-term light curve of BF Cyg}   We collected data from three large photometric measurement sets retrievable for this system, in order to reconstruct its historical light curve.   A regular photometric monitoring of BF Cyg for the years 1890-1940 is available from the Harvard plates (Jacchia, 1941). While few plates were collected in the first ten years, the data are more frequent after the year 1900.  A second data set is composed of  the photographic  measurements of Skopal et al. (1995).   A third large data set is the visual magnitude estimates collected by the American Variable Stars Association (AAVSO). These are often daily estimates of the magnitude of the system. For homogeneity we averaged the AAVSO data over a time interval of 24 days, similar to the sampling interval of the Jacchia (1941) data set. The AAVSO at our hands is updated to Dec 2004.    \\begin{figure*} \\includegraphics[width=130mm]{figure1.eps} \\caption{A 114 years light curve of the symbiotic star BF Cyg, from the year 1890 up until Dec 2004.  Dots refer to m$_{pg}$ and crosses indicate  visual data transformed to m$_{pg}$. The solid line is a 3d degree polynomial, and a 9 term harmonic wave based on 3 inependent periods, fitted to the data by least squares. See text for further explanations.} \\end{figure*}   In Fig. 1 we present the long-term light curve (LC) of BF Cyg from 1890 up to the present. In this figure, the AAVSO data have been scaled to the photographic scale, by adding a factor .61 that takes into account the (B-V) color of the system at quiescence (Munari et al. 1992) and the transformation B=m$_{pg}$ +0.11 (Allen, 1973). This curve covers the photometric behavior of BF Cyg from 1890 to Dec 2004. There is however a considerable gap in the distribution of the data points between JD 2429986 and JD 2434486 which affects significantly the spectral window function of this time series (see Section 4.1) The solid line in the figure will be explained in Section 4.2. A sample of the data is shown in Table 1. Table 1 in full  will be accessible only electronically.  \\begin{tabular}{@{}llrcc@{}} \\\\ \\multicolumn {2} {|c|}{\\bf Table 1. The m $_{pg}$ or scaled m$_{v}$ data of Fig. 1}\\\\ \\\\ \\\\ Julian day   &  magnitude \\\\ \\\\ 2411547.9&  12.22 \\\\ 2411570.8&  12.70 \\\\ 2411616.4&  12.30 \\\\ 2411684.9&  11.87 \\\\ \\\\ \\end{tabular}  The system underwent a dramatic 3.5 magnitudes brightening event around the year 1894. This outburst is followed by slow fading of the star that continues until the very present time. This behavior is reminiscent of that of the small class of symbiotic novae, that have one single major outburst recorded in their historical  light curve. Actually, there is hardly a typical light curve of symbiotic novae, and the fading from the outburst shows a large variety of behavior. It can be  gradual, such as in RR Tel or AG Peg, or characterized by large light oscillations as in V1329 Cyg  or  by a deep minimum as in PU Vul (Viotti, 1993).  In BF Cyg, a few major events of sudden brightening of the system by 1-2 magnitudes are superposed on the decline from the major 1894 outburst. The duration of such an explosive event may last a few years, and may include episodes of short term flares of brightening by a  few tenths of magnitude lasting a few days. In between outbursts, the system exhibits periods of relative quiescence that last for a few years. In the last eight years BF Cyg is in one of these quiescence states, while the system  is now back to the brightness level measured by Jacchia (1941) on the few Harvard patrol plates obtained in the year 1890  before the large 1894 outburst.   A periodic oscillation of about 754 days was already recognized by Jacchia (1941), with amplitude  $\\sim $1 magnitude. Throughout the years other values of the photometric orbital periodicity have been suggested by various investigators. For example  P=757.3  days (Pucinskas, 1970), P=756.8  days (Mikolajewska et al. 1989), P=757.2  days (Fekel et al. 2001). As already noted  by Jacchia (1941), however, there were epochs in the history of the star, during which the binary modulation was very small, sometime nearly disappearing entirely from the LC. ",
        "conclusions": "\\subsection{The outbursts cycle}  The 6 recorded outbursts of BF Cyg in the last 104 years occurred with a constant time interval of $\\sim$6376 days between them. Periodic or quasi-periodic variations in the light of stars, with periods of thousands of days and amplitudes of 2 or 3 magnitudes are known to exist in the class of semi-regular giants (Mattei et al. 1988; Kiss et al. 1999).   It seems, however, that the 6376 days periodic variability of BF Cyg is of a different nature. The structure of the LC shows the characteristics of outbursts, rather than of pulsations and so it is indeed interpreted by most researchers in the field (e.g. Gonz\\'{a}les-Riestra et al. 1990; Skopal et al. 1997).  In an attempt to understand the origin of the outbursts phenomenon and the nature of the clock that regulates their appearances, the clock of the solar cycle, with its similar time constant of $\\sim$8000 days, comes to mind. We hypothesize that the origin of the cyclic outbursts of BF Cyg lies in some magnetic activity, driven by a magnetic dynamo process in the outer layers of the cool giant component of this stellar system. It is known that the 11/22 year solar cycle modulates the mass flux of the solar wind, among other measured parameters of the Sun. The basic process originates in an interaction between differential rotation and convection motions (Babcock  1961; Ulrich \\& Boyden 2005).  Solar-like cycle in Asymptotic Giant Branch stars has been proposed by Soker (2000) as an explanation to the morphology of a few planetary nebulae. Soker (2002) has also proposed that giant stars that are members of symbiotic systems may harbor a magnetic dynamo process.   The giant in BF Cyg may indeed possess one important ingredient of such a process, namely, a relatively fast rotation. It is classified as a M5 star (Kenyon \\& Fernandez-Castro 1987; M\\H{u}rset \\& Schmid 1999). No data are available for the rotation velocity of such late type stars in the De Medeiros et al. (2000) tables, but from the general trend of the velocity distribution of cool stars, very small mean velocity (2-3 km/sec) is expected. Fekel et al (2001), on the other hand, derive a projected rotational velocity of the giant in BF Cyg of $\\sim$4.5 \\kms. This makes the giant of BF Cyg a fast rotator relative to field giants of similar spectral type. We note also that differential rotation, another important ingredient of the magnetic dynamo mechanism, has been recently measured in some active K-type giants  (Weber, Strassmeier \\& Washuettl 2005).  In the magnetic dynamo scenario, the repeating intensifications in the optical luminosity of the system that take the form of the outbursts, are due to periodic enhancement of the stellar wind from the cool giant of the system, regulated by this dynamo process. This results in an enhancement of the mass accretion rate onto the compact star of the system. The intense optical luminosity originates in the vicinity of the hot component, probably in a bloated gaseous shell around the WD star (Munari 1989).  We proposed in the past a similar solar-like cycle as an explanation for the $\\sim $ 8400 days ($\\sim $23 years) cycle that we discovered in the LC of another symbiotic star - Z And (Formiggini \\& Leibowitz 1994). A comparison between BF Cyg and Z And reveals similarities between these two systems and in their behavior both in quiescence and at outbursts. Their cool component  is a M5 star  (M\\H{u}rset \\& Schmid 1999) and the orbital period is almost the same: 758.8 days for Z And (Formiggini \\& Leibowitz 1994) and 757.3 for BF Cyg (section 4.2). Both systems belong to the classical symbiotics family, and during outbursts their hot component seems to maintain a constant bolometric luminosity while expanding in radius (Mikolajewska \\& Kenyon 1992).  One difference between the 2 systems is that in addition to its sequence of 6 \"small\" outbursts, BF Cyg underwent the dramatic 1894 event of large outburst, quite distinct from the 6 that followed. No such event has been recorded in the history of Z And. In view of the 1894 event, BF Cyg should perhaps be classified as symbiotic nova, as already suggested by Skopal et al. (1997). That event seems to be of the scale and nature that are different from those of the small outbursts. The abrupt increase by 4 magnitudes in its luminosity, and the slow decay lasting over 100 years afterwards, are probably a signature of a thermonuclear runaway process, similar to the events that characterize systems such as AG Peg, V1016 Cyg as symbiotic novae (Mikolajeswka \\& Kenyon 1992).    \\subsection {The brightness oscillations}  Near the P2=757.3 days binary period of BF Cyg we discovered the period P3=798.8 days. This could be the periodicity of oscillations of the M giant star of this system. However, we consider this interpretation unlikely for the following reason. The giant of BF Cyg is not a Mira type star. It does not show modulations in the near infrared as for Mira. In the (J-H) vs. (H-K) color diagram (Whitelock 1994) it does not lie in the region occupied by Miras.   We suggest that the P3 periodicity is the rotation period of the giant of the system. The rotation of the giant component of SS's was  already discussed in the literature (e.g. Munari 1988). In general, the giant stars in symbiotic systems are assumed to rotate synchronously with their binary revolution. This assumption is based on theoretical considerations (Zahn 1977), taking into account the synchronization time scale expected from the values of the radii, masses and binary separations that are typical of symbiotics on one hand, and the estimated ages of these binary systems on the other. The period P3 is close enough to P2 and therefore would not be an unusual exception to this commonly believed rule.   Modulation of the star light at the rotation period of the giant is indeed expected in models explaining the binary photometric variations on the basis of the reflection effect (Kenyon 1986, Formiggini \\& Leibowitz 1990). If the giant star is not strictly isotropic in its photometric characteristics, e.g. if there are dark spots on its photosphere or if there are areas of different reflectivity on its surface, rotation of this star will modulate the luminosity of the system in the rotation frequency.  Modulations at the rotation frequency are expected also in models whereby the photometric binary variations are due to variation in the optical depth towards the main light source of the system that are coupled to the orbital revolution. In such models the WD and its close vicinity are seen by the observer through different regions of the stellar wind of the giant star (Gonz\\'{a}les-Riestra et al. 1990). The stellar wind of the giant is not necessarily isotropic with respect to the giant center. The non isotropic structure of the wind may also be coupled to the giant rotation through the effect of the stellar magnetic field. In such a case, the opacity toward the main light source of the system will vary at the giant rotation frequency, in addition to its variation at the binary frequency.   The P5=14580 days periodicity is the beat period of the binary orbital period P2 and the P3 period. If the latter is  indeed the giant rotation period, P5 is the period of the tidal wave that propagates in the outer layers of the giant, in the coordinate system of the rotating star. We suggested above that a magnetic dynamo process may be the driving mechanism responsible for the outbursts cycle. It is not unreasonable to speculate that a tidal wave in differentially rotating convecting layers may also modulate the magnetic field of the star at the tidal wave frequency, in addition to the modulation by the main cycle of the dynamo. The apparent P4 periodicity, if found to be real, may be the beat period of these two magnetic cycles in the outer layers of the giant.   Finally we note that as mentioned in Section 4.2.1 the structure and the amplitude of the P2 and P3 oscillations are independent of the average luminosity of the system. In particular it appears that the amplitudes of these modulations at maximum light of outbursts are not significantly different from those during quiescence states of the system.   This fact is consistent with the reflection interpretation of the binary modulation (Kenyon 1986, Formiggini \\& Leibowitz 1990). To first approximation the reflection and the reprocessing of the hot component's radiation in the giant atmospheric layers, give rise to optical luminosity that is a given fraction of the hot component luminosity. An increase in the later will result in the same relative increase in the luminosity of the heated hemisphere of the giant, hence the independence of the amplitude of the P2 variations, expressed in magnitude units, on the luminosity of the system.  This independence is also consistent with the modulating optical depth interpretation. A given optical depth of an absorbing layer reduces the flux of the emerging radiation by a given percentage of the flux of the radiation impinging on it, i.e. it has the same amplitude in magnitude units."
    },
    "0512/astro-ph0512417_arXiv.txt": {
        "abstract": "The measurement of the Shapiro time delay in binary pulsar systems with highly inclined orbit can be affected both by the motion of the pulsar's companion because of the finite time it takes a photon to cross the binary, and by the gravitational light bending if the orbit is sufficiently edge-on relative to the line of sight. Here we calculate the effect of retardation due to the companion's motion on various time delays in pulsar binaries, including the Shaipro delay, the geometric lensing delay, and the lens-induced delays associated with the pulsar rotation. Our results can be applied to systems so highly inclined that near conjunction gravitational lensing of the pulsar radiation by the companion becomes important (the recently discovered double pulsar system J0737-3039 may exemplify such a system). To the leading order, the effect of retardation is to shift all the delay curves backward in time around the orbit conjunction, without affecting the shape and amplitude of the curves. The time shift is of order the photon orbit crossing time, and ranges from a second to a few minutes for the observed binary pulsar systems. In the double pulsar system J0737-3039, the motion of the companion may also affect the interpretation of the recent correlated interstellar scintillation measurements. Finally, we show that lensing sets an upper limit on the magnitude of the frame-dragging time delay caused by the companion's spin, and makes this delay unobservable in stellar-mass binary pulsar systems. ",
        "introduction": "}   Timing study of binary radio pulsars has provided some of the best tests of the general relativity (GR) to date.  Among the effects that have been probed in different systems are gravitational redshift, precession of the periastron and orbital decay due to gravitational wave emission. Strong evidence for the geodetic precession -- GR manifestation of the spin-orbital coupling of the binary -- has been found in several systems (e.g., \\cite{stairs}, \\cite{hotan}).  In binaries with highly inclined orbits, it has been possible to measure the Shapiro delay caused by the propagation of the pulsar signal in the companion's gravitational field (e.g., \\cite{lyne}, \\cite{splaver}, \\cite{jacoby}).  In systems that have nearly edge-on orbital orientation, like the recently discovered double pulsar system J0737-3039 \\cite{lyne}, one may also look for the effects related to the gravitational bending of the pulsar radio beam as it passes close to the companion around the moment of the pulsar's superior conjunction (\\cite{schneider}, \\cite{dorkop}, \\cite{LR05}, \\cite{RL05}). Light bending modifies the conventional Shapiro delay formula widely used in pulsar timing studies and introduces additional geometric time delay (\\cite{schneider}, \\cite{LR05}). Bending also gives rise to a {\\it lens-rotational delay} associated with the rotating beam of the pulsar and induces distortion of the pulse profile shape near conjunction \\cite{RL05}.  Our previous studies of the lensing effects in binary pulsar systems (\\cite{LR05}, \\cite{RL05}) adopted a ``static lens'' approximation, in which the companion motion is neglected. Kopeikin \\& Schafer \\cite{kopshaf} showed that the orbital motion of the pulsar companion gives rise to a $v_c/c$ correction ($v_c$ is the companion velocity) to the conventional Shapiro delay formula and such correction might be observable under some circumstances. The work of \\cite{kopshaf}, however, neglected light bending and thus cannot be applied to the situation where the light ray passes the companion within a few Einstein radii. On the other hand, only for such highly inclined systems are the lensing effects potentially observable (\\cite{LR05}, \\cite{RL05}).  In this paper we study the effect of companion motion on the time delays associated with the gravitational light bending near the companion. These include both the lensing corrections to the delays independent of the pulsar spin -- Shapiro and geometric delays (see \\S \\ref{sect:timing_no_spin}), and the lens-rotational time delays (\\S \\ref{sect:spin}). Our formulae are general and valid in both the ``strong lensing'' and the ``weak lensing'' regimes. We also provide a simple and general interpretation of the influence of the lens motion on the behavior of various timing signals in binary pulsar systems (\\S \\ref{sect:implications}) which has not been previously discussed.  Qualitatively, we expect that near binary conjunction the effect of moving lens amounts to a time shift in various time delay curves. It takes finite amount of time for the pulsar signal to travel across the binary system during which the position of its companion changes. Based on simple geometric considerations, this time shift is $(M_p/M)a_\\parallel/c$, where $M_p$ is the pulsar mass, $M$ is the total mass of the system, and $a_\\parallel$ is a separation between binary components projected along the line of sight [see eq. (\\ref{eq:a_par})]. In the limit $M_p\\to 0$ this retardation delay goes to zero since companion does not move. Our explict calculations and general expressions of various time delays confirm this simple expectation, see \\S \\ref{sect:implications}.  For completeness,  in \\S \\ref{sect:FD} we discuss the effect of frame dragging due to the companion's spin on the time delay. This type of delay was previously investigated by \\cite{laguna}, \\cite{wex}, \\cite{tart}, \\cite{rugg}, who disregarded light bending in their studies. We incorporate the lensing effect in our calculation of the frame-dragging delay, and show (\\S \\ref{subsect:fd_interp}) that this delay is unobservable in the binary pulsar systems with stellar mass companions. ",
        "conclusions": "}  In this paper we have studied the combined effects of gravitational lensing and companion motion on various time delays in binary pulsar systems. Our study improves upon previous work on the lensing effect by going beyond the static approximation (\\cite{LR05}, \\cite{RL05}). We show that for highly inclined systems, the companion motion affects the Shapiro delay and other lensing-related delays by shifting the delay curves backward in time by the amount $\\sim a_c/c$ (where $a_c$ is the semimajor axis of the companion).  To detect this effect on time delays, one should look for a systematic displacement of the peak of the Shapiro delay curve from the exact moment of the superior conjunction of the pulsar. To this end, one must be able to pinpoint the exact moment of the conjunction and to nail down the peak of the Shapiro delay curve. This may not be so easy since in practice one bins the timing data within rather large time intervals to get a good measurement of the Shapiro delay. Thus, one would need repeated measurements of many binary pulsar orbits to be able to reach the accuracy needed for measuring the retardation effect. Retardation effect on lensing delays {\\it may} be detectable in the double pulsar system J0737-3039, but the feasibility of this measurement depends on the currently uncertain inclination of the system and is compromised by the eclipse of the millisecond pulsar near its superior conjunction by the magnetosphere of its companion  (\\cite{lyne}, \\cite{mclaugh}).  In the case of the double pulsar J0737-3039 system, the retardation effect due to the companion motion not only affects timing of the millisecond pulsar but may also influence the interpretation of the scintillation measurements of this system near conjunction \\cite{coles}. Indeed, the correlated measurements of the interstellar scintillations of the two pulsars provide information about the system orientation and properties of the interstellar turbulence only when the positions of both pulsar projected on the sky plane are accurately known at each moment of time. In the static approximation the projected separation of pulsars is $\\bR$, while it becomes $\\bb_0$ [see eq.~(\\ref{eq:bb_0})] when the retardation effect is taken into account. For the currently inferred inclination angle from the scintillation measurement itself ($|i-90^\\circ|=0.29^\\circ\\pm 0.14^\\circ$, \\cite{coles}), both the retardation effect and the light bending distortion of the apparent path of the lensed millisecond pulsar \\cite{LR05} have similar orders of magnitudes and contribute at the level of $\\sim R_E$ to the projected pulsar separation at conjunction. As for this system $R_E=2550$ km is not very different from the inferred minimum projected approach distance of the two pulsars, $\\approx 4000\\pm 2000$ km, one expects that both effects may at some level affect the interpretation of the correlated scintillation measurements of \\cite{coles}. Whether this can explain the discrepancy between the values of the system's inclination obtained using the Shapiro delay ($i=88.7^\\circ\\pm 0.9^\\circ$; \\cite{lyne}, \\cite{ransom}) and scintillations is not clear at present."
    },
    "0512/astro-ph0512621_arXiv.txt": {
        "abstract": "% We show that with a Next Generation Large Telescope one can detect the accelerated motions of $\\sim 100$ stars orbiting the massive black hole at the Galactic center. The positions and velocities of these stars will be measured to astrometric and spectroscopic precision several times better than currently attainable enabling detailed measurements of the gravitational potential in the neighborhood of the massive black hole. We show that the monitoring of stellar motions with such a telescopes enables: (1)  a measurement of the Galactic center distance $R_0$ to better than 0.1\\% accuracy, (2) a measurement of the extended matter distribution near the black hole, including that of the exotic dark matter, (3) a detection of general relativistic effects due to the black hole including the prograde precession of stars and possibly the black hole spin, and (4) a detection of gravitational encounters between monitored stars and stellar remnants that accumulate near the Galactic center. Such encounters probe the mass function of the remnants. ",
        "introduction": "Near-infrared monitoring of stellar sources in the inner $1\\arcsec$ of the Galaxy with speckle imaging and adaptive optics techniques has enabled complete orbital reconstruction of several stellar sources orbiting the $\\sim 4 \\times 10^6 M_\\odot$ black hole at the Galactic center \\citep{Schoedel:02, Schoedel:03, Ghez:05}. Sources have been monitored with astrometric errors of a few milli-arcseconds, and radial velocity errors $< 50 \\textrm{ km s}^{-1}$  (Eisenhauer et al. 2003; Ghez et al. 2003a), allowing the detection of the accelerated proper motions of $\\sim 10$ stars. The star S0-2 has the shortest orbital period of $\\sim15 \\textrm{ yr}$. With current orbital data, the distance to the Galactic center $R_0$ is measured to an accuracy of $\\approx5\\%$ (\\citealt{Eisenhauer:03, Eisenhauer:05}; see \\citealt{Salim:99} for a discussion of the physical principles underlying the measurement of $R_0$ from the orbital data).  Although the stellar orbits have provided unequivocal proof of the existence of a massive black hole at the Galactic center, the matter content in the vicinity of the black hole remains largely unknown.  Due to the short radial diffusion time as compared with the age of the bulge, a large number of massive compact remnants ($5-10M_\\odot$ black holes) could have segregated into the dynamical sphere of influence of the black hole \\citep{Morris:93,Miralda:00}.  Furthermore, adiabatic growth of the massive black hole could have compressed a pre-existing distribution of cold dark matter (CDM) \\citep{Ipser:87,Quinlan:95,Gondolo:99} and stars \\citep{Peebles:72,Young:80} into a density ``spike'', although a variety of dynamical processes are capable of destroying such a spike \\citep{Ullio:01,Merritt:02,Gnedin:04,Merritt:04}. If the CDM consists of weakly-interacting massive particles (WIMPs), a sustained CDM spike might produce detectable WIMP annihilation radiation.  We show here that the finer angular resolution and increased light-gathering power of a Next Generation Large Telescope (NGLT) will enable the detection of accelerated proper motions of $\\sim 100$ stars with astrometric and spectroscopic errors several times smaller than currently possible. By fitting to mock NGLT orbital data we determine the constraints such orbits place on the gravitational potential near the massive black hole at the Galactic center (see also \\citealt{Weinberg:05}). We demonstrate that the observations will yield a measurement of the density profile in a dark matter spike. They will also reveal the lowest order general relativistic effects and possibly higher order effects such as the black hole spin. We show that $R_0$ will be measured to better than $0.1\\%$ precision, helping to place tight constraints on the shape of the Galactic dark matter halo. We also calculate the rate at which perturbations from massive remnants that accumulate near the Galactic center deflect the observed stellar orbital motions and describe how the remnants' mass function can thereby be extracted from the monitoring data. We find that if stellar-mass black holes dominate the matter density then in 10 yr of monitoring at an astrometric resolution of $0.5 \\textrm{ mas}$ approximately 10\\% of all monitored stars will experience detectable deflections in their orbital motions. ",
        "conclusions": "We have shown that the monitoring of stellar orbits around the massive black hole at the Galactic center at the high astrometric and spectroscopic resolution attainable with an NGLT enables one to probe the deep gravitational potential of the region. Many exciting measurements are achievable even for modest (factor of a few) improvements over the astrometric capabilities of current 10 meter class telescopes. The future success of Galactic center research greatly depends on advances in the astrometric capabilities of the next generation of telescopes."
    },
    "0512/hep-th0512202_arXiv.txt": {
        "abstract": "Gravitational collapse is analyzed in the Brane-World by arguing that regularity of five-dimensional geodesics require that stars on the brane have an atmosphere. For the simple case of a spherically symmetric cloud of non-dissipating dust, conditions are found for which the collapsing star evaporates and approaches the Hawking behavior as the (apparent) horizon is being formed. The effective energy of the star vanishes at a finite radius and the star afterwards re-expands and ``anti-evaporates''. Israel junction conditions across the brane (holographically related to the matter trace anomaly) and the projection of the Weyl tensor on the brane (holographically interpreted as the quantum back-reaction on the brane metric) contribute to the total energy as, respectively, an ``anti-evaporation'' and an ``evaporation'' term. ",
        "introduction": "\\label{intro} Black holes (BHs) are unstable in four (and higher) dimensions because of the Hawking effect~\\cite{hawking}, which is deeply linked to the trace anomaly of radiation fields~\\cite{birrell}. In the Randall-Sundrum (RS) Brane-World (BW) models~\\cite{RS}, a collapsing homogeneous star likewise requires a non-static exterior~\\cite{BGM}. Further, forcing a static exterior induces a trace anomaly of the same form as that of semiclassical BHs, although with opposite sign, which suggested that BH solutions of the bulk equations correspond to quantum corrected (semiclassical) BHs on the brane~\\cite{tanaka1,fabbri}, in the spirit of the holographic principle~\\cite{holography} and AdS/CFT conjecture~\\cite{AdSCFT}. \\par The junction conditions~\\cite{israel}, which preserve the regularity of (five-dimensional) geodesics, cannot allow a step-like discontinuity (e.g.~across the star surface) and a $\\delta$-like discontinuity (e.g.~across the brane) in the stress tensor at the same location, hence discontinuities in the stress tensor of brane stars are not mathematically permitted. This can be physically understood by considering that the brane thickness (of the order of the AdS length $\\ell\\sim\\lambda^{-1/2}$, $\\lambda$ being the brane tension) and the star's atmosphere cannot be both negligibly thin, and it is indeed natural to assume that the latter is much larger than $\\ell$. In Ref.~\\cite{PTP}, we employed effective four-dimensional (hydrodynamical) equations~\\cite{maart} in order to study a ``corrected'' Oppenheimer-Snyder (OS) model~\\cite{OS} in which the star is divided into three regions (see Fig.~\\ref{figura}): a homogeneous ``core'' with most of the energy; a ``transition region'' of fast density decrease which connects with a ``tail'', where the energy density slowly vanishes. The tail and transition region together form the ``BW atmosphere'', which disappears in the limit $\\lambda\\to\\infty$. \\par In general, effective equations cannot determine the brane metric uniquely unless one also knows the bulk geometry. However, for negligible dissipation and asymptotically flat brane, the evolution of the system is uniquely determined by the dynamics of the homogenous core, which corresponds to an exact five-dimensional solution~\\cite{langlois} and reproduces the trace anomaly of quantum field theory~\\cite{birrell}. Our main results are that {\\em the total energy of the system is conserved\\/} and that {\\em the collapsing star ``evaporates'' until the core experiences a ``rebound'' in the high energy regime (when its energy density is comparable with $\\lambda$), after which the whole system ``anti-evaporates''\\/}. This behaviour cannot be related to GR perturbatively (in $\\epsilon\\equiv {\\rho_0/ \\lambda}$, where $\\rho_0$ is the initial core density), but it seems in agreement with the uncertainty principle of quantum mechanics~\\cite{impBO}. Moreover, we can find a range of parameters for which the minimum radius of the collapsing core is larger than $\\ell$ (which sets the scale of Quantum Gravity in the BW), thus supporting the holographic interpretation of the model. ",
        "conclusions": "\\label{conc} Inspired by the conjecture that classical BHs in the BW may reproduce the semiclassical behavior of four-dimensional BHs, we have studied the gravitational collapse of a spherical star of dust in the RS scenario. Regularity of the bulk geometry requires continuity of the matter stress tensor on the brane and leads to a loss of mass from the boundary of the star. We found that the system of effective BW equations is closed to our level of approximation and leads to the collapsing dust star emitting a flux of energy which, at relatively low energies, approaches the Hawking behavior when the (apparent) horizon is being formed. However, no real space-time singularity forms since the star effective mass vanishes at finite star radius, thus leaving a remnant which re-expands by absorbing back the emitted radiation. In a more realistic case, one however expects that dissipation cannot be neglected and the process then becomes irreversible. \\par Let us finally point out that all the above features were obtained for BHs formed by gravitational collapse, excluding therefore primordial BHs."
    },
    "0512/astro-ph0512635_arXiv.txt": {
        "abstract": "% For centuries, our knowledge of planetary systems and ideas about planet formation were based on a single example, our solar system. During the last thirteen years, the discovery of $\\simeq170$ planetary systems has ushered in a new era for astronomy. I review the surprising properties of extrasolar planetary systems and discuss how they are reshaping theories of planet formation. I focus on how multiple planet systems constrain the mechanisms proposed to explain the large eccentricities typical of extrasolar planets.  I suggest that strong planet-planet scattering is common and most planetary systems underwent a phase of large eccentricities. I propose that a planetary system's final eccentricities may be strongly influenced by how much mass remains in a planetesimal disk after the last strong planet-planet scattering event. ",
        "introduction": "For centuries, theories of planet formation had been designed to explain our own Solar Systems, but the first few discoveries of extrasolar planetary systems were wildly different than our own. These discoveries led to the realization that planet formation theory must be generalized to explain a much wider range of planetary systems.  For example, traditional theories predicted that giant planets would form at several AU and beyond, where temperatures are cold enough for ices to initiate the growth of grains and planetesimals (Lissauer 1993, 1995).  Now, we know of over 70 giant planets inside 1 AU and 40 inside 0.1AU (http://www.obspm.fr/planets). Theorists have proposed numerous possible mechanisms to explain the existence of these planets.  Typically, they assume that the giant planet formed beyond a few AU, but then migrated inwards through a protoplanetary or planetesimal disk to their currently observed locations (e.g., Goldreich \\& Tremaine 1980; Lin et al.\\ 1996; Ward 1997; Murray et al.\\ 1998; Cionco \\& Brunini 2002) and stop before being accreted on the star (e.g., Trilling et al.\\ 1998; Ford \\& Rasio 2006).  Similarly, it had long been assumed that planets formed in circular orbits due to strong eccentricity damping in the protoplanetary disk and remained on nearly circular orbits (i.e., eccentricity $\\le$0.1; Lissauer 1993, 1995).  However, over half of the extrasolar planets beyond 0.1AU have eccentricities $\\ge$0.3, and one is as large as $\\simeq$0.95.  Theorists have suggested numerous mechanisms to excite the orbital eccentricity of giant planets (e.g., Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Lin \\& Ida 1997; Holman et al.\\ 1997; Murray et al.\\ 1998; Ford, Havlickova \\& Rasio 2000; Kley 2000, 2004; Chiang \\& Murray 2002; Lee \\& Peale 2002; Marzari \\& Weidenschilling 2002; Ford, Rasio \\& Yu 2003; Adams \\& Laughlin 2003; Veras \\& Armitage 2004; Namouni 2005).  In recent years, improved observations of a few multiple planet systems have allowed theorists to determine their current orbital configuration and use that to place strong constraints on the formation of a few planetary systems (Lee \\& Peale 2002; Ford, Lystad \\& Rasio 2005).  We review some of the mechanisms proposed to explain orbital migration in disks in \\S2 and eccentricity excitation in \\S3.  In \\S4, we review the current knowledge of three particularly well-studied multiple planet systems.  We conclude with a discussion of the implications of these multiple planet systems for theories of orbital migration in \\S5. ",
        "conclusions": ""
    },
    "0512/astro-ph0512129_arXiv.txt": {
        "abstract": "We use a semi-analytical approach to simulate absorption spectra of QSOs at high redshifts with the aim of constraining the cosmic reionization history.  We consider two physically motivated and detailed reionization histories: (i) an Early Reionization Model (ERM) in which the intergalactic medium is reionized by PopIII stars at $z\\approx 14$, and (ii) a more standard Late Reionization Model (LRM) in which overlapping, induced by QSOs and normal galaxies, occurs at $z\\approx 6$.  From the analysis of current Ly$\\alpha$ forest data at $z < 6$, we conclude that it is impossible to disentangle the two scenarios, which fit equally well the observed Gunn-Peterson optical depth, flux probability distribution function and dark gap width distribution.  At $z>6$, however, clear differences start to emerge which are best quantified by the dark gap and peak width distributions.  We find that 35 (zero) per cent of the lines of sight within $5.7< z <6.3$ show dark gaps widths $>50$\\AA \\ in the rest frame of the QSO if reionization is not (is) complete at $z \\gtrsim 6$.  Similarly, the ERM predicts peaks of width $\\sim 1$\\AA \\ in 40 per cent of the lines of sight in the redshift range $6.0-6.6$; in the same range, LRM predicts no peaks of width $>0.8$\\AA. We conclude that the dark gap and peak width statistics represent superb probes of cosmic reionization if about ten QSOs can be found at $z > 6$. We finally discuss strengths and limitations of our method. ",
        "introduction": "After the recombination epoch at $z\\sim 1100$, the universe remained almost neutral until the first generation of luminous sources (stars, accreting black holes, etc) were formed. The photons from these sources ionized the surrounding neutral medium and once these individual ionized regions started overlapping, the global ionization and thermal state of the intergalactic gas changed drastically. This is known as the reionization of the universe, which has been an important subject of research over the last few years, especially because of its strong impact on the formation and evolution of the first cosmic structures (for a comprehensive review on the subject of reionization and first cosmic structures, see \\citeNP{cf05a}).  Calculations based on the hierarchical structure formation in cold dark matter (CDM) models, predict that reionization should naturally occur somewhere between $z \\sim 6-15$ \\cite{co00,gnedin00}.  Recent observational progresses, however, seem to indicate that the reionization history at high redshifts might have been a complicated process.  To start with, WMAP observations of the temperature-polarization cross correlation at large scales for the CMB suggest a Thomson optical depth as high as $\\tau_e\\sim 0.17$, which implies that reionization might have occurred at redshifts as high as $z\\sim 15$ \\cite{ksb++03,svp++03}. On the other hand, the spectroscopy of the Ly$\\alpha$ forest for QSOs at $z > 6$ discovered by the Sloan Digital Sky Survey (SDSS; \\citeNP{fnl+01,fss++03,f05}) seems to indicate that the ionization state of the intergalactic medium (IGM) might be very different along different lines of sight.  For example, the analyses of the spectrum of the most distant known quasar (SDSS J1148+5251) show some residual flux both in the Ly$\\alpha$ and Ly$\\beta$ troughs, which when combined with Ly$\\gamma$ region \\cite{fo05}, imply that this flux is consistent with pure transmission.  The presence of unabsorbed regions in the spectrum corresponds to a highly ionized IGM along that particular line of sight.  However, \\citeN{bfw++01} detected a complete Gunn-Peterson trough in the spectrum of SDSS J1030+0524 ($z=6.28$), where no transmitted flux is detected over a large region (300 \\AA) immediately blueward of the Ly$\\alpha$ emission line.  This result has been shown to be consistent with a hydrogen neutral fraction $f_{\\rm HI} \\gtrsim 10^{-3}$ (\\citeNP{fnswbpr02}, hereafter F02).  There have been other different approaches to investigate the neutral hydrogen fraction. \\citeN{wlc05} and \\citeN{wl04} have estimated the sizes of the ionized regions around 7 QSOs at $z > 6$ (which included the above cited QSO).  Following that they considered the neutral gas surrounding the QSO as a function of different parameters: the Str\\\"omgren sphere size $R_S$, the quasar's production rate of ionizing photons $\\dot{N}_{\\rm phot}$, the clumping factor of the gas $C$ and the age of the QSO $t_{\\rm age}$. According to their arguments, the small sizes of the HII regions ($\\sim 10$ physical Mpc) imply that the typical neutral hydrogen fraction of the IGM beyond $z\\sim 6$ is in the range 0.1 - 1.  However, this approach is weighted down by several uncertainties.  For example, one of the uncertainties is the quasar's production rate of ionizing photons $\\dot{N}_{\\rm phot}$ as it depends on the shape of the spectral template used. Moreover it is implicitly assumed in the modelling of clumping factor that the formation of quasars and galaxies were simultaneous. This in turn implies that quasars ionize only low density regions and hence the clumping factor which regulates the evolution of the HII regions is low. If, instead, stars appears much earlier than QSOs, the quasars have to ionize high density regions, which means that one should use a higher value of clumping factor in the calculations \\cite{yl05}.  \\citeN{mh04} have used a different approach based on the damping wings of the neutral hydrogen. Using density and velocity fields obtained by hydrodynamical simulation, they computed the Ly$\\alpha$ absorption as a function of wavelength. In this case the neutral hydrogen fraction, $\\dot{N}_{\\rm phot}$ and $R_S$ are treated as free parameters, constrained by matching the optical depth observed in the QSO SDSS J1030+0524. Also in this case they find a neutral hydrogen fraction larger than 10 per cent, i.e. the IGM is significantly more neutral at $z\\sim 6$ than the lower limit directly obtainable from the GP trough of the QSO spectrum ($10^{-3}$).  However this result is based only on one quasar. Moreover the observational constraints on the optical depth are very uncertain and can introduce errors in the estimates of $f_{\\rm HI}$.  The fact that we find transmission along some lines of sight while the medium seems quite neutral along others possibly has been interpreted that the IGM ionization properties are different along different lines of sight at $z \\gtrsim 6$, thus suggesting that we might be observing the end of the reionization process.  Finally, the evolution of the luminosity function of Ly$\\alpha$ emitters \\cite{maro04,fzh05,hcen05} suggests that the neutral fraction of hydrogen at $z=6.5$ should be greater than 50 per cent \\cite{maro05}.  From a theoretical point of view several efforts have been made in the last few years in order to reconcile WMAP and SDSS measurements (e.g. \\citeNP{cfo03,hh03,gnedin04}). In particular, Choudhury \\& Ferrara (2005, hereafter CF05) \\nocite{cf05} have showed that a self-consistent model for reionization, which agrees with various sets of observations over a wide redshift range, predicts a highly ionized universe at $z \\approx 6$.  According to the model, the rise in the GP optical depth towards $z=6$ is achieved by assuming a drop of the photoionization rate caused by the disappearance of first generation metal free (hereafter PopIII) stars.  To explore this idea, it becomes important to verify whether the difference in the GP trough along different lines of sight at $z > 6$ can be explained by assuming a highly ionized universe. In other words we pose the question: is it possible that the IGM is overall in a highly ionized state and the differences observed in QSO spectra arise simply due to the cosmic variance in density fluctuations?   Since at redshifts larger than 5, the transmission in the spectra is extremely low, the mean transmitted flux being $F_{\\rm mean} \\lesssim 0.182$ (\\citeNP{songaila04}, hereafter S04), it is not possible to study the properties of the Ly$\\alpha$ forest through the usual Voigt profile analysis.  Alternative quantities have been used to characterize the spectra and to compare them with models, such as the probability distribution function (PDF) of the flux (F02; \\citeNP{sc02}, hereafter SC02) and more importantly, the distribution of the dark gap widths.  It turns out that the length of dark gaps at high redshifts can be quite large ($\\sim 80$ Mpc comoving, which corresponds to $\\sim 30$\\AA \\ in the rest frame wavelength at redshift 6), and hence the modelling requires a large sample of very long lines of sight (say, $\\gtrsim 100$ Mpc).  Such requirement is beyond the reach of current numerical simulations and can only be fulfilled through semi-analytical calculations.  In this paper we use semi-analytical techniques to produce artificial spectra of the Ly$\\alpha$ forest in the redshift range $5.0<z<6.6$. We have considered different reionization scenarios, namely, (i) an early reionization model (CF05), characterized by an highly ionized IGM for $z \\lesssim14$, and (ii) a late reionization model, in which the overlapping of ionized regions is completed only by $z\\sim 6$.  The main goal of this paper is to identify a statistics for the Ly$\\alpha$ forest at $z > 6$ which can be used as an effective tool for discriminating between the two scenarios above.  The outline of this paper is as follows. Section 2 presents the formalism to simulate the Ly$\\alpha$ flux from the IGM density and peculiar velocity fields. In addition, we also describe the two reionization scenarios in detail. Section 3 contains the comparison of our models with observational data at $z < 6$.  We introduce various statistical tools for the Ly$\\alpha$ forest and use our models to make testable predictions for $z \\gtrsim 6$. Finally we summarize our conclusions in Section 4. ",
        "conclusions": " (i) the Ly$\\alpha$ forest observations at $z<6$ are unable to discriminate early vs. late reionization scenarios; (ii) the same data cannot exclude that reionization took place as early as by $z\\approx 14$.  In order to make progress higher redshift quasar spectra are necessary, which are likely to become soon available as SDSS is expected to find $\\sim 20$ luminous quasars in the redshift range $6<z<6.6$. By extending our model predictions to higher redshifts we find that: (i) The mean and the PDF of the transmitted flux are essentially useless to constrain the ionization state at $z \\gtrsim 6$ as most of the pixels are consistent with zero transmission (independent of the ionization state), i.e. in practice these statistics probe the noise distribution; (ii) the dark gap width distribution (DGWD) is very sensitive to the reionization history. We expect at least 30 per cent of the lines of sight (accounting for statistical errors) in the range $z= 5.7-6.3$ to have dark gaps of widths $>50$ \\AA \\ (in the QSO rest frame) if the IGM is in the pre-overlap stage at $z \\gtrsim 6$, while no lines of sight should have such large gaps if the IGM is already ionized. The constraints become more stringent at higher redshifts. We find that in order to discriminate between early and late reionization scenarios 10 QSOs should be sufficient for the DGWD to give statistically robust results. (iii) The statistics of the peaks in the spectra represents an useful complement to the dark gaps and can put additional constraints on the ionization state. As for the DGWD, we find that this statistics constrains reionization models more efficiently at high redshifts. In particular, if the universe is highly ionized at $z\\sim 6$, we expect to find peaks of width $\\sim 1$\\AA \\ in 40 per cent of the lines of sight, in the redshift range $6.0-6.6$; on the contrary, the LRM predicts no peaks larger than 0.8 \\AA.  As an independent check of the models, we have extended all the above statistics to Ly$\\beta$ regions. It turn out that this diagnostics is less powerful than the analog Ly$\\alpha$ one to probe the ionization state of the IGM. Moreover, since the Ly$\\beta$ cross section is 5.27 times smaller than Ly$\\alpha$ one, the flux is always higher in the Ly$\\beta$ region than in the Ly$\\alpha$ forest. This implies that to obtain Ly$\\beta$ constraints as stringent as those from Ly$\\alpha$, requires the analysis of QSOs spectra for $z > 6.6$.   We would like to comment on some additional issues concerning LRM. As discussed in the text, the hydrogen distribution in the LRM for low density IGM is characterized by two distinct phases at $z\\gtrsim 6$, namely an ionized and a neutral phase. To model this two-phase IGM we have studied different topologies of neutral regions. Interestingly, the main conclusions of our work remain unchanged (see for instance Figure \\ref{pmax63_comp}) irrespective of whether we assume that the positions of the neutral regions are completely random (LRM) or we correlate the HI regions along different lines of sight with the density field (LRMd). This result is basically due to the damping wings of neutral regions, which are able to suppress the flux in regions of the spectra that are fully ionized (See Figure \\ref{zoom}). On the other hand if the suppression of the flux does not necessarily correspond to the presence of neutral regions, it implies that QSO spectra might not be very useful to study in details the topology of the neutral hydrogen.  However it is still possible to get some idea about the clustering of the neutral regions {\\it provided we know the evolution of the volume filling factor of ionized regions reasonably well}. We have studied an alternative distribution of the neutral regions, called LRMc, where we assume that neutral regions form the largest possible coherent structure along the line of sight (sometimes as large as 100 Mpc comoving which corresponds to almost $1/3$ of the box). Because of such high clustering, large volumes of IGM are left ionized, resulting in a large fraction of lines of sight which do {\\it not} encounter any neutral region at all. Consequently, the distribution of the largest dark gap widths is biased towards lower widths compared to LRM. This means that the statistics of the largest dark gaps could also give an idea of the clustering in the HI regions. Moreover, as is well known, the 21 cm signal from neutral hydrogen is sensitive to distribution of the HII regions \\cite{fhz04,fzh04}. Hence 21 cm maps could be promising to study the correlation between neutral regions and to obtain a more detailed and quantitative analysis of the size of neutral regions.  Comparing the LRM and LRMc, we also find that {\\it a large gap does not necessarily correspond to a large neutral region}. In fact smaller regions of neutral hydrogen (of sizes $\\lesssim 1$ comoving Mpc) dispersed along the line of sight are more effective in suppressing the flux (because of damping wings) and thus creating large dark gaps in the absorption spectra compared to the larger clustered regions. Probing such small regions is quite difficult with cosmological simulations as they are close to the resolution limits, thus semi-analytical studies can be more helpful in such cases. Our method, in fact,  does not suffer of spurious resolution effects. At high redshift, the length of dark gaps can be $\\gtrsim 60$ Mpc and hence the analysis requires a large sample of very long lines of sight. In order to create realizations of such long lines of sight, numerical simulations typically sample different regions of the box more than once (the so-called ``oversampling'' effect; \\citeNP{pn05}) or combine various spectra of smaller sizes end-to-end (F02). It is difficult to obtain the distribution of very large gaps (which are much larger than the box sizes) from such procedures as multiple ray passages through the same box could produce spectacular spurious artifacts in the gap statistics. For example, we find a much better match with the observations of dark gap width distribution when compared to the simulations of \\citeN{pn05}, who have used a box of size 6.8 $h^{-1}$Mpc.  However, our method suffers from some limitations which are worth noting. First, we are not able to tackle the non-linearities in any self-consistent formalism -- instead we assume a density distribution for the baryons (lognormal, in this case). Since the Ly$\\alpha$ and Ly$\\beta$ forests in the QSO absorption spectra arise from mostly quasi-linear regime, the approximation should be reasonable for computing the transmitted flux. Second, it is nearly impossible to include full radiative transfer effects in our computation of the distribution of the neutral regions and also we are not able to take into account the clustering of sources which is crucial to understand the properties of ionized bubbles. However it is most likely that the location of ionizing sources might not be significantly correlated with neutral regions, particularly when one is dealing with high values of filling factor, as in our case (see, for example, the maps in Figure 1 of \\citeNP{cfw03}). Anyway it would be interesting to combine our approach in the distribution of neutral regions with radiative transfer simulations for a more detailed analysis of the absorption spectra, particularly in the vicinity of the QSO (which corresponds to a highly non-linear structure which is beyond the validity of the lognormal approximation).  In the future, it would be most interesting to check our predictions against a large sample of high signal-to-noise QSO data at redshift $>6$. Our results show that in that case the dark gap statistics would provide a robust and independent probe of the reionization history."
    },
    "0512/astro-ph0512403_arXiv.txt": {
        "abstract": "Kashlinsky et al. (2005) find a significant cosmic infrared background fluctuation excess on angular scales $\\simgt 50$ arcsec that cannot be explained by instrumental noise or local foregrounds. The excess has been tentatively attributed to emission from primordial very massive (PopIII) stars formed $\\le 200$~Myr after the Big Bang. Using an evolutionary model motivated by independent observations and including various feedback processes, we find that PopIII stars can contribute  $< 40$\\%  of the total background intensity ($\\nu J_\\nu \\sim 1-2$ nW m$^{-2}$ sr$^{-1}$ in the 0.8-8 $\\mu$m range) produced by all galaxies (hosting both PopIII and PopII stars)  at $z\\ge 5$. The infrared fluctuation excess is instead very precisely accounted by the clustering signal of galaxies at $z\\ge 5$, predominantly hosting PopII stars with masses and properties similar to the present ones. ",
        "introduction": "Observations of the infrared background provide important information on the emission of cosmic luminous sources throughout the history of the Universe. It has been suggested (Santos, Bromm \\& Kamionkowski 2003; Salvaterra \\& Ferrara 2003) that a large fraction of the measured Near-InfraRed (1-10 $\\mu$m) cosmic Background (NIRB) arises from redshifted Ly$\\alpha$ line photons and nebular emission produced by the first very massive metal-free stars. This hypothesis, however, is very demanding in terms of the required conversion efficiency of baryons into stars (Madau \\& Silk 2005). A large NIRB contribution from such stars has more recently been rejected by the paucity ($\\le 3$) of $z\\sim 10$ candidate sources in Hubble Space Telescope ultra-deep observations (Salvaterra \\& Ferrara 2005). Nevertheless, a more modest contribution from very high redshift galaxies, whose clustering should leave a distinct signature on small-scale angular fluctuations of the background light (Magliocchetti, Salvaterra \\& Ferrara 2003; Kashlinsky et al. 2004; Cooray et al. 2004), is still possible.  Kashlinsky et al. (2005) have recently found significant NIRB fluctuations in deep exposure data obtained with Spitzer/IRAC (Fazio et al. 2004a, 2004b) in four channels (3.6, 4.5, 5.8, and 8 $\\mu$m), after Galactic stars and galaxies bright enough to be individually resolved by the instrument have been carefully subtracted. With the only exception of the 8 $\\mu$m channel, the shape and amplitude of the power spectrum cannot be reproduced by either contributions from intervening dusty, Galactic neutral hydrogen gas (cirrus) or from local interplanetary dust (zodiacal light). Ordinary galaxies ($z \\lsim 5$) produce fluctuations due to their clustering and shot-noise. The faint flux limits ($\\geq 0.3$ $\\mu$Jy) of Spitzer data allow to push their residual clustering contribution below the level of the excess signal at relatively large ($\\gsim 50$~arcsec) angular scales (Kashlinsky et al. 2005). The shot noise component, estimated directly from galaxy counts, fits the observed fluctuations at smaller angular scales, and rapidly fades away at larger angles. The residual large scale signal has been ascribed by Kashlinsky et al. (2005) as coming from very distant ($z\\ge 5$) sources provided their total flux contribution is $> 1$ nW m$^{-2}$ sr$^{-1}$. The aim of this Letter is to show that this is indeed the case.  The layout of the paper is as follows: in Section 2 we will briefly describe the adopted model, while in Section 3 we provide predictions for the NIRB intensity and fluctuations and compare the latter ones with the results of  Kashlinsky et al. (2005). Section 4 summarizes our conclusions. ",
        "conclusions": "Using a physically-motivated, observationally-tested model of the early Universe (Schneider et al. 2005), we compute the expected background radiation in the NIR by sources forming when the Universe was $<1$ Gyr old. We find that the background intensity, $\\nu J_\\nu \\sim 1-2$ nW m$^{-2}$ sr$^{-1}$, is almost constant in the 0.8-8 $\\mu$m range. PopII galaxies dominate the NIRB in the entire wavelength range, while PopIII galaxies contribute at most 40\\% of the total intensity (at $\\lambda\\sim 1.5\\;\\mu$m), via their strong Ly$\\alpha$ line emission. Finally, we found that the infrared fluctuation excess on angular scales $\\geq 50$ arcsec detected by Spitzer/IRAC (Kashlinsky et al. 2005) is accounted very precisely by the clustering signal of galaxies at $z\\ge 5$ predominantly hosting stars with masses and properties similar to the present ones.  Two additional points are worth noticing: (i) very massive stars ($M \\geq 100 \\msun$) do not need to be invoked to explain NIRB fluctuations and reionization history; (ii) because of their small contribution ($P(q) \\leq 10^{-10}$ nW$^2$ m$^{-4}$ sr$^{-1}$) to the observed power spectrum in all channels, extracting the signal of the (very) first PopIII stars is extremely challenging. Future instruments (as the James Webb Space Telescope) will be able to directly identify these sources up to $z=10$ or above. Finally, the intensity of the NIRB provided by $z\\ge 5$ galaxies falls short of accounting for the excess measured by IRTS (Matsumoto et al. 2005) and DIRBE (Hauser \\& Dwek 2001) experiments. The origin of this component remains very puzzling (Salvaterra \\& Ferrara 2005) and might require either a revision of current model of zodiacal light subtraction or the existence of a large population of faint galaxies located at $z\\simeq 2-3$ (or both). Important insights on these issues are expected from the upcoming CIBER experiment (Bock et al. 2005), that will be able to simultaneously measure the total NIRB intensity and fluctuation power spectrum in the poorly known wavelength range 0.8-2 $\\mu$m. Such instrument, in addition, will allow a clear separation of the cosmological signal from local foregrounds (i.e. zodiacal light)."
    },
    "0512/astro-ph0512259_arXiv.txt": {
        "abstract": "We investigate the dissociation equilibrium of $\\rm H_2$ in very cool, helium-rich white dwarf atmospheres. We present the solution of the non-ideal chemical equilibrium for the dissociation of molecular hydrogen in a medium of dense helium. We find that at the photosphere of cool white dwarfs of $T_{\\rm eff}\\rm=4000 \\, K$, the non-ideality results in an increase of the mole fraction of molecular hydrogen by up to a factor of $\\sim 10$, compared to the equilibrium value for the ideal gas. This increases the $\\rm H_{2}-He$ CIA opacity by an order of magnitude and will affect the determination of the abundance of hydrogen in very cool, helium-rich white dwarfs. ",
        "introduction": "Several very cool white dwarfs with suspected $T_{\\rm eff}\\rm \\wig< 4500 \\, K$ have been discovered recently \\citep{FAR,KM,G,OPP,Harris01,HOD,Ibata00,HR2}. Most of them are thought to posses helium-rich atmospheres with an very high $\\rm He/ H \\wig> 10^{3}$ ratio \\citep{BAL,KM,G,BL,Bergeron01,OPP,HOD}. In most cases, however, current atmosphere models fail to reproduce the observed spectra and photometry of these peculiar stars.  The reason, and there may be more than one, for this shortcoming of the models is presently unknown. However, current models predict extreme physical atmospheric conditions for such stars, reaching densities of up to $\\rm 2-3 \\ g/cm^{3}$.  Under these conditions, the mostly ideal gas constitutive physics used in published atmosphere models is demonstrably inadequate. A careful look at the dense matter effects on the equation of state, chemistry, opacities, and radiative transfer is necessary to compute physically realistic models of these stars. Several of these effects have been studied previously, such as refractive radiative transfer \\citep{KS04}, the effects of fluid correlations on $\\rm He^-$ free-free and He Rayleigh scattering \\citep{KSM05,IRS}, and the ionization of warm, dense helium \\citep{KSM05,BSW}. In this contribution we present an additional correction that arises in the dense fluid: The solution for the dissociation of molecular hydrogen in dense fluid helium, in the limit $\\rm He/H>>1$. The relative importance of these corrections varies considerably, even more so when they are combined. As several more dense matter effects remain unexplored, it is premature to ponder their implications for the analysis of the coolest white dwarfs known and whether they will result in models that reproduce the data.  Nonetheless, incorporating adequate constitutive physics in atmosphere models is a necessary step to reach a proper understanding of these peculiar stars.  We introduce non-ideal effects into the equilibrium dissociation of molecular hydrogen through a modification of the chemical potentials of $\\rm  H \\, \\textrm{\\scriptsize{I}}$ and $\\rm H_2$ (section 2). We find that the strong interactions in the dense, helium-rich atmosphere results in a significant decrease in the dissociation fraction of molecular hydrogen, with a corresponding change in the $\\rm H_{2}-He$ Collision-Induced Absorption (CIA) opacity, which is linear function of $n_{\\rm H_{2}}$. In section 3, we illustrate the impact of the interactions on the $\\rm H_{2}/H \\,\\textrm{\\scriptsize{I}}$ ratio on a sequence of white dwarf atmosphere models with $T_{\\rm eff}\\rm=4000 \\, K$, a gravity of $\\rm log \\ \\it g \\rm = 8$ (cgs), and a homogeneous composition of ${\\rm He/H} =10^{2},10^{4}, \\ \\rm and \\ 10^{6}$, where He/H is the number abundance ratio. ",
        "conclusions": "Recent discoveries of a number of cool white dwarfs with peculiar spectral energy distribution represent a challenge in the modeling of very cool white dwarf atmospheres, as the spectra of these stars cannot be fitted with existing models. We believe that the main reason of this shortcoming of the models is the poorly explored, extreme physical regime found inside these atmospheres. We presented a correction to the abundance of $\\rm H_2$ that arises from the strong interaction of hydrogen molecules and atoms in a dense, fluid helium medium. We have found that in the dense, helium-rich, cool white dwarf atmospheres the formation of $\\rm H_2$ is more favorable than in the ideal gas description. For white dwarfs of $T_{\\rm eff}\\rm=4000 \\, K$ and $\\rm He/H\\wig>10^{3}$ the abundance of molecular hydrogen increases by an order of magnitude, with a corresponding increase in the $\\rm H_2-He$ CIA opacity by the same factor. This improvement is a new, significant effect that must be included in realistic modeling of very cool white dwarf atmospheres.  We have shown that the non-ideal effects affect strongly the abundances of even trace species in the atmosphere. This strongly suggest the necessity of revising the abundances of other trace species with significant opacity for similar effects. We expect that the study of their abundances and absorption processes in fluid helium will significantly improve our understanding of the atmospheric physics, composition, and evolution of the oldest and coolest white dwarfs.  I thank D. Saumon for useful discussions and the referee, P. Bergeron, for suggestions that improved the clarity of the manuscript. This research was supported by the United States Department of Energy under contract W-7405-ENG-36."
    },
    "0512/astro-ph0512545_arXiv.txt": {
        "abstract": "I comment (in a review fashion) on a few selected topics in the field of extragalactic globular clusters with strong emphasis on recent work.   The topics are: bimodality in the colour distribution of cluster systems, young massive clusters, and the brightest old clusters. Globular cluster research, perhaps more than ever, has  lead to important (at least to astronomers) progress and problems in galaxy structure and formation. ",
        "introduction": "\\begin{figure}[h] \\begin{center} \\includegraphics[width=10cm,angle=-00]{einstein.eps} \\end{center} \\caption{The heading of Einstein's paper on M13. The title reads: A simple application of the Newtonian law of gravitation to globular star clusters. } \\label{einstein} \\end{figure}   Since 2005 is the Einstein-year, a talk  on globular clusters can honour it by citing Einstein as a pioneer in globular cluster dynamics. In  his paper on M13 (Einstein 1921) he concluded that the non-luminous mass contributes no higher order of magnitude to the total mass than does the luminous mass (Fig.~\\ref{einstein}). To my knowledge this has been Einstein's only contact with globular clusters. As in other issues, his claim still holds.  \\\\ Globular clusters (GCs) may have the reputation of looking quite similar to each other, but they are a very inhomogeneous species regarding their intrinsic properties. In terms of mass, metallicity, age, density, they span orders of magnitudes. They are found in all types of galaxies, provided that the host galaxy's mass is sufficient. Different from what was thought  decades ago, GCs are not only survivors from the early Universe, but have also been formed regularly in galaxies and still form today.\\\\ \\indent The objective of this contribution is to highlight several recent topics in extragalactic GC research with some bias towards our own work. It is by no means exhaustive and will bravely face the usual fate, namely to be outdated soon. A more complete review is in preparation (Brodie \\& Strader 2006). The field is very extensive, ranging from stellar populations to cosmology, so only a few points can be illuminated. In the following I review/comment on these topics: colour bimodality in globular cluster systems (GCSs), young massive clusters, and the brightest old clusters. ",
        "conclusions": "The field of GC research, of which a few topics have been presented here, is as exciting as ever. One cannot yet claim that the formation of globular clusters and their relation to their host galaxies are well understood,  neither does it constitute a total mystery. Globular clusters may even have formed in different ways as entire objects or by coagulation of smaller units. Their potential for understanding galaxy formation and evolution is just on the verge of being exploited. The determination of ages, metallicites, internal dynamics, and kinematic and dynamical properties of cluster systems requires lots of work. However, it is fun by itself and promises deep insights into the evolution of structure. Einstein's notion that M13 does not host large amounts of dark matter touched already upon one of the fundamentals of structure formation and today even appears more revelant than 85 years ago."
    },
    "0512/gr-qc0512028_arXiv.txt": {
        "abstract": "The description of extreme-mass-ratio binary systems in the inspiral phase is a challenging problem in gravitational wave physics with significant relevance for the space interferometer LISA.   The main difficulty lies in the evaluation of the effects of the small body's gravitational field on itself.  To that end, an accurate computation of the perturbations produced by the small body with respect the background geometry of the large object, a massive black hole, is required. In this paper we present a new computational approach based on Finite Element Methods to solve the master equations describing perturbations of non-rotating black holes due to an orbiting point-like object.   The numerical computations are carried out in the time domain by using evolution algorithms for wave-type equations.   We show the accuracy of the method by comparing our calculations with previous results in the literature.  Finally, we discuss the relevance of this method for achieving accurate descriptions of extreme-mass-ratio binaries. ",
        "introduction": "Extreme-Mass-Ratio Binaries (EMRBs) in the inspiral stage of their evolution are considered to be a primary source of gravitational radiation~\\cite{Barack:2003fp,Gair:2004ea} to be detected by the proposed laser interferometric space antenna LISA~\\cite{Vitale:2002qv,Danzmann:2003ad,Danzmann:2003tv,Prince:2003aa}. They consist of a ``small'' object, such a main sequence star, a stellar mass black hole, or a neutron star, with mass $m$ ranging from $1 M_\\odot$ to $10^2 M_\\odot$, orbiting a massive black hole (MBH) with mass $M$ ranging from $10^3 M_\\odot$ (if we consider the case of intermediate mass black holes in globular clusters) to $10^9 M_\\odot$ (the case of big supermassive black holes sitting in the center of galaxies). This translates to EMRBs with mass ratios, $\\mu = m/M\\,,$ in the range $10^{-1} - 10^{-9}\\,$. In order to exploit this type of systems through LISA, it is crucial to have a good theoretical understanding of their evolution, good enough to produce accurate waveform templates in support of data analysis efforts. Because there is no significant coupling between the strong curvature effects produced by the MBH and its companion, relativistic perturbation theory is a well suited tool to study EMBRs.   Clearly, the accuracy of this approximation depends on the smallness of the mass ratio $\\mu$.  The challenge in modeling EMRBs is to compute the perturbations generated by the small body in the (background) gravitational field of the MBH, and how these perturbations affect the motion of the small body itself.   This problem has been known in literature as the {\\em radiation reaction} problem.  This is an old problem and several approaches to deal with it have been proposed (see the recent review by Poisson~\\cite{Poisson:2004lr} and the contributions to~\\cite{Lousto:2005cq}). The most extended approach consists in modelling the small object by using a point-like description and then, to describe the {\\em radiation reaction} effects on the dynamics as the action of a local {\\em self-force} that is responsible for the deviations from geodesic motion.  A consistent derivation of the equations of motion coming out from this set up was given by Mino, Sasaki and Tanaka~\\cite{Mino:1997nk}, and later, adopting an axiomatic approach, by Quinn and Wald~\\cite{Quinn:1997am} (see also~\\cite{Detweiler:2002mi}). However, these works only provide a formal prescription for the description of the orbital motion.  For the practical calculations of the self-force some techniques have been proposed: the {\\em mode-sum} scheme~\\cite{Barack:1999wf,Barack:2000eh,Barack:2001bw,Barack:2001gx}, and a regularization scheme based on zeta-function regularization techniques~\\cite{Lousto:1999za} (see~\\cite{Lousto:2005co} for a recent progress report).   The computation of the self-force and waveforms, and any other physical relevant information related to the inspiral due to radiation reaction constitute the main challenge of this problem.  One possible way is to resort to analytic techniques by adding extra approximations to the problem, similar to those from post-Newtonian methods.   However, the results may not be applicable to situations of physical relevance involving highly spinning MBHs and very eccentric orbits.  To make computations without making further simplifications of the problem, numerical techniques appear to be a necessary tool. It is important to distinguish between frequency-domain and time-domain calculations. The frequency domain approach has been used for a long time; it provides accurate results for the computation of quasinormal modes and frequencies~\\cite{Vishveshwara:1970vi,Chandrasekhar:1975qn}. However, the frequency-domain approach has more difficulties when we are interested in computing the waves originated from highly eccentric orbits since one has to sum over a large number of modes to obtain a good accuracy.  In this sense, calculations in the time-domain can be more efficient for obtaining accurate waveforms for the physical situations of relevance.   However, the time-domain numerical approach has to face a challenge, which consists of dealing with the different physical scales (both spatial and temporal) present in the problem and that expand over several orders of magnitude. Specifically, one needs to handle not only large wavelength scales comparable to the massive black hole, but also to resolve the scales in the vicinity of the small object where radiation reaction effects play a crucial role.    The conclusion is that we need to incorporate adaptive schemes in our numerical algorithms in order to provide the resolution that every region in the physical domain requires.  Since the small object is going to be moving through the domain (unless we choose a very particular coordinate system), it is convenient to allow the adaptive scheme to change in time to distribute properly the resolution.  Our choice to deal with these issues is the Finite Element Method (FEM), which is a numerical technique where adaptivity can be implemented in a natural way.   The FEM has other properties, which we will discuss in this paper, that make it very suitable to be used for the description of EMRBs and also for other physical systems that are the subject of investigation in Numerical Relativity.   In a recent work~\\cite{Sopuerta:2005rd}, we have already tested Adaptive Mesh Refinement techniques intrinsic to the FEM in a toy model consisting of a particle orbiting a black hole in the context of scalar gravity, where we have shown how an adaptive scheme can provide better accuracy than a non-adaptive scheme with an equivalent computational cost.   In this paper we use the FEM to perform time-domain simulations of a point-like object orbiting in geodesics (no radiation-reaction) around a non-rotating MBH, and compute physically relevant quantities like energy and angular momentum emitted in gravitational waves and waveforms.  This type of calculations constitute a good touchstone to evaluate the Finite Element (FE) techniques that we present in this paper, specially in relation to use this type of computations for the evaluation of the self-force on the particle. Since the MBH is non-rotating MBH the problem can be reduced to solve the one-dimensional Partial Differential Equations (PDEs) of black hole perturbations theory.  These equations, in the Regge-Wheeler gauge, reduce to a master equation from which the metric perturbations can be fully recovered.   The master equation for axial modes is known as the Regge-Wheeler equation, and for polar modes as the Zerilli-Moncrief equation. In this paper, instead of using the Regge-Wheeler function, we use a modification originally proposed by Cunningham, Price and Moncrief~\\cite{Cunningham:1978cp} that puts the axial modes on an equal footing with polar modes, as described by the Zerilli-Moncrief function, in relation to computing energy and angular momentum luminosities, and waveforms.    The plan of the paper is the following: In section~\\ref{bhpert} we summarize the main results from (non-rotating) black-hole perturbation theory that we need in our computations, including the explicit form of the source terms coming from the particle energy-momentum tensor. As far as we know, the expressions we present here for the sources associated with the axial modes, described by the Cunningham-Price-Moncrief master function, are new. We also perform an analysis of the discontinuities in the master equations due to the Dirac delta distributions that the source terms exhibit. In section~\\ref{femformulation} we describe the numerical framework.  We use a FE discretization for the spatial domain and a Finite Difference discretization in time. We start with the discretization of the domain, consisting of dividing the computational domain into disjoint subdomains (the {\\em elements}). Then, we describe the FE functional spaces, which are finite-dimensional functional spaces used to approximate locally (at each element) our solution.  The next step is the derivation of the {\\em weak} form of the master equations, which is an integral form.  It is important to remark that FE algorithms are derived from the integral form of the equations, in contrast with other numerical techniques where the differential form is used to obtain a discretization. From the weak form, we obtain the spatial discretization by imposing the vanishing of the residuals of our equations, which basically means to impose the vanishing of the components of the equations with respect to a basis of functions constructed from the FE functional spaces.  This process leads to a coupled system of Ordinary Differential Equations (ODEs) which has a close analogy with the equations governing the behaviour of a system of coupled oscillators.  A very important point in the discretization process is the fact that, because the FE formulation is based on an integral form of the equations, we obtain automatically a discretization of the sources containing Dirac delta distributions and its first derivative without having to resort to sequences of functions approaching in some limit the Dirac delta, we just use the properties of these distributions at the analytic level in the {\\em weak} form of the equations.  To solve the resulting ODEs we introduce a collection of evolution algorithms to solve the equations in second-order form and which have parameters that allow us to control the appearance of spurious high-frequency modes, which are common in systems like the one we are studying having a very localized source.  We finish this section by discussing the structure of the mesh, in particular how adaptivity is implemented and how we can change in time this structure as the particle moves.   In section~\\ref{results} we discuss the performance of the FE numerical code we have developed and compare results regarding the computations of energy and angular momentum radiated with previous works in the literature, showing in this way the accuracy that this method is able to achieve.  We conclude in section~\\ref{discussion}, where we discuss the convenience of using the FEM for the simulations of EMRBs in the light of the results of this paper and describe possible ways to proceed in the future to make this goal a reality.  Finally, we have included two appendices:  In Appendix~\\ref{geodesic}, we summarize the geodesic equations of motion for the particle, and in Appendix~\\ref{GaussLegendre}, we briefly describe the Gauss-Legendre quadrature method for evaluating numerically some of the integrals that appear in the FE discretization of our equations.   The conventions that we follow throughout this work are: Greek letters are used to denote spacetime indices; capital Latin letters are used for indices in the time-radial part of the metric;  lower-case Latin indices are used for the spherical sector of the metric.  We use physical units in which $G = c = 1$. ",
        "conclusions": "\\label{discussion} In this paper we have presented a new method for computing the gravitational radiation emitted by a point like object orbiting a non-rotating black hole.  We have shown that the method is accurate by comparing it with previous results in the literature obtaining an agreement with relative errors of the order of $1$\\%, in many cases even of the order of $0.1$\\% or below.  We also have shown that these numerical techniques provide sufficiently smooth waveforms, which is one of the goals of these calculation in relation with gravitational-wave data analysis efforts. These results together with the particular feature of the computational method presented suggest that it is a suitable method to be use in self-force calculations for inspiralling EMRBs and in the posterior waveform calculations at the next perturbative order. Our numerical calculations are based on the FEM and related techniques. The main features of the FEM that makes it suitable for the study of EMRBs, and perhaps also for problems that Numerical Relativity deals with, are the following: (i) {\\em Proper description of the Computational Domain}.  This is particularly relevant when we want to solve the perturbative equations in a 2D or 3D setup (see~\\cite{Sopuerta:2005rd}), as it is the case if we want to consider a rotating Kerr black hole, which is the astrophysically relevant case.  It would be also relevant for the study of black hole spacetimes in Numerical Relativity.  In this scenario, the spacetime geometry may involve holes (inner boundaries arising from black hole singularity excision) and we may wish to use a spherical-type outer boundary to allow gravitational radiation leave the domain smoothly.    All these geometric issues have usually caused a number of problems in Finite Differences techniques, but can be handled in a natural way using the FEM.   In this respect, the FEM has already shown its capabilities in solving problems in other scientific areas that involve much more complicated domains than the ones we can face in General Relativity. (ii) {\\em Imposition of Boundary Conditions}. In close connection with the previous point, the underlaying philosophy in the FEM is that one should use the mesh that adapts best to the geometric characteristics of the problem we want to solve.  In  particular, to the boundary conditions, since it is not equally simple and convenient to impose outgoing radiation conditions in a rectangular boundary than in a spherical one.  This also has an impact when we perform the FEM discretization, since it is based on the {\\em weak} form of our equations, which can have built in the the boundary conditions.  In the case of problems in 2D or higher dimensions, if the boundary is {\\em natural} (adapted to the problem), the implementation of the boundary conditions becomes trivial (see, e.g.~\\cite{Sopuerta:2005rd}). This has advantages even in 1D problems, like the one we have studied in this paper, where the imposition of boundary conditions like Sommerfeld or von Neumann is simpler than in a Finite Differences framework.  This paper illustrates this fact. (iii) {\\em Treatment of distributions.}  Many description of EMRBs treat the small body as a point-like object, which despite being somehow unnatural in General Relativity, allows us to perform computations in a consistent way.   The consequence of having a point-like object is that the equations that we have to solve contain source terms where Dirac delta distributions and its derivatives (up to second derivatives in the case we were solving the Teukolsky equation sourced by a point-like object) appear.   To deal with this kind of distributions in a Finite Difference framework is not an easy task, and the different ways in which one can handle them involved not trivial {\\em a priori} regularizations of the distributions. Instead, in the FEM, the fact that the discretization is based on the {\\em weak} form of the equations, an integral formulation, is a key point.  We can evaluate the integrals that involve Dirac distributions analytically by using the properties of the distributions, without the need of using any regularization of those distributions.  A sample of this has been given in this paper, where we used the {\\em weak} formulation of the problem to discretize a source term containing the Dirac delta distribution and its first derivative.   Then, the type of discretization we would get is in this sense analogous to the one proposed by Price and Lousto~\\cite{Lousto:1997wf,Martel:2001yf,Martel:2003jj}, where they also used an integral form of the equations to discretize them. Therefore, using the FEM provides an additional advantage for the study of EMRBs where the small object is treated as a point-like object. (iv) {\\em Adaptivity.}   This is a key ingredient for the simulations of EMRBs.  The calculations presented in this paper do not necessarily require adaptivity, but they are an excellent benchmark to test these techniques.   However, for the case of rotating massive black hole adaptivity may be the only way of performing physically realistic simulations.  The FEM is a natural choice to achieve the high level of adaptivity required, both in the construction of the mesh and later by using any of the robust techniques of mesh refinement available (see~\\cite{Sopuerta:2005rd} and references therein).   Apart from these specific reasons, there are other motivations in favour of using the FEM.   In this sense it is important to mention that because the FEM is based on piecewise (polynomial) approximations, it lies on sound mathematical grounds (see, e.g.~\\cite{StrFix:73,Ciarlet:1978bo}). From the point of view of building numerical codes based on the FEM, it is important to emphasize the high degree of independence of the different ingredients of the FEM discretization process~\\cite{StrFix:73,Hughes:1987tj,Babuska:2001bs,Zienkiewicz:77oc}, which makes it very suitable for modular programming.  In addition, the FEM has been widely used in many areas of scientific research and, as a consequence, a number of FEM packages and tools are available for scientific computation.   There are a number of ways of extending this work in order to improve the computational framework in order to simulate EMRBs, and in particular to evalute the radiation-reaction effects.  From the computational side we can introduce higher-order elements (by using FE functional spaces with higher-order polynomials), which will improve the accuracy of the computations. From the physical point of view, we can change the description of the gravitational field, meaning the formulation of the perturbative scheme.  In this sense, to compute the metric perturbations using the Lorentz gauge, as it has been recently proposed by Barack and Lousto~\\cite{Barack:2005nr}, appear to be a very convenient choice for a number of reasons (see~\\cite{Barack:2005nr} for a detailed discussion).  Among the advantages of this approach it is worth to mention the following ones: (i)  Because one is working with pure metric perturbations the sources do not contain derivatives of the Dirac delta distribution, and hence the solution of the equations is continuous at the particle location, which will improve the accuracy of the computations. (ii) Moreover, in contrast with the computations in the Regge-Wheeler gauge, we do not need a metric perturbation reconstruction procedure (just algebraic computations) to evaluate the self-force. (iii) The regularization procedures to obtain the self-force have only been given in the Lorentz gauge.   It also has some disavantages:  We need to solve a coupled system of equations instead of single wave-type equations, and there are constraints that need to be satisfied along the evolution.   In the astrophysically motived EMRBs, the MBH is highly rotating and therefore it is desirable to be able to repeat these calculations by using the Kerr solution as the background spacetime.  This is a more difficult problem since it involves three-dimensional PDEs (or two-dimensional if we factor out the dependence in the polar angle). In this sense, it is important to mention that the FEM techniques that have been presented and used in this paper can be transfered to the higher-dimensional problem of computing Kerr perturbations.  For the same reasons that have been pointed out before, a promising approach may be to solve for metric perturbations of the Kerr black hole in the Lorentz gauge.        \\[ \\] {\\bf Acknowledgements:} The authors acknowledge the support of the Center for Gravitational Wave Physics funded by the National Science Foundation under Cooperative Agreement PHY-0114375, and support from NSF grant PHY-0244788 to Penn State University. They also wish to thank the Information Technology Services at Penn State University for the use of the LION-XO computer cluster in some of the calculations presented in this paper, and to Uli Sperhake, Pengtao Sun, and Jinchao Xu for fruitful discussions.   \\appendix"
    },
    "0512/astro-ph0512015_arXiv.txt": {
        "abstract": "Analysis of INTEGRAL Core Program and public Open Time observations performed up to April 2005 provides a sample of 62 active galactic nuclei in the 20-100 keV band above a flux limit of $\\sim$ 1.5$\\times$10$^{-11}$ erg cm$^{-2}$ s$^{-1}$. Most(42) of the sources in the sample are Seyfert galaxies, almost equally divided between type 1 and 2 objects, 6 are blazars and 14 are still unclassified.  Excluding the blazars, the average redshift of our sample is 0.021 while the mean luminosity is Log(L) = 43.45. We find that absorption is present in 65\\% of the objects with 14\\% of the total sample due to Compton thick active galaxies. In agreement with both Swift/BAT team results and 2-10 keV studies, the fraction of absorbed objects decreases with the 20-100 keV luminosity. All Seyfert 2s in our sample are absorbed as are 33\\% of Seyfert 1s. The present data highlight the capability of INTEGRAL to probe the extragalactic gamma-ray sky and to find new and/or absorbed active galaxies. ",
        "introduction": "The extragalactic gamma-ray sky is still poorly explored with only a few surveys having been performed so far above 10 keV: the all-sky survey conducted in the 1980's by the HEAO/A4 instrument in the 13-80 keV band (Levine et al. 1984) and more recently those of RXTE/PCA and Swift/BAT, in the 8-20 and 14-195 keV bands respectively (Revnivtsev et al. 2004, Markwardt et al. 2005). The latter, characterized by a positional uncertainty of $\\leq$3\\arcmin and a flux detection limit of $\\sim$10$^{-11}$ erg cm$^{-2}$ s$^{-1}$, is the most accurate and sensitive survey to date at high energies. It covers the high latitude sky providing a sample of 44 active galactic nuclei (AGN), most of which are previously known from X-ray studies.  Despite being so rare, gamma-ray surveys are an efficient way to find  AGN as they probe heavily obscured regions/objects, i.e. those that could be missed in optical, UV, and even X-ray surveys. Indeed, 64\\% of the non-blazar sources found by Swift/BAT have absorption in excess of 10$^{22}$ atoms cm$^{-2}$ and the overall column density distribution is bimodal. While none of the sources brighter than 3$\\times$10$^{43}$ erg s$^{-1}$ shows high column densities, almost all weaker objects are absorbed. Quantifying the fraction of AGN missed by low energy surveys is necessary if we want to provide input parameters for synthesis models of the X-ray background and to understand the accretion history of the Universe.  A further step in this field is provided by the imager (IBIS) on board INTEGRAL, which, like BAT, is surveying a large fraction of the sky above 20 keV with similar sensitivity and positional accuracy. Analysis of the first year of INTEGRAL observations covering largely the galactic plane and centre has provided a first sample of 10 gamma-ray selected AGN (Bassani et al. 2004); more recently the second IBIS survey has listed 32 such objects (Bird et al. 2005). Detailed analysis of INTEGRAL observations also suggests that AGN are becoming a major constituent of the IBIS source population (Revnivtsev et al. 2005, Beckmann et al. 2005); at least 15\\% of the objects in the second IBIS catalogue are active galaxies. Here, we present a further step in our all-sky survey project, limiting the search to extragalactic objects. We have analysed $\\sim$11300 INTEGRAL pointings and detected 62 AGN in the 20-100 keV energy range. ",
        "conclusions": "It must be remembered that the present survey is highly inhomogeneous in sky exposure and coverage, thus the present sample is far from being  complete. Nevertheless a few interesting  considerations can be drawn by a simple statistical analysis. For objects with known distance, we plot in figure 1 the gamma-ray luminosity against redshift, to show the large range in these parameters sampled by the present survey. From this figure it is also evident that our sensitivity limit is around 1.5$\\times$10$^{-11}$ erg cm$^{-2}$s$^{-1}$ (straight line in the figure).  Within the overall sample, 42 objects are classified as Seyfert galaxies, 6 are blazars and 14 are still unclassified. Within the sample of Seyfert galaxies, 23 objects are of type 1-1.5 while 19 are of type 2 thus illustrating the power of gamma-ray surveys to find narrow line AGN. Excluding the blazars, the average redshift of our sample is 0.021$\\pm$0.017 (1$\\sigma$) while the mean luminosity in Log is 43.45$\\pm$0.71. Assuming 10$^{22}$ atoms cm$^{-2}$ to be the dividing line between absorbed and unabsorbed sources, we find that absorption is present in 65$\\pm$17\\% of the sample. This result is in line with the Swift findings (Markwardt et al. 2005) and above that found in X-ray surveys for bright objects (La Franca et al. 2005, Comastri 2004). It is also interesting to note that 14$\\pm$7\\% of the total sample is due to Compton thick objects which is about three times the frequency found by Swift.  In figure 2 we show the column density of the sources as a function of IBIS luminosity. Although from this figure there is little evidence of any strong correlation, the few high luminosity objects in the survey all have low absorption. We can further investigate the relationship between N$_{H}$ and Log(L)  by forming the ratio between the fraction of absorbed sources above a given luminosity and that below this value. It can be seen from figure 3 that below log(L)$\\sim$43.5 the ratio is practically constant while above this luminosity it decreases sharply implying that at high IBIS luminosities  there are very few absorbed sources. This is in agreement with that found by the Swift/BAT team and obtained from 2-10 keV studies (La Franca et al. 2005).  Within the sub-sample of 17 Seyfert 2s with known N$_{H}$, we find that all are absorbed and almost 30\\% are Compton thick; this is in line with previous estimates of the column density distribution of type 2 objects based on X-ray data (Risaliti et al. 1999, Bassani et al. 1999). Interestingly we also find that 33\\% of type 1 objects are absorbed.  More in depth studies of the present sample require optical classification of all objects  and a detailed analysis of their broad band behaviour (particularly in the X-ray to gamma-ray band) in order to understand the role of absorption and the relation of extragalactic gamma-ray surveys to those in other wavebands. In the meantime this catalogue testifies to the power of INTEGRAL/IBIS to probe the extragalactic gamma-ray sky, discover new active galaxies and find absorbed objects."
    },
    "0512/astro-ph0512365_arXiv.txt": {
        "abstract": "We present 3D hydrodynamical simulations of ram pressure stripping of massive disc galaxies in clusters. Studies of galaxies that move face-on have predicted that in such a geometry the galaxy can lose a substantial amount of its interstellar medium. But only a small fraction of galaxies is moving face-on. Therefore, in this work we focus on a systematic study of the effect of the inclination angle between the direction of motion and the galaxy's rotation axis.  In agreement with some previous works, we find that the inclination angle does not play a major role for the mass loss as long as the galaxy is not moving close to edge-on (inclination angle $\\lesssim 60\\degree$). We can predict this behaviour by extending Gunn~\\& Gott's estimate of the stripping radius, which is valid for face-on geometries, to moderate inclinations.  The inclination plays a role as long as the ram pressure is comparable to pressures in the galactic plane, which can span two orders of magnitude. For very strong ram pressures, the disc will be stripped completely, and for very weak ram pressures, mass loss is negligible independent of inclination. We show that in non-edge-on geometries the stripping proceeds remarkably similar. A major difference between different inclinations is the degree of asymmetry introduced in the remaining gas disc.  We demonstrate that the tail of gas stripped from the galaxy does not necessarily point in a direction opposite to the galaxy's direction of motion. Therefore, the observation of a galaxy's gas tail may be misleading about the galaxy's direction of motion. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512153_arXiv.txt": {
        "abstract": "Two evolutionary scenarios are proposed for the formation of extreme helium stars: a post-AGB star suffering from a late thermal pulse, or the merger of two white dwarfs. An identification of the evolutionary channel for individual objects~has~to rely on surface abundances. We present preliminary results from a non-LTE analysis of CNO, Mg and S for two unique objects, V\\,652\\,Her and HD\\,144941. Non-LTE abundance corrections for these elements range from negligible values to $\\sim$0.7\\,dex. Non-LTE effects typically lead to systematic shifts in the abundances relative to LTE and reduce the uncertainties. % ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512479_arXiv.txt": {
        "abstract": "We discuss the production of cosmogenic neutrinos on extragalactic infrared photons in a model of its cosmological evolution. The relative importance of these infrared photons as a target for proton interactions is significant, especially in the case of steep injection spectra of the ultrahigh energy cosmic rays. For an E$^{-2.5}$ cosmic ray injection spectrum, for example, the event rate of neutrinos of energy above 1 PeV is more than doubled. ",
        "introduction": "The assumption that the ultra high energy cosmic rays (UHECR) are nuclei (presumed here to be protons) accelerated in powerful extragalactic sources provides a natural connection between these particles and ultra high energy neutrinos. This was first realized by Berezinsky and Zatsepin~\\cite{BerZat69} soon after the introduction of the GZK effect~\\cite{GZK}. The GZK effect is the modification of the UHE proton spectrum from energy losses by photoproduction interactions with the 2.7K microwave background radiation (MBR). In the case of isotropic and homogeneous distribution of UHE cosmic ray sources, the GZK effect leads to a cut-off of the cosmic ray spectrum below 10$^{20}$ eV. The charged mesons generated in these interactions initiate a decay chain that results in neutrinos. Since the mesons and muons do not lose energy before decay, the high energy end of the spectrum of these neutrinos follows the injection spectrum of UHECR, while below the interaction threshold it is flat \\cite{Stecker73}, \\cite{Stecker79}. The neutrinos which are produced by photomeson producing interactions of UHECR nuclei are sometimes referred to as {\\em cosmogenic neutrinos}.  Several calculations of of the fluxes of UHE photomeson neutrinos were published in the 1970s \\cite{WTW,Stecker73,BerSmi75,BerZat77,Stecker79}, Hill and Schramm~\\cite{HS85,HSW86} used the non-detection of such neutrinos to place an upper limit on the cosmological evolution of the sources of UHECR.  The problem has been revisited several more times~\\cite{YT93,PJ96,ESS01}.  In 2004 Stanev~\\cite{TS04} considered interactions of UHECR with photons of the extragalactic infrared and optical background (IRB), pointing out that this process generates non-negligible cosmogenic neutrino fluxes. This suggestion was quickly followed by a confirmation in Ref.~\\cite{BMM04} which emphasized the importance of the IRB as interaction target. This idea was further developed in Ref.~\\cite{BK05}.  Ref.~\\cite{TS04} gave an estimate of the cosmogenic neutrino flux generated in interactions on the IRB, but did not account correctly for the cosmological evolution of the infrared background. In this paper we perform a calculation using a realistic empirically based  model of the cosmological evolution of the spectral energy distribution of the extragalactic IR-UV  background given in Ref. \\cite{SMS05} which will be referred to as SMS05. The aim is to estimate correctly the role of these extragalactic photons, particularly the infrared photons which are by far the most numerous, as targets for UHE proton interactions.  The paper is organized as follows: in Section II we discuss the model of the infrared background and its cosmological evolution. In Section III we describe the calculation. Section IV gives the results of the calculation and Section V contains the discussion of the results and the conclusions from this research. ",
        "conclusions": "The evolution models of SMS05 do not give the highest IRB-generated neutrino flux. We compared the IRB photon  density in this model with the models of Refs.~\\cite{F01,KMH02}. Both of these models have higher IRB density in the relevant energy range between 3$\\times$10$^{-3}$ and 1 eV. The total IRB densities in the this range are 1.27~\\cite{F01} and 1.12~\\cite{KMH02} compared with the density of 1.03 used here. In addition, Ref.~\\cite{KMH02} shows somewhat faster cosmological evolution. The use of any of these models would have increased somewhat the calculated flux of cosmogenic neutrinos. The uncertainty in the IRB flux is of the order of 30\\%~\\cite{SMS05}, while the biggest uncertainty in this calculation is in the UHECR flux and its cosmological evolution.  The IRB contribution to the total cosmogenic neutrinos flux can be slightly increased as protons of energy below 10$^{18}$ eV can interact with IRB photons and generate lower energy neutrinos. If such interactions were included the IRB spectrum would be wider than the MBR one, especially at energies below 10$^{16}$ eV.  It is difficult to compare our results with those of Refs.~\\cite{BMM04,BK05} because of the different astrophysical input in these calculations. Qualitatively the results of these calculations are similar to ours and certainly agree within a factor of 2.  In conclusion, we calculated the flux of cosmogenic neutrinos from interactions of UHECR protons with IRB photons using the recent calculations of IR photon spectra densities as a function of redshift by SMS05. Our calculations show that UHECR interactions with IRB photons produce a significant flux of cosmogenic neutrinos, one which is comparable to interactions with MBR photons. This is especially true in the case of assumed steep injection spectra of the ultrahigh energy cosmic rays. The total neutrino event rates at energies above 1 PeV increase by more than a factor of 2 in the case of injection spectra with $\\gamma$ = 1.5. Because of the much lower mean free path of protons above 10$^{20}$ eV in the MBR interactions with IRB photons do not increase the higher energy end of the cosmogenic neutrino spectrum. The total cosmogenic fluxes, however, are still not detectable with conventional neutrino telescopes such as IceCube~\\cite{IceCube} or the European km$^3$ telescope~\\cite{km3}. A reliable detection is only expected from radio~\\cite{radio} and acoustic neutrino detectors.   {\\bf Acknowledgments} We thank Tanja Kneiske for sharing with us the tables of cosmological evolution of the IR background from Ref. \\cite{KMH02}. This work was supported in part by NASA grants ATP03-0000-0057 and  ATP03-0000-0080."
    },
    "0512/astro-ph0512246_arXiv.txt": {
        "abstract": "We have proposed a new High Resolution Shock Capturing (HRSC) scheme for Special Relativistic Hydrodynamics (SRHD) based on the semidiscrete central Godunov-type schemes and a modified Weighted Essentially Non-oscillatory (WENO) data reconstruction algorithm. This is the first application of the semidiscrete central schemes with high order WENO data reconstruction to the SRHD equations.  This method does not use a Riemann solver for flux computations and a number of one and two dimensional benchmark tests show that the algorithm is robust and comparable in accuracy to other SRHD codes. ",
        "introduction": "Gas flows at (ultra)relativistic speeds are an integral component of many astrophysical phenomena. Some significant examples of these include accretion disks around compact objects, Gamma ray bursts, collapse of (super)massive stars, pulsar wind nebula/supernova remnant interactions, Active Galactic Nuclei (AGN), X-ray binaries, superluminal jets and many others. As it stands, there is a massive amount of astronomical data from many of these sources that to date, have only been partially understood and which hold clues to many problems related to the phenomena mentioned above. Correlating these data with the theory in many of these phenomena would require doing simulations that would include modelling (ultra)relativistic gas/fluid flows.  Such flows are described by non-linear hyperbolic conservation laws that have shock waves as possible solutions. Shock waves are difficult to approximate numerically using standard finite difference techniques as they usually lead to spurious oscillations. High Resolution Shock Capturing (HRSC) schemes are a class of numerical methods devoted specifically to approximating hyperbolic conservation laws. Over the past fifteen years HRSC research have progressed immensely and they have been applied to a wide variety of problems involving classical and relativistic hydrodynamics.  For a pedagogical introduction to some modern HRSC schemes and their applications, see \\cite{leveque2,toro,pen}.  To date,  the most commonly used HRSC schemes in relativistic hydrodynamics have been ones that have not incorporated many of the recent advances in HRSC research.  In this work, we will apply a new scheme to the SRHD equations that is simpler than most of the previous approaches and incorporates a number of recent advances from HRSC research. In Rahman $\\&$ Moore (2005), this scheme was applied to the multidimensional classical hydrodynamics code.  The work presented here should be considered an extension of that work. Before describing the motivations behind the algorithm used in this work, we provide a brief review of numerical methods used for the SRHD equations and that of HRSC schemes.  Wilson J. R. (1972), Wilson et al. (1979), Hawley et al. (1984), and Centrella $\\&$ Wilson (1984) were the first to apply numerical methods to approximate the SRHD equations. They discretized the SRHD equations using an explicit finite difference scheme with artificial viscosity and monotonic transport.  Since their work,  this technique has been used to study a number of astrophysical phenomena including stellar collapses, accretion disks, cosmology etc.  Most SRHD schemes up to the late eighties used the artificial viscosity technique of Wilson et al. (1979) to handle shock waves. However,  a major breakthrough occurred when HRSC schemes were applied to the SRHD equations (\\cite{marti91,marquina92,eulderink93,eulderink95}). One of the most significant of these new schemes was the relativistic Piecewise Parabolic Method (PPM) by Mart\\'i $\\&$ M\\\"uller (1996) (MM) which was based on the PPM reconstruction scheme of Colella $\\&$ Woodward  (1985).  For a recent review of modern SRHD schemes, we refer the reader to Mart\\'i $\\&$ M\\\"uller (2003).  The use of HRSC schemes for SRHD equations essentially shifted the paradigm for solving these equations from artificial viscosity based methods to this approach.  Modern HRSC schemes are based on two conventional approaches for solving hyperbolic conservation laws.  These are the so-called upwind (\\cite{Harten,vanleer,roe,toro}) and central methods (\\cite{LxF,lax54,NT}). A common aspect of most modern HRSC-SRHD schemes is that they are solved using upwind schemes that use so-called Riemann solvers to approximate the solution.  A drawback of using Riemann solvers is that they are computationally expensive and difficult to implement because they need computations of eigenvalues and eigenvectors of the flux matrix and uses flux splittings etc., to compute the numerical fluxes. On the other hand, central schemes do not use Riemann solvers and are simpler and easier to implement. Until recently, central schemes were not considered for wide ranging applications because a trade-off for their simplicity was loss of accuracy. This is because central schemes were generally more dissipative than their upwind counterparts.  However, recent progress in HRSC research has addressed this issue and some newer formulations of central schemes have been shown to be comparable performance-wise to the upwind approach. The work presented here is based on some of these newer central schemes.  Before discussing them,  we provide below a brief overview of central schemes and discuss the most important aspect of HRSC schemes known as non-oscillatory data reconstructions methods.  Modern central schemes are based on the first and second order shock capturing algorithms developed by Lax-Friedrichs (LxF) (\\cite{LxF,lax54}), and Nessyahu $\\&$ Tadmore (1990) (NT), respectively. Since NT, there have been a number of extensions of central HRSC schemes to higher orders and multidimensions. Some of these recent extensions are summarized below. The NT scheme was extended to 3$^{rd}$ order by Liu $\\&$ Tadmor (1998), and to multidimensions by Jiang $\\&$ Tadmore (1998). These were followed by a new formulation of central HRSC schemes known as semidiscrete central schemes.  The semidiscrete schemes were designed to address the dissipation issue mentioned above and were proposed in Kurganov $\\&$ Tadmore (2000) (KT).  KT showed that semidiscrete schemes, which are formulated on non-staggered grids and use more accurate information of local speeds of propagation, are less dissipative than their staggered grid counterparts.  Hence, these schemes retained the simplicity of the central approach and were comparable in accuracy to other Riemann solver based upwind approaches.  The inception of semidiscrete central schemes precipitated a great deal research on extending this formalism which included higher order formulations in multidimensions,  genuinely multidimensional formulations and unstructured grid formulations among others(\\cite{LT,GT,LPR1,LPR2,KP1,KNP,KL}).  Besides the method used to advance the solution,  another aspect of HRSC schemes that is equally important for any algorithm is the so called non-oscillatory data interpolation method. Data interpolation is used in all HRSC schemes.  Non-oscillatory data reconstruction is the piecewise continuous polynomial interpolation of the data (which may contain discontinuities) over the computational domain.  Most of the HRSC schemes mentioned above use the so called Essentially Non-Oscillatory (ENO) (\\cite{HEOC}) data reconstruction algorithm for data interpolation. This methods works by interpolating the data using the smoothest stencil from a number of choices.  A modern extension of ENO schemes is the Weighted  Essentially Non-Oscillatory data reconstruction (WENO) algorithm (\\cite{LOC,JS}) that has a number of advantages over its predecessor.  In order to take advantage of the developments in central approaches and high order data reconstruction, another class of HRSC schemes were developed as an extension of the NT scheme by Levy et al. (\\cite{LPR1,LPR2}) called the Central Weighted Non-Oscillatory (CWENO) schemes. Recently, Kurganov $\\&$ Levy (2000) (KL) have combined the semidiscrete central schemes with the WENO reconstruction method and proposed yet another, better HRSC scheme.  The work presented here is based on the KL HRSC scheme which we will describe below.  Despite the many advances mentioned above, there have been relatively few attempts at applying central-type methods to computational astrophysics.  However this is rapidly changing.  Among the most significant attempts to apply central schemes for computational astrophysics, the following are noteworthy.  Del Zanna $\\&$ Bucciantini (2002), have developed an algorithm that is a variation of the conventional finite volume central type approach and applied it to multidimensional SRHD and SRMHD equations.  They have also used these codes to study Pulsar Wind Nebula (PWN) interactions with the interstellar medium \\cite{delzanna04} and supernova remnants \\cite{delzanna03}.  Anninos $\\&$ Fragile (2003,2004) have developed a new central type scheme for relativistic hydrodynamics in fixed, curved spacetime and have applied their scheme to study accretion disks around kerr black holes. Lucas-Serano et al. (2004) have assessed the applicability of central schemes to SRHD equations using the PPM data reconstruction scheme \\cite{WC2}.  Their results were the first to demonstrate that central type schemes of KT can be used for SRHD. Recently, Shibata $\\&$ Font (2005) have successfully tested the suitability of central type schemes for general relativistic simulations.  We turn our attention now to applications of the WENO data reconstruction methods in computational astrophysics. Balsara (\\cite{balsara}) have used the WENO methods extensively in magnetohydrodynamics.  Feng et al. (2004), have recently developed a cosmological hydrodynamics code using the WENO reconstruction scheme.  Zhang $\\&$ MacFadyen (2005), have recently used the WENO scheme in their adaptive SRHD code.  The work mentioned above point to a promising future for the application of central and WENO schemes in astrophysical hydrodynamics.  Indeed, the work presented here was done concurrently with many of the recent work mentioned above.  However, the algorithm presented here is the first to combine the WENO and the semidiscrete central approach for applications in SRHD.  We describe below this algorithm and its novel aspects.  In a previous paper, \\cite{rahman1}, we have described in detail the incentive behind our algorithm for multidimensional classical hydrodynamics.  Many of the same reasons apply for the SRHD equations as well.  The algorithm is based on combining the semidiscrete approach of Kurganov $\\&$ Levy (2000) with high order WENO data reconstruction methods.  To ensure robustness we have added the flattening, steepening and monotonicity preserving techniques of the PPM reconstruction scheme by Colella $\\&$ Woodward (1985) to the data reconstruction scheme. Essentially, the simplicity of central schemes, the accuracy WENO reconstructions and the robustness of the piecewise parabolic method have been combined to propose a new robust central scheme.  Building on the success of this algorithm for classical hydrodynamics,  it has been applied to the SRHD equations. This work is original in several respects.  It is the first attempt at solving the SRHD equations using the semidiscrete central scheme with WENO data reconstructions.  The algorithm has been tested using both 3$^{rd.}$ and 4$^{th.}$ order data reconstruction techniques. It is also a dimensionally unsplit algorithm and to our knowledge is the only unsplit multidimensional algorithm in SRHD. The work presented here can be considered as laying the foundation for the development of a multi-purpose, multi-dimensional SRHD code that could be used to study a wide variety of astrophysical phenomena.  The paper is organized as follows.  Sec. 2 presents the SRHD equations.  Sec. 3 describes the new algorithm for numerically approximating the SRHD equations.  In Sec. 4, we present the results of a number of one dimensional benchmark tests.  In Sec. 5, the two dimensional tests are presented.  Finally we conclude this paper in Sec. 6. ",
        "conclusions": "\\label{section:summary}  We have developed a new multidimensional relativistic hydrodynamics code based on the semidiscrete central and WENO reconstruction approach, and some elements of the PPM method.  This is the first such multidimensional relativistic code combining these elements. A dimensionally unsplit scheme was used to advance the solution.  We also carried out a number of one and two dimensional tests of the code using both 3$^{rd}$ and 4$^{th}$ order reconstruction methods.  The test results are comparable to a number of other codes currently being used in computational astrophysics.  In all cases,  the tests showed good agreements with codes currently in use.  The work presented here should be considered as the first few steps towards the development of a multi-purpose relativistic hydrodynamics code.  Some of the developments mentioned above on semidiscrete central schemes could well extend the performance of our code.  These include the genuinely multidimensional formulation,  the formulation on unstructured grid or higher order WENO reconstructions. From the point of applications,  we are working to extend our code to three dimensions, provide it with adaptive capabilities and apply it for parallel architecture using MPI."
    },
    "0512/astro-ph0512070_arXiv.txt": {
        "abstract": "The exact nature of weak \\MgII~absorbers (those with $W_r(2796) < 0.3~\\Ang$) is a matter of debate, but most are likely related to areas of local star formation or supernovae activity outside of giant galaxies.  Using 18 QSO spectra obtained with the Ultra-Violet Echelle Spectrograph (UVES) on the Very Large Telescope (VLT), we have conducted a survey for weak \\MgII~absorbers at $1.4 < z < 2.4$.  We searched a redshift path length of $\\Delta z = 8.51$, eliminating regions badly contaminated by atmospheric absorption so that the survey is close to 100\\% complete to $W_r(2796) = 0.02~\\Ang$.  We found a total of 9 weak absorbers, yielding a number density of absorbers of $dN/dz = 1.06 \\pm 0.12$ for $0.02 \\leq W_r(2796) < 0.3~\\Ang$.  \\citet{Nar05} found $dN/dz = 1.00 \\pm 0.20$ at $0 < z < 0.3$ and \\citet{Church99} found $dN/dz = 1.74 \\pm 0.10$ at $0.4 < z < 1.4$.  Therefore, the population of weak \\MgII~absorbers appears to peak at $z \\sim 1$.  We explore the expected evolution of the absorber population subject to a changing extragalactic background radiation (EBR) from $z = 0.9$ to $z = 1.78$ (the median redshift of our survey), and find that the result is higher than the observed value.  We point out that the peak epoch for weak \\MgII\\ absorption at $z\\sim1$ may coincide with the peak epoch of global star formation in the dwarf galaxy environment. ",
        "introduction": "\\label{sec:intro}  Quasar absorption line systems (QALS) are unique and powerful tools for studying the chemical content, kinematics, ionization state, and overall structure of galaxies as well as of the intergalactic medium over $0 < z < 6$.  \\MgII~absorption is an important feature in many QALS, because it relates to metal forming processes.  The \\MgIIdblt~doublet is a strong transition and is easily detected in low ionization QALS, making it an extremely useful tool for probing galaxies and their environments and for studying the structure of the intergalactic medium.  Weak \\MgII~absorbers are defined to be those absorbers with $W_r(2796) < 0.3~\\Ang$, and at least some of them represent a unique population(s) of \\MgII~absorbers that appears to be much different from stronger absorbers \\citep{Rig02, Nest05}.  Unlike strong \\MgII~absorbers, which are typically associated with luminous galaxies (within $\\sim38 h^{-1} (L/L^*)^{0.15}$ kpc) \\citep{Berg91, Berg92, LeBrun93, Steid94, Steid97, Steid95}, weak \\MgII~absorbers at $0.4 < z < 1.4$ are not typically found within $50 h^{-1}$ kpc of such a galaxy \\citep[but see \\citet{Church05} for some exceptions]{Rig02}, raising questions about their processes of formation.  Through both statistical means \\citep{Church99, Rig02} and through investigation of individual systems \\citep{Church00}, weak \\MgII~absorbers are found to arise in sub-Lyman limit systems ($15.8 < \\log{N(\\HI)} < 16.8~[\\cmsq]$), comprising at least 25\\% of \\Lya~forest clouds in that column density range \\citep{Rig02}.  Weak \\MgII~absorbers are also found to have high metallicity; at least 10\\% solar, but in some cases even solar or supersolar \\citep{Rig02, Char03,Simcoe05}.  This is a striking result because, as noted above, luminous galaxies are rarely found within a $50 h^{-1}$ kpc impact parameter of the quasar.  Furthermore, some absorbers have \\FeII/\\MgII~ratios that do not allow for $\\alpha$-enhancement solely from Type II supernovae, suggesting that metals are also produced ``\\textit{in situ}'' by Type Ia supernovae and not ejected by nearby, luminous galaxies.  It is possible that some weak \\MgII~absorbers arise in areas of intergalactic star formation, or are related to supernovae remnants in dwarf galaxies \\citep{Rig02}.  Their incidence is therefore likely to be related to global star formation rates and supernovae activity in these environments.  \\citet{Milni05} have found that a flattened geometry for weak \\MgII~absorbers is most consistent with their number statistics and kinematics, and suggest that such structures could be part of the cosmic web.  Photoionization modeling suggests that the \\MgII~absorption arises in a region $\\sim 1-100$ pcs thick \\citep{Char03,Simcoe05}.  This thin region that produces \\MgII~absorption is most likely surrounded by a lower density region that produces high ionization \\CIV~absorption, with both regions centered at the same velocity \\citep{Rig02}.  Components offset in velocity by tens of \\kms~also often exist in \\CIV~and could represent additional surrounding low density filaments \\citep{Milni05}.  Previous surveys for weak \\MgII~absorbers have been conducted at $0 < z < 0.3$ by \\citet{Nar05} and at $0.4 < z < 1.4$ by \\citet{Church99}.  \\citet{Nar05} used E140M Space Telescope Imaging Spectrograph \\textit{Hubble Space Telescope} archive data to search for \\SiII~$\\lambda 1260$ and \\CII~$\\lambda 1335$ as tracers of \\MgII, and found a number density of absorbers $dN/dz = 1.00 \\pm 0.20$.  \\citet{Church99} used the High Resolution Spectrograph on the Keck I Telescope and observed the \\MgIIdblt~doublet directly, finding $dN/dz = 1.74 \\pm 0.10$.  These values imply that the total cross section for weak \\MgII~absorption is twice that of strong \\MgII~absorbers \\citep{Steid92}.   The findings of \\citet{Nar05} for $0 < z < 0.3$ weak \\MgII~absorbers are consistent with cosmological evolution of the $z \\sim 1$ weak \\MgII~absorber population, based upon $\\Lambda$CDM cosmology \\citep{Church99}.  However, when the effect of the evolution of the extragalactic background radiation (EBR) on the absorber population is taken into account, $dN/dz$ is expected to be higher at $z \\sim 0$ than observed.  This is partly because the less intense ionizing background (compared to $z \\sim 1$) leads to an increase in the low-ionization \\MgII~phase, i.e. absorbers that were below the equivalent width threshold at $z \\sim 1$ are detectable at $z \\sim 0$.  Also, there is a significant contribution at $z \\sim 0$ from \\MgII~absorption in a lower density phase that gave rise only to higher ionization absorbers at $z \\sim 1$ \\citep{Nar05}.  It appears that the processes that lead to weak \\MgII~absorber formation are less active at $z \\sim 0$ then at $z \\sim 1$.  Determining the evolution of $dN/dz$ at higher redshifts is an important step toward identifying the formation processes of weak \\MgII~absorbers, and thus served as the motivation for this work.  In order to push observational measurements of $dN/dz$ for weak \\MgII~absorbers to higher redshifts, we surveyed the spectra of 18 QSOs obtained with the Ultra-Violet Echelle Spectrograph (UVES) on the Very Large Telescope (VLT).  Such a measurement will shed light on the nature of the absorber population, and may indicate which factors affect the evolution of the absorbers.  The data and survey method are outlined in \\S~\\ref{sec:survey}.  Our survey results are summarized in \\S~\\ref{sec:results}.  We also simulated the evolution of the $z = 0.9$ absorber population back to $z = 1.78$, subject to the changing EBR in order to predict the expected $dN/dz$ evolution.  These results are presented in \\S~\\ref{sec:evol} and compared with our survey results.  In \\S~\\ref{sec:conc} we summarize our findings and end with a discussion. ",
        "conclusions": "\\label{sec:conc}  We have conducted a survey of 18 QSO spectra in the ESO archive obtained with the UVES on the VLT, searching for weak \\MgII~absorbers with $0.02 \\leq W_r(2796) < 0.3~\\Ang$.  In order to test our survey completeness we randomly added simulated \\MgIIdblt~transitions to the spectra.  These simulated lines were at the low end of the equivalent width limit.  Because the spectra have high signal-to-noise, the survey is nearly 100\\% complete in the regions that we searched, having eliminated regions strongly affected by telluric absorption.  We found 9 systems in a redshift path length of $\\Delta z = 8.51$ over the range $1.4 < z < 2.4$.  This yields a number density per unit redshift of $dN/dz = 1.06 \\pm 0.12$ at $\\langle z \\rangle = 1.78$.  For comparison, \\citet{Nar05} find $dN/dz = 1.00 \\pm 0.20$ at $\\langle z \\rangle = 0.15$ and \\citet{Church99} find $dN/dz =  1.74 \\pm 0.10$ at $\\langle z \\rangle = 0.9$.  We conclude that the weak \\MgII\\ absorber population peaks at $z \\sim 1$, and that the objects that generate this type of absorption are less common both at higher and at lower redshifts.  A decrease in $dN/dz$ from $z=1.78$ to $z=0.9$ would be expected simply on the basis of cosmological evolution.  The strength of \\MgII\\ absorption coming from a population of objects is also expected to change subject to the changing EBR. In order to understand the balance of these effects, we calculated the expected $dN/dz$ at $z = 1.78$ based upon the known population at $z = 0.9$.  For the purpose of this estimate, we assumed that the absorber population was static or that destruction was balanced by the formation processes between these two epochs, i.e. we only accounted for evolution due to the changing EBR and due to cosmological evolution.  Our prediction of $dN/dz = 1.71-2.14$ is markedly higher than our observed value.  Therefore, we conclude that one of two general scenarios must apply.  First, the structures that give rise to weak \\MgII\\ absorption could be long-lived, and there could be a gradual build-up of these structures over time, such as through hierarchical structure formation.  Second, the structures could be transient, such that the rate of their formation process determines $dN/dz$.  An obvious example of such a process is star formation.  This second possibility was advocated by \\citet{Nar05} who considered the evolution of $dN/dz$ from $z=0.9$ to $z=0$.  In that case, starting with the observed $z=0.9$ population of weak \\MgII\\ absorbers, the evolving EBR along with cosmological evolution would predict a significantly larger population at $z=0$ than is observed.  This issue was complicated by the incidence of a new type of low redshift weak \\MgII\\ absorber descended from absorbers from which only \\CIV\\ absorption was observed at $z=0.9$.  Nonetheless, the trend for an overall decrease in the weak \\MgII\\ absorber population is consistent with the decrease in the global star formation rate over that redshift interval.  This suggests a possible relationship between star formation and weak \\MgII\\ absorbers, which is also consistent with their high metallicities and (in at least some cases) iron abundances \\citep{Rig02}.  We should therefore consider how the global star formation rate relates to the results of our $\\langle z \\rangle = 1.78$ survey.  The global star formation rate is roughly constant over the redshift range $1 \\leq z \\leq 4$, and then decreases to lower redshifts (see \\citet{Gab04} and references therein).  If production of weak \\MgII\\ absorbers was directly correlated with the global star formation rate we would expect a $dN/dz$ at $z=1.78$ consistent with a constant creation rate, subject to the EBR and to cosmological evolution.  Our observation of a smaller $dN/dz$ than this expectation would suggest that the global star formation rate should be smaller at $z=1.78$ than at $z=0.9$.  This is not the case.  However, we need to consider that weak \\MgII\\ absorbers are preferentially found in certain environments.  Theoretically, in the context of hierarchical structure formation and empirically on the basis of the morphology--density relationship, it is expected that star formation peaks later in low mass galaxies \\citep{Kauffmann04}. More directly, \\citet{Bauer05} found a shift over the range $0 < z < 1.5$ in the relative contributions of low and high mass galaxies to the global star formation rate, with low mass galaxies contibuting more at low redshift.   Although measurements of the dwarf galaxy contribution have not yet extended to higher redshifts, it seems likely that the star formation rate in dwarfs peaks at a lower redshift than that for giants.  Given the constant global star formation rate at $1 < z < 4$ \\citep{Gab04}, if the giants are contributing relatively more at high redshift, the dwarfs must contribute less.  Given previous suggestions that weak \\MgII\\ absorbers are related to dwarfs, it is intriguing that the peak epoch of weak \\MgII\\ absorbers may coincide with the peak of star formation in the dwarf environment.  \\citet{Rig02} estimate that weak \\MgII\\ absorbers account for approximately 25\\%-100\\% of \\Lya\\ forest clouds with $15.8 < \\log N(\\HI) < 16.8~[{\\rm cm}]^{-2}$ at $z\\sim1$.  At first glance it seems surprising that the redshift path density for \\Lya\\ forest clouds would sharply decrease with decreasing $z$ until $z=1.5$ ($dN/dz \\propto (1+z)^{2.47\\pm0.18}$) \\citep{Kim02,Wey98}, while we find the weak \\MgII\\ absorber population to increase over the same interval. However, if the weak \\MgII\\ absorbers are closely related to star formation in dwarfs, as we have suggested, then there is no reason to expect a direct correlation of their evolution with the evolution of the \\Lya\\ forest, the latter which is explained by cosmological expansion.  This work was funded by the National Science Foundation grant NSF AST-07138 and by an NSF REU supplement.  We are indebted to the ESO Archive for making this work possible.  The authors wish to express their gratitude to Chris Churchill for his valuable insight and discussions during this project.  Thanks also to Mike Eracleous for his assistance with our analysis.  We also acknowledge Joe Masiero for providing us with tools for our data processing and Andrew Mshar for technical assistance.  We also wish to thank an anonymous referee for invaluable and helpful suggestions."
    },
    "0512/astro-ph0512300_arXiv.txt": {
        "abstract": "{In the framework of the Chandrasekhar-mass deflagration model for Type Ia supernovae (SNe Ia), a persisting free parameter is the initial morphology of the flame front, which is linked to the ignition process in the progenitor white dwarf. Previous analytical models indicate that the thermal runaway is driven by temperature perturbations (``bubbles'') that develop in the white dwarf's convective core. In order to probe the conditions at ignition (diameters, temperatures and evolutionary timescales), we have performed hydrodynamical 2D simulations of buoyant bubbles in white dwarf interiors. Our results show that fragmentation occurring during the bubble rise affects the outcome of the bubble evolution. Possible implications for the ignition process of SNe Ia are discussed. ",
        "introduction": "\\label{intro}  Advances in computer simulations of type Ia supernovae (SN Ia) have led to a rough consensus over the nature of these explosions, even if several issues remain controversial. SNe Ia are believed to be the outcome of thermonuclear explosions of carbon-oxygen white dwarfs (CO WDs) in binary systems.  Among possible explosion scenarios (see \\citealt{hn00} for a review of them) the single degenerate scenario fares best in explaining the observed homogeneity of SNe Ia \\citep{l00}. In this model a CO WD approaches the Chandrasekhar mass by mass transfer from a non-degenerate companion and is disrupted by a thermonuclear explosion.  In this paper we wish to study the processes that lead to the ignition of SNe Ia. The mechanism that ignites the thermonuclear flame in the WD core is the link between the late stages of the progenitor evolution and the development of the initial flame morphology. From a practical point of view, it provides the initial conditions for the explosion simulations. A large set of different initial conditions has been tested in the literature of SN Ia simulations, with results ranging from mildly energetic \\citep{rhn02b,gko03,thr04,rh04} to failed explosions \\citep{cpv04} in pure deflagration explosion models. The initial flame setup in multi-dimensional simulations of SN Ia explosions is considered basically a free parameter, though constrained by some analytical works (\\citealt{gsw95,wwk04,ww04}; Sect.~\\ref{bubbles}). Consequently most simulation schemes are tested against different sets of initial flame shapes, allowing a comparison of the results for various setups \\citep{rhn02b,rh04,gsb05,lah05}.  Among the possible initial conditions the most intuitive one, also for its use in 1D simulations, is the centrally ignited flame \\citep[e.g.][]{nh95a,rhn99b,rhn02a,gko03}. A different ignition condition is provided by a number of spherical flame kernels placed in the innermost part of the WD and detached from the center.  Early works in 2D \\citep{nhw96,rhn99b} implemented only a few blobs per quadrant, while newer simulations \\citep{rhn02b,thr04} have a finer spatial resolution and allow the implementation of more seeds. Nevertheless, the maximum number of blobs that constitute the initial flame is still constrained by resolution rather than physics. This situation may be cured by the use of co-expanding computational grids \\citep{r04}. A hybrid approach between centrally ignited flame and multi-spot scenario is presented by \\citet{rh04} in a full-star 3D simulation.  The influence of different initial conditions on the outcome of the explosion has been discussed in several papers. There is general agreement that an initially larger number of blobs produces more vigorous explosions because the initial flame surface is relatively larger and can provide more flame acceleration. Also the comparison of different initial conditions in \\citet{rh04} shows that the hybrid ``foamy'' initialization provides more seeds for the development of instabilities and gives a larger total energy and production of $^{56}\\mbox{Ni}$ with respect to the model ignited at the center. In the same work, the comparison with the failed explosion found by \\citet{cpv04} is addressed. Their simulation starts from a single spherical flame seed, displaced slightly off-center.  Since the features of SN Ia explosions depend on the number of igniting points in the WD core, an interesting question pertains to an estimate of this number. In the papers cited above the study of this problem was conducted by means of an \\emph{a posteriori} evaluation of the explosion features as a function of the initial conditions, without exploring the physics of the ignition process itself (with the exception of \\citealt{wwk04}). A successful ignition theory has to provide a physical basis for the choice of favored initial conditions in SNe Ia explosion models.  In this paper, a new approach to the ignition problem is based on 2D simulations performed with the FLASH code \\citep{for00}. We will use an indirect approach: the features of the ignition process will be explored by studying temperature perturbations in the WD core (``bubbles'') that are supposed to trigger the ignition (cf.~Sect.~\\ref{bubbles}). Using the results of the bubble simulations, we are able to obtain important clues about the ignition physics.  The paper is structured as follows: in Sect.~\\ref{ignition} the basic ideas of the physics of ignition are introduced. Section \\ref{numerical} describes the setup of the simulations. Section \\ref{simulations} focuses on the bubbles, presenting the relevant physics and the issues related to the 2D simulations. Some important parameters for the bubble evolution are highlighted and their role is explored in Sect.~\\ref{parameter}. In Sect.~\\ref{final-discussion} we use the results of the simulations in order to study possible implications for the ignition of SNe Ia. Our conclusions are summarized in Sect.~\\ref{conclusions}. ",
        "conclusions": "We have studied the ignition process in SNe Ia by simulating the evolution of buoyant bubbles in the WD core. After discussing the bubble physics, the role of the relevant parameters was studied by means of a grid of thirteen simulations. This parameter study shows that floating bubbles may either ignite almost in place or be dispersed by fragmentation. The previous results found an application in the study of ignition process in SNe Ia. The discussion in Sect.~\\ref{implications} provides some clue that the multi-spot ignition model is the favored one. This bubble distribution is consistent with the initial conditions assumed in many recent 3D simulations of SNe Ia explosions.  As pointed out e.g.~by \\citet{hn00}, a successful model for SNe Ia should explain, both, the homogeneity of the class and, hopefully as a function of one or few parameters, the diversity of observational features. The present study indicates that the ignition process contributes to the homogeneity of SNe Ia. The explanation comes from examining the second scenario described in Sect.~\\ref{implications}. According to the PDFs (no matter which one), there is some probability for the occurrence of one (or a few) relatively hot bubble among the the mild temperature fluctuations. In principle, a very hot bubble could lead to a runaway before the colder ones, resulting in premature ignition in one or very few points. The physics of the bubbles leads to a sort of self-regulation because the efficiency of the disruption increases with the bubble temperature (Sect.~\\ref{parameter-t}) and the hottest bubbles are disrupted more effectively. Therefore, one can conclude that a large range of possible ignition conditions is unlikely.  The singly-ignited initial model proposed by \\citet{cpv04} and further analyzed by \\citet{pcl04} could be interpreted as an ignition model coming from a single-bubble runaway. This kind of ignition is not favored by the previous probability arguments.  Of course, this homogeneity will only apply if the underlying convective patterns in the WDs are homogeneous. The diversities in the convective flow \\citep{ww04} are potentially able to affect noticeably the ignition conditions. Their role in producing the observed range of diversity in SNe Ia has not yet been explored.  As far as WD central temperature, bubble temperature and radius of the ignition zone are concerned, the results for the ignition are essentially in agreement with the findings of \\citet{wwk04} and \\citet{ww04}. While our study and the work by \\citet{wwk04} and \\citet{ww04} have the same theoretical background, our approach and methods are widely different.  We discussed the simplified assumptions that this work is based on, and their role in the analysis of the results. In particular, possible limitations are introduced by the uniform background hypothesis (Sect.~\\ref{turb}) and by neglecting adiabatic cooling (Sect.~\\ref{remarks}).  In SNe Ia simulations that implement multi-point ignition, the diameter of the flame seeds (and consequently their number) is set by the spatial resolution of the simulation. In \\citet{rh04} this diameter is $7.0\\ \\mbox{km}$, and future simulations, performed with better computational resources, will be able to resolve progressively smaller scales and allocate more bubbles. The dependence of the explosion features on the number of bubbles, when this number is larger than about 100, opens new insight on the explosion mechanism of SNe Ia \\citep{rhn05}. At this stage the increased spatial resolution will not help to improve the results unless there is a better understanding of the link between the progenitor evolution and the early explosion phase. Future SN simulations should also take into account the short ($\\sim 0.1\\ \\mbox{s}$) temporal evolution of the ignition process, and its interplay with the ongoing explosion."
    },
    "0512/astro-ph0512136_arXiv.txt": {
        "abstract": "Binary mergers involving black holes and neutron stars have been proposed as major sources of gravitational waves, r--process nucleosynthesis, and gamma ray bursters. In addition, they represent an important, and possibly unique, observable that could distinguish between normal and self--bound neutron stars. These two families of stars have distinctly different mass--radius relationships resulting from their equations of state which can be revealed during their mergers if stable mass transfer ensues.  We consider two cases of gravitational-radiation induced binary mergers: (i) a black hole and a normal neutron star, and (ii) a black hole and a self-bound strange quark matter star.  We extend previous Newtonian analyses to incorporate the pseudo-general relativistic Paczy\\'nski-Wiita potential or a potential correct to second--order post-Newtonian order in Arnowitt--Deser--Misner coordinates. These potentials are employed to study both the orbital evolution of the binary and the Roche lobe geometry that determines when and if stable mass transfer between the components is possible. The Roche lobe geometry with pseudo-general relativistic or post-Newtonian potentials has not heretofore been considered.  Our analysis shows that differences in the evolution of normal neutron stars and strange quark matter stars are significant and could be detected in gravity waves. Both the amplitude and frequencies of the wave pattern are affected. In addition, details of the equation of state for either normal neutron stars or strange quark stars may be learned. A single merger could reveal one or two individual points of the mass-radius relation, and observations of several mergers could map a significant portion of this relation. ",
        "introduction": "Mergers of compact objects in binary systems, such as a pair of neutron stars (NS-NS), a neutron star and a black hole (NS-BH), or two black holes (BH-BH), are expected to be prominent sources of gravitational radiation \\citep{THORNE1}.  The gravitational-wave signature of such systems is primarily determined by the chirp mass $M_{chirp}=(M_1M_2)^{3/5}(M_1+M_2)^{-1/5}$, where $M_1$ and $M_2$ are the masses of the coalescing objects. The radiation of gravitational waves removes energy which causes the mutual orbits to decay. For example, the binary pulsar PSR B1913+16 has a merger timescale of about 250 million years, and the pulsar binary PSR J0737-3039 has a merger timescale of about 85 million years \\citep{LYNE04}, so there is ample reason to expect that many such decaying compact binaries exist in the Galaxy.  Besides emitting copious amounts of gravitational radiation, binary mergers have been proposed as a source of the r-process elements \\citep{LSCH2} and the origin of the shorter-duration gamma ray bursters \\citep{EICHLER89}.   The expected rates of binary mergers have been estimated by either utilizing the observational information from known binaries or else through theoretical models of binary stellar evolution (population synthesis).  The coalescence rate for the NS-NS case has been recently revised upward in view of the newly discovered binary PSR J0737-3039 \\citep{LYNE04}. Due to its low luminosity and relatively short lifetime, the estimated NS-NS merger rate has been increased from prior estimates by almost an order of magnitude \\citep{BURGAY03, KALOGERA03}.  The merger rates can be coupled with expected characteristics of future gravity wave detectors to make predictions of the observed merger rate.  The updated analysis predicts the observed NS-NS merger rate to be $1.8\\times 10^{-4}$ yr$^{-1}$ galaxy$^{-1}$ with an uncertainty of about a factor of 3 to 10.  For the advanced LIGO detector, it is estimated that mergers up to distances of 350 Mpc \\citep{FINN2001} could be observed, with a total detection rate of up to approximately $1$ day$^{-1}$ anticipated \\citep{KALOGERA03}.  In the NS-BH case the merger rate has been estimated to be about $10^{-5}-10^{-4}$ yr$^{-1}$ galaxy$^{-1}$ \\citep{BETHE98,PORTEGIES98}, or up to 6 times larger than the NS-NS merger rate.  It seems likely that this rate might also be adjusted upward by the discovery of PSR J0737-3039.   A neutron star in a binary merger with a larger mass companion will be tidally disrupted.  If the tidal disruption occurs far enough outside the innermost stable circular orbit (ISCO) of the binary, it is expected that an accretion disc will form or that mass transfer to its companion will occur \\citep{CLARK1,Jaranowski92,PORTEGIES98b}. In either case, the gravity wave signal is expected to be significantly different than if tidal disruption occurs after the star penetrates the ISCO.  Tidal disruption will occur at larger separations for larger neutron star radii, so some aspects of the neutron star equation of state (EOS) could be determined in the case in which tidal disruption occurs outside the ISCO.  Depending upon the equation of state (EOS) of dense matter and the mass ratio of the binary, tidal disruption could lead to the less massive star overflowing its Roche lobe and transferring its mass through the inner, or first, Lagrange point to its more massive companion.  If mass loss from the lighter star occurs quickly enough, and if conservation of mass and orbital angular momentum can be assumed, the two stars will then begin to spiral apart.  Although this increases the Roche lobe volume, the radius of the neutron star also increases in response to the mass loss.  Mass transfer will continue in a stable fashion if the lighter star can expand sufficiently fast such that it is able to continuously fill its Roche surface.  We refer to mass transfer under such conditions as {\\em stable mass transfer}. Because the orbital separation now increases, the gravity wave amplitude will decrease.  The signature of stable mass transfer in gravity waves should therefore be strikingly different than for an amorphous tidal disruption or a direct plunge.  As we show below, important information about the radius of a neutron star and the underlying equation of state could be contained in the gravity wave signal. The longer timescale produced by stable mass transfer might also extend the duration of an associated gamma ray burst \\citep{PORTEGIES98b}. In the absence of mass transfer, this timescale might be of order milliseconds, the orbital period of the binary near the last stable orbit.  In contrast, stable mass transfer extends over a period of at least several tenths of a second.  Another effect of stable mass transfer would be to modify the amount of material potentially ejected from the system.  Matter from a tidally disrupted neutron star, which could be accelerated to escape velocities from the binary \\citep{LSCH2}, undergoes decompression which results in heavy nuclei and an intense neutron flux leading to the copious production of r-process elements \\citep{LATTIMER77,MEYER89}. In the case in which stable mass transfer occurs, sudden disruption of the neutron star near the last stable orbit is avoided, but mass could be transferred unstably at later times and larger separations when the neutron star approaches its minimum stable mass \\citep{COLPI93}.     One of the most significant aspects of the equation of state that stable mass transfer could reveal is whether or not the neutron star is actually a strange quark matter star.  It has been suggested that if strange quark matter is the ultimate ground state of matter (i.e., has a lower energy at zero pressure than iron) compression of neutron star matter to sufficiently high density triggers a phase transition which converts virtually the entire neutron star to strange quark matter \\citep{MADSEN99}.  Such a star is self-bound as opposed to being gravitationaly bound as is the case of a normal neutron star.  It has so far proved very difficult to find venues from astrophysical observations that could unambiguously distinguish strange quark stars from normal neutron stars.  This is because self-bound stars have similar radii, moments of inertia, and neutrino emissivities and opacities to that of moderate mass normal neutron stars.  Therefore, it may be unlikely that photon or neutrino observations, or radio binary pulsar timing measurements, will be able to differentiate these cases, especially if strange quark stars have a small hadronic crust, supported perhaps by electrostatic forces.  In that case, the effective temperatures and radii of solar-mass-sized strange quark stars and normal neutron stars would tend to be similar.  Even during the proto-neutron star stage, which is observable through neutrino emissions \\citep{Burrows86}, these two types of stellar configurations likely yield similar neutrino signals until such late times that the low luminosities prevent an unambiguous discrimination \\citep{Prakash01}.  However, major differences in the evolutions of normal neutron stars and strange quark matter stars emerge during the final stages of binary mergers if stable mass transfer occurs. These differences would be prominent in both the amplitudes and frequencies of gravitational wave emisions.    The Newtonian equations for orbital motion become progressively inadequate as compact objects spiral inward.  For example, the existence and location of the innermost stable circular orbit (ISCO) are not predicted by Newtonian gravitation.  Furthermore, in calculations performed to date, the gravitational equipotential surfacws and the size of the Roche lobe have been computed using a Newtonian background \\citep{KOPAL1,PACZYNSKI1,EGGLETON1}). Although a solution in full GR is not yet possible, systematic post-Newtonian corrections have been developed to explore regions close to compact objects (see, for example, \\citep{BLANCHET3}).  To extend the Roche lobe overflow model for mass transfer, we will employ recently calculated corrections to the Newtonian results by using the second order post--Newtonian approximation in Arnowitt--Deser--Misner coordinates \\citep{ROCHE2PN}.   We adopt a semi-analytic approach to describe the final stages of NS-NS or BH-NS mergers, extending earlier treatments \\citep{Prakash03}.  While the evolution of such mergers has been previously calculated with numerical techniques involving three-dimensional Newtonian hydrodynamics \\citep{KLUZNIAK1}, only a few simulations have implemented GR corrections to the orbital evolution or have considered the case in which the inspiralling stars have unequal masses \\citep{DAVIES2005,BEJGER05,SHIBATA05a,SHIBATA05b}.  In the final phase of a merger, Newtonian orbital mechanics cannot be expected to hold and general-relativistic (GR) corrections become important. In the few cases in which unequal masses were considered, the mass ratios were severely restricted because of numerical difficulties \\citep{Shibata03}.   This work is organized as follows.  In \\S \\ref{sec:General}, we provide the general considerations upon which this work is based on and discuss the mechanism of Roche lobe overflow. The three main ingredients: (1) the evolution of orbital angular momentum, (2) the evolution of the Roch lobe radius, and (3) the EOS parameter that governs these evolutions, are then presented. The differential equations that govern the evolution of the separation distance $a$ and the mass ratio $q=M_1/M_2$ are derived in this section along with the condition for stable mass transfer.  In \\S 3, the EOS's of normal neutron stars and self-bound quark stars used in this work are summarized.  \\S 4 provides the evolution equations for mergers using the second order post--Newtonian potential.  We summarize here the recently calculated results for Roche lobes in the second order post--Newtonian approximation \\citep{ROCHE2PN}.  In \\S 5, we introduce, in addition to the Newtonian and 2PN analyses, two pseudo--general relativistic potentials for comparisons of binary evolution simulations. We first introduce the potential due to \\citet{PACZYNSKI2}, and then we modify it to construct a hybrid potential that correctly incorporates some post--Newtonian features.  We fit its parameters to match predictions for the position of the innermost stable circular orbit (ISCO) from \\citet{BLANCHET2003}, who evaluated the effective gravitational potential up to order $(v^2/c^2)^3$ in post-Newtonian order.  We extend the Newtonian analysis of Roche lobes to include the Paczy\\' nski--Wiita \\citep{PACZYNSKI2} and hybrid potentials and provide simple fitting results.  We consider the corresponding changes to the orbital dynamics that affects the gravitational radiation reaction and the location of the ISCO.  Regions of stable mass transfer in the $M_1-M_2$ plane are identified in \\S 6 for all potentials considered.  In \\S \\ref{sec:evolution}, we present results of merger simulations for both a normal star and a self-bound star orbiting around another more massive object (BH or NS).  A variety of equations of states with varying stiffness at supranuclear density for both cases are considered.  Our results pertaining to gravity wave signals are presented in \\S 8.  Conclusions are contained in \\S \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have extended Newtonian models of binary orbital evolution and Roche lobe geometry to second post-Newtonian order to examine the final stages of compact binary mergers in which the lighter star is a neutron star or a strange quark matter star.  In our simulations, binary mergers can be categorized into two distinct classes, depending upon whether or not stable mass transfer occurs.  We find that similar, non-negligible, regions in mass space ($M_1-M_2$) allow for the possibility of stable mass transfer for both normal neutron stars and strange quark matter stars.  Mass transfer is not expected to occur if (1) the binary mass ratio $q=M_1/M_2$ is too close to unity, or (2) the total mass of the system is too large ($M\\gtrsim10$ M$_\\odot$).  The limit on $q$ for normal stars is $q\\gtrsim0.75$, but for strange quark matter stars it is $q\\gtrsim 0.9$.  Binary mergers in which stable mass transfer is able to occur behave dramatically different than those in which plunge occurs. Qualitatively, normal stars and SQM stars follow a similar evolution: stable mass transfer initiates a period of reverse, or outwards, spiralling, characterized by diminishing frequencies and gravitational wave amplitude, that has a duration of 5--10 s in the case of normal stars and lasts essentially forever in the case of SQM stars.  In contrast, the gravitational wave emission of a plunge is characterized by a single burst of high-frequency radiation.  These differences should be distinguishable from gravitational radiation observations.  In the case of stable mass transfer, large differences in the gravity wave signal following the onset of mass transfer between the cases of normal neutron stars and strange quark matter stars are expected. This signature may be unique in its ability to distinguish between these stellar models from astrophysical observations.  Whereas normal neutron star evolutions will be characterized by rapidly diminishing frequencies, strange quark matter star evolutions will have a radiation frequency that reaches an asymptotic value.  In addition, stable mass transfer from a normal neutron star will have a finite duration of order 5--10 s, concluding when the star expands quickly and overfills its Roche lobe when its mass decreases to about 0.16 M$_\\odot$, which will be quickly (of order a few ms) followed by the violent decompression of the star when it reaches its minimum mass (about 0.09 M$_\\odot$).  In contrast, mass transfer should continue essentially forever in the case of an SQM star, with a gravitational wave amplitude that decreases with time as $t^{-1}$.  Several observational constraints become possible from mergers in which stable mass transfer occurs.  Most important is the ability to distinguish between the star being normal or composed of strange quark matter.  In either case, the possibility exists to extract the radius-mass relation.  Not only will a distinct $M-R$ point be potentially measureable from a single system, corresponding to the onset of stable mass transfer, but in the case that the star is normal, an additional point is possible corresponding to the termination of stable mass transfer.  The overall signal may also allow estimation of the function $\\alpha(M)=d\\ln r/d\\ln M$ for intermediate masses.  Moreover, observations of many different events will allow sampling of a corresponding number of $M-R$ points.   Research support of the U.S. Department of Energy under grant number DOE/DE-FG02-87ER-40317 for all authors and grant number DOE/DE-FG02-93ER-40756 for Madappa Prakash is gratefully acknowledged."
    },
    "0512/astro-ph0512650_arXiv.txt": {
        "abstract": "We report on high resolution Echelle spectroscopy of 20 giant stars in the Galactic old open clusters NGC 6791 obtained with Hydra at the WIYN telescope. High precision radial velocity allow us to isolate 15 {\\it bona fide} cluster members. From 10 of them we derive a global [M/H]=+0.39$\\pm$0.05. We therefore confirm that NGC 6791 is extremely metal rich, exhibits  a few marginally sub-solar abundance ratios, and within the resolution of our spectra does not show evidences of spread in metal abundance. With these new data we re-derive the cluster fundamental parameters suggesting that it is about 8 Gyr old and 4.3 kpc far from the Sun. The combination of its chemical properties, age, position, and Galactic orbit hardly makes NGC 6791 a genuine Population I open cluster. We discuss possible interpretations of the cluster peculiarities suggesting  that the cluster might be what remains of a much larger system, whose initial potential well could have been sufficient to produce high metallicity stars, and which has been depopulated by the tidal field of the Galaxy. Alternatively, its current properties may be explained by the perturbation of the Galactic bar on an object originated well inside the solar ring, where the metal enrichment had been very fast. ",
        "introduction": "NGC~6791 is an extremely interesting and intriguing open cluster.  The combination of old age, small distance and high metal abundance makes this cluster very attractive, and indeed in the last 40 years it has been the target of intensive and numerous studies (Carney et al. 2005, and references therein).  A large number of optical photometric studies (Stetson et al. 2003 and references therein) has been recently complemented by the deep ACS/HST investigation by King et al. (2005), and the near IR study by  Carney et al. (2005).\\\\ Since the pioneering study of Kinman (1965) it was clear that NGC 6791 is a very old and very metal rich cluster.  Its age was measured several times by using different sets of isochrones (Carraro et al. 1999 and references therein), and it is probably confined in the range from 8 to 12 Gyr, depending on the cluster precise metal abundance. Taylor (2001, and references therein) critically reviewed all the available metallicity estimates, concluding that the [Fe/H] for NGC 6791 should probably lie in the range +0.16 to +0.44 dex. This combination of age and metallicity is unique in the Milky Way open cluster population, and has been recently questioned by Bedin et al. (2005), whose HST study of the White Dwarf  cooling sequence supports a much younger age. As these authors comment, this age discrepancy may arise from defects in the White Dwarf current models, or  from the poorly known cluster metal abundance.\\\\  Noteworthy, the cluster is also known to harbour a number of sdB/sdO stars (Landsman et al. 1998, Buson et al. 2005), which may be explained by a  scenario of a high metallicity driven wind in the Red Giant Branch phase of the progenitors of these stars or, more simply, by the binarity hypothesis (Green et al. 2005) .  The UV upturn (namely the abrupt rise in the UV continuum emission shortward of $\\lambda \\approx 2000\\AA$)  similar to that typical of any ellipical galaxy (Landsman et al. 1998) and the highly eccentric orbit, unusual for a Population I object, contribute to make this cluster even more intriguing .\\\\ In an attempt to substantially improve our knowledge of NGC ~6791, we carried out a spectroscopic campaign to provide radial velocities and accurate metallicities of a statistically significant number of stars in  the cluster. In fact, current abundance determinations either lack sufficient resolution or are restricted to a very small number of stars.\\\\ This new set of abundance estimates coupled with the high quality of existing photometry (Stetson et al 2003) allow us to significantly improve on the fundametal parameters of this cluster, and better clarify its intriguing nature. ",
        "conclusions": "The hardest point with NGC~6791 is how inside this cluster such high metallicity stars could have been produced. In fact, this cluster does not have any counterpart in the Milky Way.  We note here that Kinman (1965) originally identified NGC~6791 as a globular cluster.  Even if this interpretation were to be adopted, the high metallicity of NGC~6791, much higher than that of any Galactic globular cluster or nearby dwarf galaxy, remains mysterious. \\\\   With Berkeley~17 and Collinder~261, NGC 6791 is one of the oldest open clusters of the Galaxy (Carraro et al. 1999), but its metal abundance is incomparably higher. Moroever, NGC 6791 is one of the most massive open clusters (4000 $M_{\\odot}$ at least). It lies at 1 kpc above the Galactic plane, inside the solar ring. This combination of mass and position is hard to explain, since the interaction with the dense Galactic environment should strongly depopulate a typical open cluster. \\\\   NGC 6791 is routinely considered in the studies of the chemical evolution of the Galactic disk, and occupies a unique position in the Galactic disk radial abundance gradient (see Fig.~9).  By including NGC 6791, the slope of the gradient changes from -0.05 (solid line) to -0.07 (dashed line).  Besides, if one considers the slope defined only by clusters older than 4 Gyr (Friel et al. 2002, Fig.~3, upper panel), the slope doubles, from -0.06 to -0.11.\\\\ In the same figure, the horizontal solid line indicates the epicyclical amplitude of NGC 6791 orbit (see below, and Carraro \\& Chiosi 1994).  One can readily see how NGC 6791 is quite an exotic object. If, by chance, at the present time the cluster would be at different orbit phase, which would put it, e.g., beyond 12 kpc, there would be a drastic change and even an inversion of the slope of the Galactic disk abundance gradient.  Finally, the position of this cluster in the Galactic disk Age-Metallicity relationship (Carraro et al. 1998) is puzzling as well, since the cluster significantly deviates from the mean trend.\\\\  In Fig~10 we present NGC~6791 Galactic orbit. This was obtained by integrating back in time ( 1 Gyr) the cluster from its present position and kinematics using the Galaxy N-body/gasdynamical model by Fux (1997, 1999). The adopted radial velocity and proper motions come from Geisler (1988) and Cudworth (private comunication), whereas the Galactocentric rectangular  initial conditions (positions and velocities) was derived as in Carraro \\& Chiosi (1994).\\\\ \\noindent Intererstingly enough, this plot shows that the cluster moves from the outer disk regions of the Milky Way, more than 20 kpc far away from the Galactic center, and enters the Solar Ring going as close as 6 kpc from the Galactic center.\\\\  The eccentricity (e=0.59) of this orbit is quite high for a Population I star cluster (Carraro \\& Chiosi 1994), and it is much more similar to a globular cluster/dwarf galaxy orbit.   A plausible scenario  is that NGC 6791 is what remains (the nucleus) of a much larger system, which underwent strong tidal disruption. This would explain the cluster orbit, and provide a reasonable explanation for the high metallicity of its stars, which  could have been produced only inside a deep potential well.  However, within the observational errors, we did not find any significant abundance spread. This would mean that the bulk of the stars in the cluster was produced in a single burst of star formation. This fact makes more difficult the capture interpretation since Local Group Galaxies normally exhibit spreads in metal content and possess lower metal abundance (Mateo 1998). We stress however the fact that our results are based on only ten stars, and that only larger spectroscopic surveys can better address this particular problem.\\\\   An alternative more conservative scenario is that the cluster was born in the inner side of the Galaxy, close to the bulge, where the metal enrichment has been  fast. Grenon (1999) studied the kinematics of a group of old (10 Gyrs) metal rich $[M/H] \\geq 0.30$ stars and suggested that they formed close to the bulge and then migrated at large Galactocentric distance due to the perturbation of the Galactic bar.  The orbit we calculated actually includs the effect of the bar, and NGC 6791 indeed moves well outside the solar circle. NGC 6791 is very concentrated for an open cluster, and spent most of its time at moderate Galactic latitude. This might help to explain its survival."
    },
    "0512/astro-ph0512466_arXiv.txt": {
        "abstract": "We present MAMBO 1.2~mm observations of five $BzK$-pre-selected vigorous starburst galaxies at z~$\\sim2$.  Two of these were detected at more than 99.5\\% confidence levels, with 1.2~mm fluxes around 1.5~mJy.  These millimeter fluxes imply vigorous activity with star-formation rates (SFRs) $\\approx$500--1500$~M_{\\sun}~yr^{-1}$, confirmed also by detections at 24$\\mu$m with the MIPS camera on board of the Spitzer satellite.  The two detected galaxies are the ones in the sample with the highest SFRs estimated from the rest-frame UV, and their far-IR- and UV-derived SFRs agree reasonably well. This is different from local ULIRGs and high-z submm/mm selected galaxies for which the UV is reported to underestimate SFRs by factors of 10--100, but similar to the average $BzK$--ULIRG galaxy at $z\\sim2$.  The two galaxies detected at 1.2~mm are brighter in $K$ than the typical NIR-counterparts of MAMBO and SCUBA sources, implying also a significantly different $K$-band to submm/mm flux ratio. This suggests a scenario in which $z\\sim2$ galaxies, after their rapid (sub)mm brightest phase opaque to optical/UV light, evolve into a longer lasting phase of $K$-band bright and massive objects.  Targeting the most UV active $BzKs$ could yield substantial detection rates at submm/mm wavelengths. ",
        "introduction": "In the last few years, selection techniques at different wavelengths have revealed a variety of, at first glance, different high redshift source populations (e.g., Steidel et al. 1996; 2004; Smail et al. 1997; Kurk et al. 2000; Franx et al. 2003; Daddi et al. 2004a; Cimatti et al. 2004).  Highly debated questions are how these objects are linked to each other; how much they overlap with each other; and what the evolutionary path is, if any, between them and with the different types of low redshift galaxies.  Observations in the local universe are indicating \\citep[e.g.,][]{tho05,nel05} that the stars in massive early-type galaxies formed in general at high redshift, $z>>1$, with short formation timescales, suggesting high peak star-formation rates (SFRs). Populations of high redshift galaxies, among those mentioned above, having large stellar masses and high SFRs have indeed also been found. However a detailed characterization of the formation of early-type galaxies is still missing.  Recently, \\citet{dad04b} presented a technique based on optical/near-IR photometry in the $B$, $z$ and $K$ bands that, at least to $K_{Vega}<20$, allows to obtain virtually complete samples of galaxies at redshift $1.4<z<2.5$, including star-forming and passively evolving galaxies. This technique is reddening independent for $z\\sim2$ star-forming galaxies, that in $K<20$ samples have SFRs estimated between $50-1000$ solar masses per year. These SFRs require large and possibly uncertain reddening corrections.  The average long-wavelength emission properties of $BzK$ star-forming galaxies ($sBzK$s hereafter) at $z=2$ have been measured accurately by \\citet{dad05}, finding that the typical $sBzK$ is an Ultra Luminous Infra-Red Galaxy (ULIRG) with $SFR\\sim200-300~M_{\\sun}~yr^{-1}$.  As the $BzK$ technique is substantially complete for $z\\sim2$ galaxies, one expect that also submm/mm selected galaxies (SMGs, see Blain et al. 2002 for a review) should fulfill the $BzK$ criteria, as many of them lie at $z\\sim2$ \\citep[e.g.,][]{cha05}, and that the most active $sBzK$ galaxies could actually be SMGs.  In this letter we present observations with the 117--element {\\it Max-Planck Millimeter Bolometer Array} \\citep[MAMBO;][]{kre98} of five luminous ($K_{Vega}\\simlt19$) $sBzK$ galaxies, with high expected SFRs, based on their optical/UV properties.  We report the detection of two of the five galaxies, and discuss some implications on the relation between $sBzK$ and SMGs, and on the assembly process of massive galaxies at high redshifts.  We assume a Salpeter initial mass function (IMF) from 0.1 to 100 $M_\\odot$, and a cosmology with $\\Omega_\\Lambda, \\Omega_M = 0.73, 0.27$, and $h = H_0$[km s$^{-1}$ Mpc$^{-1}$]$/100=0.71$ \\citep{spe03}. ",
        "conclusions": "The fluxes measured at 1.2mm (close to the IR-peak of the dust emission) and at 24$\\mu$m with Spitzer (in the mid-IR rest-frame region) imply that the two detected $sBzKs$ are active IR galaxies. To obtain accurate estimates of their total IR luminosities (8--1000$\\mu$m) would require further observations between 24$\\mu$m and 1.2mm, and especially in the range 60--100$\\mu$m rest-frame, that could constrain the temperature of the dust. For $z=2$, obtaining these data to the required depths is beyond the possibilities of the available observing capabilities. However, we can derive constraints on the total IR luminosities (and thus SFRs) by assuming these 2 detections have IR SEDs similar to those of well studied $z=2$ populations, i.e. the average $sBzK$ galaxies (Daddi et al. 2005) or SCUBA selected galaxies (Chapman et al. 2005), or by assuming that local empirical correlations between mid-IR or submm luminosities and $L_{IR}$ (e.g. Chary \\& Elbaz 2001) holds also at $z=2$. We find that this suggests IR luminosities in the range of $L_{IR}\\approx3$--$10\\times10^{12}L_\\odot$, which correspond to SFRs in the range 500--1500$M_{\\sun}~yr^{-1}$, the factor of 3 uncertainty being indicative of the range of estimates from the different methods.  We considered the possibility that the detected mm fluxes are due to AGN emission. The submm-to-X-ray spectral slope $\\alpha_{SX}$ (Alexander et al. 2003) is an indicator of AGN activity.  For the two MAMBO-detected $sBzK$s we derive an upper limit of $\\alpha_{SX}\\simlt1.20$\\footnote{We assume ratios of $2.5-3$ between 850$\\mu$m and 1.2~mm fluxes, that are the factors spanned by the models at the luminosities of our objects.} that, if compared to Fig.~7 in \\citet{ale03}, points toward the starburst part of the diagram occupied, e.g., by SCUBA galaxies in HDFN and Arp~220, and exclude both luminous unobscured (3C273) and obscured quasars. Of course, this does not allow us to completely discard the presence of Compton thick AGNs (NGC~6240).  However, the lack of any AGN features in the UV spectra, together with the reasonable agreement between UV- and far-IR- estimated SFRs, suggest that the far-IR activity of these sources is powered by star formation.   An interesting result of our observations is indeed that the IR luminosities and related SFRs inferred for our objects from the mm fluxes are in quite a good agreement, at most a factor 2--3 larger than those estimated from the dust-reddening corrected UV luminosities (Fig.~\\ref{fig:sfr}, top panel). This agreement, already noticed by \\citet{dad05} for the typical $sBzK$ galaxy, is somewhat surprising. A fairly different situation seemingly holds for submm/mm selected galaxies (Chapman et al. 2005; Swinbank et al. 2004) and for local far-IR selected ULIRGs (Goldader et al. 2002, with only a notable exception), where the observed IR luminosities and far-IR estimated SFRs exceed by 1--2 orders of magnitude the same quantities estimated from the dust corrected UV luminosities (Fig.~\\ref{fig:sfr}, top panel).  Extending this comparison further (Fig.~\\ref{fig:sfr}, bottom panel), we notice that $sBzKs$ galaxies in this paper and in Daddi et al. (2005) are generally brighter in the $K$ band ($\\sim18-20$~mag) than the typical counterparts of MAMBO galaxies \\citep{ber00,dan02,dan04} and SCUBA sources \\citep[e.g.,][]{bla02,sma04}, although they lie at a similar redshift $z\\sim2$. The mm/submm fluxes of SMGs are instead generally higher than those of $sBzKs$ (Fig.~\\ref{fig:sfr}).  This implies a significantly different $K$-band to submm/mm flux ratio between far-IR selected vs. optically selected $z\\sim2$ galaxies.  Similarities also exist between $sBzK$s and far-IR selected populations both in their irregular morphologies (Conselice et al. 2004; Trentham et al. 1999; Daddi et al. 2004a; this paper, see Fig.~\\ref{fig:mambob}) and in their red observed UV slopes (i.e. high apparent reddening; Chapman et al. 2005; Daddi et al. 2004b; 2005; Goldader et al. 2002; this paper, see Table~1).  The differences (submm/mm flux levels, $K$-band to submm/mm ratios and $SFR_{UV}$ to $SFR_{IR}$ ratios) and similarities (morphologies, UV slopes) are most likely effects induced by the different selection of the samples ($K$-band versus far-IR), and can be reasonably understood by assuming that the far-IR emitting region of most SMGs is observed behind a thick dust screen, nearly opaque to the optical and near-IR radiation. This can be fit into an evolutionary scenario in which the starburst galaxies are initially deeply embedded into a large column density of dust, during the submm/mm brightest phase when substantial obscuration of a few magnitudes or more is present even at observed $K$-band (rest-frame $5000-8000$~\\AA\\ for $z=2-3$). The UV light is almost completely attenuated behind the dust screen so that it is impossible to measure the ongoing SFR from the residual UV light transmitted. As time passes by, SFR and L$_{IR}$ are reduced and the dust is diffused and ejected out of the main star-formation regions (presumably by galactic winds or other kinds of feedback), or simply destroyed. Thus the resulting galaxy is brighter in the near-IR and transparent enough in the rest frame UV that the transmitted light can account for the ongoing SFR after reddening correction.  The submm brightest phase has presumably a short duration \\citep[$\\sim40-200~Myr$;][]{sma04,gre05}, while the subsequent {\\em normal} $sBzK$ phase is estimated to last quite longer ($\\sim0.5$--1~Gyr; Daddi et al. 2005), which would explain why SMGs have a factor of 10 or more lower spatial density \\citep{sco02,dad04a,dad04b}, while having similar masses \\citep[e.g.,][]{gen03,gre05} and perhaps spatial clustering \\citep{bla04,dad04a,ade05}.  This picture is consistent with the evolutionary scenario proposed by Vega et al. (2005).  Chapman et al. (2005) also suggest that SMGs evolve into near-IR bright galaxies.  We conclude by suggesting that the follow-up of the most UV active $sBzK$ could be a promising alternative for finding submm/mm sources, possible complementary to the radio pre-selection method presented by \\citet{cha01} for a similar redshift range $z\\sim1.5-2.8$, or the Spitzer$+$MIPS pre-selection (Lutz et al 2005).  \\smallskip  We would like to thank the IRAM staff at Pico Veleta for their support of these observations; Fabian Walter for his careful reading of an earlier version of this manuscript and his insightful and helpful comments; Axel Wei\\ss\\ for useful conversation about the MAMBO data; and Ian Smail for useful discussions.  ED gratefully acknowledges NASA support through the Spitzer Fellowship Program, under award 1268429. This work is partly supported by a Grant-in-Aid for Scientific Research (No. 16540223) by the Japanese Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0512/astro-ph0512185_arXiv.txt": {
        "abstract": "We present simultaneous UV-G-R-I monitoring of 19 M~dwarfs that revealed a huge flare on the M9~dwarf % with an amplitude in the UV of at least 6 magnitudes. This is one of the strongest detections ever of an optical flare on an M~star and one of the first in an ultracool dwarf (UCD, spectral types later than about M7). Four intermediate strength flares ($\\Delta m_{\\rm UV}<4\\,{\\rm mag}$) were found in this and three other targets. For the whole sample we deduce a flare probability of $0.013$ (rate of $0.018\\,{\\rm hr}^{-1}$), and $0.049$ ($0.090\\,{\\rm hr}^{-1}$) for 2M1707+64 alone. Deviations of the flare emission from a blackbody is consistent with strong $H_{\\alpha}$ line emission. We also confirm the previously found rotation period for 2M1707+64 \\citep{Rockenfeller2006a} and determine it more precisely to be $3.619\\pm 0.015\\,{\\rm hr}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512520_arXiv.txt": {
        "abstract": "We present high-resolution ($1\\farcs8$ -- $4\\farcs5$) CO data of the Virgo spiral galaxy NGC 4569, obtained using the Nobeyama Millimeter Array. We found that the molecular gas is highly concentrated in the circumnuclear region with two off-center peaks. A CO image with the highest angular resolution of $2\\farcs0 \\times  1\\farcs8$ shows that six blobs likely form a circumnuclear elliptical ring (CER) with a semi-major axis radius of 720 pc. The CER shows a strongly twisted velocity field, and the position--velocity (PV) diagram shows significant forbidden velocity components. These kinetic features are understood as being non-circular motion due to a bar-potential. We found that the CER coincides with the H$\\alpha$ bright central core and that the mass ratio of the molecular gas to the dynamical mass is about 18\\%. These results support a gaseous inflow scenario induced by a weak bar potential and self-gravity of the gas. ",
        "introduction": "NGC 4569 is a member of the Virgo cluster galaxies and its nucleus is classified as a starburst-dominated LINER \\citep{alo00}. The mechanism for causing nuclear activity is still unveiled. A CO line observation is useful to explore the physical conditions of the central regions of spiral galaxies to understand this issue, because the molecular gas is a dominant component among neutral gases in the central region.  \\citet{sak99} observed this galaxy using the Nobeyama Millimeter Array (NMA) at the Nobeyama Radio Observatory (NRO) in the $^{12}$CO($J=1$--$0$) line, and obtained a resolution of $3\\farcs9 \\times 3\\farcs6$. They showed the gas concentration in the central region, two significant peaks on the central concentration, and strong non-circular motions. \\citet{hel03} obtained $^{12}$CO($J=1$--$0$) line data with a resolution of $6\\farcs6 \\times 5\\farcs5$ in the BIMA SONG project. They found extended CO emission up to about $1\\arcmin$ radius.   In order to investigate the more detail kinematics in the central region of spiral galaxies, we achieved the Virgo high-resolution CO-line survey (ViCS) during a long-term project at the NMA from 1999 December to 2000 March to obtain the highest resolution CO data (Sofue et al., 2003, Paper I hereafter ). We presented CO data of NGC 4569 with low-resolution ($4\\farcs5 \\times 3\\farcs1$) in Paper I. In this paper as Part V of the series of the ViCS, we present the highest resolution ($2\\farcs0\\times 1\\farcs8$) CO data, and discuss the distribution of the interstellar matter (ISM) in the central region and the relationship to the nuclear activity.  An optical $B$-band image of this galaxy is shown in figure \\ref{dss}, which was taken from STScI Digitized Sky Survey (DSS). The properties of this galaxy are summarized in table \\ref{prop-n4569}. NGC 4569 is assumed to be at a distance of $16.1$ Mpc, which is the distance of the Virgo cluster member NGC 4321 \\citep{fer96}. The angular size of $1\\arcsec$ corresponds to 78 pc. {A position angle (P.A.) of $18\\arcdeg$ and an inclination of $63\\arcdeg$ are adopted \\citep{nak04}. }   \\vskip 3mm \\centerline{--- Figure \\ref{dss} ---} \\vskip 3mm  \\vskip 3mm \\centerline{--- Table \\ref{prop-n4569} ---} \\vskip 3mm ",
        "conclusions": "\\label{sec-discussion} The highest resolution CO map {shows that six molecular blobs are likely to be the CER.} There is a bright H$\\alpha$ core at the center of the CER. The existence of the CER is common and is sometimes found in spiral galaxies (e.g., Kohno et al., 2003). The formation of the CER may be related to the dynamical condition, which we discuss in the following subsection. Moreover, we discuss the existence of molecular blobs and the bright H$\\alpha$ core accompanying the CER.  \\subsection{Formation of the CER} The velocity field (figure \\ref{MOM1}) and the PV diagram (figure \\ref{PV}) show the existence of a strong non-circular motion. \\citet{nak04} showed that the two-dimensional orbits of the ISM of this galaxy are in fact oval. This non-circular motion implies that there is an asymmetric potential (bar structure). Although the optical image of NGC 4569 does not show a strong bar, the non-circular motion can be reproduced even by a weak-bar potential  (e.g., Wada, Habe, 1992; Koda, Wada, 2002). The central gas dynamics is largely related to the radii of the inner Lindlbad resonances (ILRs). Wada and Habe (1992) have shown that the gas forms a CER around ILRs. NGC 4569 might have ILRs and the CER has been formed as a result.   {We compare the observational data with a Smoothed Particle Hydrodynamics (SPH) simulation of a gas disk in a bar potential. Figures \\ref{OBS-SIM} a, b, and c show snapshots of the gas distribution, and velocity field, and a PV diagram from model E of Koda and Wada (2002), respectively. We rotated the gas distribution and the velocity field so that their major axes are aligned with the horizontal axes of figures \\ref{OBS-SIM}a and b using a position angle of $18\\arcdeg$ and an inclination of $63\\arcdeg$. The major axis of the elliptical ring is assumed to be tilted counterclockwise by $31\\arcdeg$ on the plane of the galaxy disk (see subsubsection \\ref{sec-uni}). There is a strong concentration of the gas on a central elliptical ring, and the gas is moving along the ring (figure \\ref{OBS-SIM}a).  Figure \\ref{OBS-SIM}b shows that the global velocity gradient is not parallel to the minor axis of the main disk, which implies that the gas does not circularly rotate. A PV diagram of the CER is shown in figure \\ref{OBS-SIM}c, which shows forbidden velocity components as blue-shifted components in the left-hand side and red-shifted components in the right-hand side. The PV diagram is similar to a parallelogram. As shown in figure \\ref{PV}, the gas in the CER moves along this parallelogram.  Corresponding figures are presented in the right-hand side of figure \\ref{OBS-SIM}; figures \\ref{OBS-SIM}A, B, and C are the gas distribution, velocity field, and PV diagram obtained from observations, respectively.  The observed features for NGC 4569 are well reproduced by the SPH simulation; the observed ring structure in the higher resolution map (figure \\ref{OBS-SIM}A) is similar to the elliptical ring of figure \\ref{OBS-SIM}a. The $S$-shaped twists in the observed velocity field (figure \\ref{OBS-SIM}B) are similar to that of the model (figure \\ref{OBS-SIM}b). The parallelogram-like feature seen in the observed PV diagram (figure \\ref{OBS-SIM}C) is also reproduced by the model (figure \\ref{OBS-SIM}c).}  In the next subsection, we discuss the gaseous self-gravity and the formation of large molecular blobs.  \\vskip 3mm \\centerline{--- Figure \\ref{OBS-SIM} ---} \\vskip 3mm  \\subsection{Large Molecular Blobs and Central Starburst}  Wada and Habe (1992) discussed that the CER fragments and collapses due to the self-gravitational instability of gas when the ratio of the gas mass to the stellar mass is larger than 10\\%. They described the process of decay of the CER as follows: (1) If the CER is dense enough, it fragments and collapses into blobs due to the self-gravity. (2) Then, these gas blobs collide with each other. Kinetic energy is dissipated due to the collisions. (3) Massive blobs are formed by accumulating smaller blobs, losing the kinetic energies. (4) Consequently, the large blobs fall into the center.  The observed results are likely to support this scenario. First, the CER and non-circular motion are consistent with the gas distribution and velocity field caused under a weak bar potential in numerical hydrodynamic simulations. Second, the large molecular blobs observed in this galaxy are also found in simulations. The largest blob A is the closest to the galactic center, which can be understood by the scenario: small blobs grow into a larger one while losing the kinetic energies due to collisions; then, the larger blob falls into the central region because of the loss of kinetic energy. Wada and Habe (1992) discussed that a gravitational instability occurs when the gas-to-stellar mass ratio within a radius is greater than 10\\%. This condition is satisfied in this galaxy as follows. The rotational velocity is $\\sim 130$ km s$^{-1}$ at the radius of $20\\arcsec$, adopting an inclination angle of $63\\arcdeg$ (see figure \\ref{PV}). The dynamical mass within $20\\arcsec$ is calculated to be $6.1 \\times 10^{9} M_\\odot$. Using the equation (\\ref{eq-exp}), the gaseous mass within $20\\arcsec$ was $1.1 \\times 10^{9} M_\\odot$. The gas-to-stellar mass ratio reaches $18\\%$. Therefore, the self-gravity of the gas is important and a self-gravitational instability can occur.   Thus, the observable results support the scenario that the gas becomes clumpy due to its self-gravity, accumulated into the center, and is leading to the central starburst (Wada, Habe, 1992). The central bright H$\\alpha$ core might be caused by the gravitational instability of the CER and consequent gas inflow."
    },
    "0512/astro-ph0512534_arXiv.txt": {
        "abstract": "We present high resolution spectro-astrometry of a sample of 28 Herbig Ae/Be and 3 F-type pre-main sequence stars. The spectro-astrometry, which is essentially the study of unresolved features in long-slit spectra, is shown from both empirical and simulated data to be capable of detecting binary companions that are fainter by up to 6 magnitudes at separations larger than $\\sim$ 0.1 arcsec. The nine targets that were previously known to be a binary are all detected.  In addition, we report the discovery of 6 new binaries and present 5 further possible binaries. The resulting binary fraction is 68$\\pm$11\\%. This overall binary fraction is the largest reported for any observed sample of Herbig Ae/Be stars, presumably because of the exquisite sensitivity of spectro-astrometry for detecting binary systems. The data hint that the binary frequency of the Herbig Be stars is larger than for the Herbig Ae stars. The appendix presents model simulations to assess the capabilities of spectro-astrometry and reinforces the empirical findings.  Most spectro-astrometric signatures in this sample of Herbig Ae/Be stars can be explained by the presence of a binary system. Two objects, HD~87643 and Z CMa, display evidence for asymmetric outflows. Finally, the position angles of the binary systems have been compared with available orientations of the circumprimary disc and these appear to be co-planar. The alignment between the circumprimary discs and the binary systems strongly suggests that the formation of binaries with intermediate mass primaries is due to fragmentation as the alternative, stellar capture, does not naturally predict aligned discs.  The aligment extends to the most massive B-type stars in our sample. This leads us to conclude that formation mechanisms that do result in massive stars, but predict random angles beween the binaries and the circumprimary disks, such as stellar collisions, are also ruled out for the same reason. ",
        "introduction": "In contrast to the low mass stars, which are commonly accepted to form due to the gravitational collapse of a dusty cloud and the subsequent magnetically controlled accretion via a disc (e.g. Bertout 1989), the formation of the most massive stars is still very much uncertain. Part of this is due to the extreme rarity of these objects and the fact that they are very embedded in their molecular clouds, rendering them obscured. To make headway, it is important to study the optically visible Herbig Ae/Be stars as these have masses intermediate between the two above cases.  The Herbig Ae/Be stars are now also found to mark the transition between the different formation mechanisms of low and high mass stars. Recent studies indicate that the Herbig Ae stars present similar phenomena as T Tauri stars (Vink et al. 2003, 2005, see also Hubrig, Sch\\\"oller \\& Yudin 2004), whereas the Herbig Be stars are similar to the more embedded Massive Young Stellar Objects (MYSOs, e.g. Drew et al. 1997).  To learn more about their formation, it is necessary to probe the (circum-)stellar environments of Herbig Ae/Be stars at the small scales where many important features for the further evolution such as discs, outflows, and binaries are found. Millimetre observations have detected discs in the size range of $\\sim$ 200 to 3000 AU around Herbig Ae/Be stars (Mannings \\& Sargent, 1997 \\& 2000; Pi\\'{e}tu, Dutrey \\& Kahane 2003), while near-IR coronagraphic observations of the Herbig Ae star HD 150193 (Fukagawa et al., 2003) and the Herbig Be star HD 100546 (Grady et al., 2001) reveal discs extending to $\\sim$ 190 and 515 AU respectively. Using spectropolarimetry, Oudmaijer \\& Drew (1999) and Vink et al. (2002) found evidence for even smaller scale discs, with sizes of the order of several stellar radii, around Herbig Ae/Be stars.  Pre-main sequence binary fractions have been determined from IR imaging (Bouvier \\& Corporon, 2001; Leinert, Richichi \\& Haas 1997; Pirzkal, Spillar \\& Dyck 1997; and Li et al., 1994) and optical spectroscopy (Corporon \\& Lagrange, 1999).  These studies are based on limited samples of Herbig Ae/Be stars, but from extrapolations of the detections, the inferred binary frequency is almost double that of solar type Main Sequence stars for the period range probed (Leinert et al. 1997), and  of the same order as the T Tauri stars (K\\\"{o}hler \\& Leinert, 1998).  In this paper we apply the technique of spectro-astrometry to study Herbig Ae/Be stars. This method is particularly suited to detect binary companions of emission line stars, and as we will show, can detect secondaries as much as 6 magnitudes fainter than the primary. As such it is among the most powerful methods to investigate binary frequencies, while it has the potential to probe even closer to the star and detect small scale structures in the circumstellar material.  The basic principle of spectro-astrometry is to measure the relative spatial position of spectral features from a longslit spectrum (Bailey, 1998). This is achieved by obtaining high spectral resolution long-slit spectra of an object and determining the centre of the point spread function (PSF) at each wavelength. In his pioneering paper, Bailey (1998) already illustrated that spectro-astrometry is effective at finding binaries and outflow structures around Herbig Ae/Be stars, focusing on the strong H$\\alpha$ emission of these objects. Later studies have found outflows and binary signatures from a sample of T Tauri stars (e.g. Takami, Bailey \\& Chrysostomou 2003 and references therein) and an outflow from a  Brown Dwarf (Whelan et al.  2005).   Returning to the detection of binary stars, a system where only one of the components emits H$\\alpha$ displays a simple and easily detectable signature in the photo-centre of the spectrum. As the spatial profile is the sum of both stars, the peak is not located at the position of either star, but somewhere between the components, depending on the intensity ratio - analogous to the centre of motion of binaries. Because the H$\\alpha$ line intensity ratio changes across the line, the peak position shifts towards the H$\\alpha$ dominant star and the binary is revealed by an excursion in the positional spectrum.  We also expect to see a decrease in the full-width-at-half-maximum (FWHM) of the spatial profile across H$\\alpha$ as it is dominated by only one star. This paper is the first to present and use this additional information that can be straightforwardly retrieved from the data.  We will demonstrate later that the FWHM spectra have a strong diagnostic potential in assessing the properties of the objects under investigation.   In principle the method can be applied across any spectral line in which the intensity ratios of the two stars change from the continuum ratio; in the case of an absorption line the peak position will move away from the star displaying the absorption.  Previously we reported on data of the pre-main sequence binary HK Ori (Baines et al. 2004), the aim of the present work is to further develop the technique and apply it to a large sample of Herbig Ae/Be stars.  We observed 28 young Herbig Ae/Be stars and 3 F-type pre-main sequence stars spectro-astrometrically.  The observations are centred on H$\\alpha$ because the H$\\alpha$ emission line is usually by far the strongest feature in the spectrum of these objects. It is therefore particularly suitable to study the geometry of the circumstellar material, or as in this paper, binarity.  In Section \\ref{Observations} we describe the observations and the targets. The method used to extract information from such data is outlined here. Section \\ref{Results4} contains our results, which is followed by a discussion on our findings in Section 4. Section 5 summarizes the main conclusions and in the Appendix we discuss the strengths and limits of spectro-astrometry by simulating data.    \\begin{sidewaystable*} \\centering  \\begin{minipage}{240mm} {{\\bf Table 1.} Log of the observations. Columns 2 to 4 list the RA, Dec and spectral type of the objects respectively. The $V$ magnitudes in column 5 are taken from SIMBAD. Column 6 lists the dates of the observations and column 7 the exposure times. The average seeing for each object is given in column 8 and the number of photons per pixel in the stellar continuum in column 9. The 1$\\sigma$ detection limit of both the position spectra and FWHM spectra are given in columns 10 and 11 respectively. These are the root-mean-square variations measured in line-free regions in the stellar continuum. Properties of the lines are listed in the final columns. The emission line profiles are classified as doublepeaked (D), singly peaked (S) or as showing (inverse) P Cygni (P) profiles. The equivalent width, with estimated errors of 10\\%, is also listed.  References: for the spectral types: 1. Mora et. al. (2001), 2. SIMBAD, 3. Th\\'e et. al. (1994). \\label{log} }  \\begin{tabular}{lcclrccrrrrrrrcc} \\hline Object    &     RA     &     DEC     &     Spec.                    &{\\it V}&  Date    & t$_{exp}$& Seeing & Photons &    rms$_{pos}$ &     rms$_{FWHM}$ &     H$\\alpha$ W$_\\lambda$ &     H$\\alpha$ \\\\ &     (2000) &     (2000)  &     type                      &      &  &    (s) &     (arcsec) &     (x10$^4$) &     (mas) &     (mas) &     ($\\rm \\AA$) &     Profile \\\\  \\hline XY Per    & 03 49 36.3 & +38 58 55.5 & A2{\\sc i}{\\sc i}$^2$      & 9.4  & 21-09-02 & 6 x 3000 & 2.2 &  6.7 & 2.0 &  4.3 & -6.6 & D   \\\\ ''       & ''         & ''          & ''                        & ''   & 23-09-02 & 8 x 900  & 2.1 &  6.1 & 4.6 & 21.0 & -4.5 & D   \\\\ AB Aur    & 04 55 45.8 & +30 33 04.3 & A0$^2$                    & 7.1  & 29-01-02 & 4 x 180  & 2.4 & 16.0 & 2.5 &  4.5 & -45  & P   \\\\%O, no Fe\\\\ ''       &  ''        & ''          & ''                        & ''   & 19-09-02 & 8 x 510  & 1.8 & 32.0 & 0.9 &  1.7 & -21  & P   \\\\%O, no Fe\\\\ HD 31648  & 04 58 46.3 & +29 50 37.0 & A5{\\sc v}$^1$             & 7.7  & 29-01-02 & 4 x 300  & 2.6 & 16.5 & 2.3 &  5.0 & -17  & D   \\\\ ''       & ''         & ''          &  ''                       & ''   & 19-09-02 & 4 x 1440 & 1.6 & 27.0 & 1.1 &  2.7 & -14  & D/P \\\\%no O,invP Fe\\\\ HK Ori    & 05 31 28.0 & +12 09 10.3 & A4/G1{\\sc v}$^1$          & 11.9 & 29-01-02 & 8 x 300  & 2.4 &  1.6 & 7.7 & 19.2 & -47  & D   \\\\%O, Fe\\\\ HD 244604 & 05 31 57.3 & +11 17 41   & A3$^2$                    & 9.4  & 28-01-02 & 4 x 600  & 1.9 & 10.5 & 2.1 &  4.5 & -1.7 & D   \\\\ T Ori     & 05 35 50.4 & -05 28 35   & A3{\\sc i}{\\sc v}$^1$      & 9.5  & 28-01-02 & 4 x 600  & 1.6 &  6.7 & 2.1 &  5.3 & -8   & D   \\\\%O, no Fe\\\\ ''       & ''         & ''          & ''                        & ''   & 29-01-02 & 4 x 600  & 2.0 &  4.9 & 2.6 &  6.5 & -8   & D   \\\\%O, no Fe\\\\ CQ Tau    & 05 35 58.5 & +24 44 54.1 & F5{\\sc i}{\\sc v}$^1$      & 10.7 & 23-09-02 & 4 x 2400 & 1.8 &  4.7 & 3.4 &  8.2 & -1.6 & invP\\\\%\\\\ V380 Ori  & 05 36 25.4 & -06 42 57.7 & A0$^2$                    & 10.0 & 29-01-02 & 4 x 300  & 2.5 &  1.4 & 8.3 & 19.1 & -80  & S/D \\\\%O, strong Fe\\\\ MWC 137   & 06 18 45.5 & +15 16 52.4 & B0$^3$                    & 11.2 & 22-09-02 & 8 x 600  & 1.2 &  1.7 & 2.3 & 11.2 & -370 & S   \\\\%O, Fe\\\\ HD 45677  & 06 28 17.4 & -13 03 11.1 & B3$^3$                    & 8.0  & 29-01-02 & 8 x 300  & 2.8 & 13.8 & 2.8 &  6.8 & -57  & D   \\\\%O, Fe\\\\ MWC 147   & 06 33 05.2 & +10 19 20.0 & B6$^2$                    & 8.8  & 28-01-02 & 4 x 600  & 2.1 & 20.5 & 1.7 &  4.2 & -56  & D   \\\\%O, weak Fe\\\\ MWC 158   & 06 51 33.4 & -06 57 59.4 & B9$^2$                    & 6.6  & 28-01-02 & 8 x 45   & 1.6 & 22.0 & 1.3 &  3.2 &  -71 & D   \\\\%O, no Fe\\\\ ''       & ''         & ''          & ''                        & ''   & 20-09-02 & 24 x 120 & 1.4 & 36.0 & 0.9 &  1.7 &  -77 & D   \\\\%O, Fe?invP?\\\\ GU CMa    & 07 01 49.5 & -11 18 03.3 & B2{\\sc v}$^2$             & 6.6  & 28-01-02 & 4 x 300  & 1.6 & 85.0 & 1.4 &  1.7 &  -7  & D   \\\\ Z CMa     & 07 03 43.2 & -11 33 06.2 & B5$^3$                    & 9.9  & 28-01-02 & 4 x 600  & 1.8 & 11.0 & 1.5 &  4.7 & -68  & D   \\\\%O, strong Fe\\\\ MWC 166   & 07 04 25.5 & -10 27 15.7 & B0{\\sc i}{\\sc v}$^2$      & 7.0  & 28-01-02 & 4 x 240  & 1.8 & 33.0 & 1.1 &  2.4 & -1   & D   \\\\ HD 58647  & 07 25 56.1 & -14 10 43.6 & B9{\\sc i}{\\sc v}$^2$      & 6.8  & 29-01-02 & 4 x 180  & 2.6 & 23.5 & 2.0 &  5.3 & -6.0 & D   \\\\ HD 85567  & 09 50 28.5 & -60 58 03.0 & B5{\\sc v}$^3$             & 8.6  & 29-01-02 & 4 x 450  & 2.1 & 16.5 & 2.0 &  4.8 & -40  & S   \\\\%O, Fe\\\\ HD 87643  & 10 04 30.3 & -58 39 52.1 & B2$^2$                    & 8.9  & 28-01-02 & 8 x 120  & 1.2 & 10.0 & 1.2 &  3.0 & -145 & P,D \\\\%O, strong Fe\\\\ HD 95881  & 11 01 57.6 & -71 30 48.4 & A1/2{\\sc iii}/{\\sc iv}$^2$&8.3   & 29-01-02 & 4 x 300  & 1.9 & 15.0 & 2.0 &  3.9 & -11.5& D   \\\\ HD 97048  & 11 08 03.3 & -77 39 17.5 & A0$^2$                    & 8.5  & 28-01-02 & 4 x 450  & 1.7 & 22.7 & 1.4 &  3.0 & -25  & S   \\\\%O, no Fe\\\\ HD 98922  & 11 22 31.7 & -53 22 11.5 & B9{\\sc v}$^2$             & 6.8  & 29-01-02 & 4 x 120  & 1.8 & 24.7 & 1.2 &  3.2 & -17.5& P   \\\\ HD 100546 & 11 33 25.4 & -70 11 41.2 & B9{\\sc v}$^2$             & 6.7  & 28-01-02 & 4 x 60   & 1.4 & 15.3 & 1.4 &  3.0 & -35  & S   \\\\%O, no Fe\\\\ HD 104237 & 12 00 05.1 & -78 11 34.6 & A4$^3$                    & 6.6  & 29-01-02 & 4 x 120  & 2.0 & 25.4 & 1.6 &  3.6 & -26  & D/P \\\\ HD 135344 & 15 15 48.4 & -37 09 16.0 & A0/F4{\\sc v}$^3$          & 8.6  & 29-01-02 & 4 x 600  & 1.7 & 25.4 & 1.3 &  3.0 & -6.5 & S/D \\\\ HD 142527 & 15 56 41.9 & -42 19 23.3 & F6{\\sc i}{\\sc ii}$^2$     & 8.3  & 29-01-02 & 4 x 450  & 1.7 & 30.0 & 1.2 &  3.0 & -5.8 & P   \\\\ 51 Oph    & 17 31 25.0 & -23 57 45.5 & B9.5{\\sc iii}$^1$         & 4.8  & 29-01-02 & 4 x 30   & 2.0 & 28.0 & 1.4 &  3.5 & +4   & D   \\\\%no O, weak Fe?\\\\ HD 190073 & 20 03 02.5 & +05 44 16.7 & A2{\\sc i}{\\sc v}$^1$      & 7.8  & 21-09-02 & 8 x 1250 & 2.0 & 30.0 & 1.1 &  2.7 & -24  & P   \\\\%O, Fe\\\\ MWC 361   & 21 01 36.9 & +68 09 47.8 & B2{\\sc v}$^2$             & 7.4  & 19-09-02 & 16 x 500 & 1.3 & 49.0 & 0.9 &  1.9 & -59  & D   \\\\%O?, Fe\\\\ BH Cep    & 22 01 42.9 & +69 44 36.5 & F5{\\sc i}{\\sc v}$^2$      & 10.8 & 22-09-02 & 4 x 2700 & 1.5 &  3.6 & 2.7 &  7.0 & 0.0  & D/T \\\\ SV Cep    & 22 21 33.3 & +73 40 18   & A0$^2$                    & 10.1 & 22-09-02 & 4 x 1800 & 1.7 &  3.3 & 3.8 &  7.2 & +1.2 & D   \\\\ MWC 1080  & 23 17 26.1 & +60 50 43   & B0$^2$                    & 11.6 & 23-09-02 & 4 x 2400 & 1.8 &  2.1 & 4.5 &  8.7 & -140 & P/D \\\\%Strong\\\\ \\hline \\end{tabular} \\hspace*{\\fill} \\end{minipage} \\end{sidewaystable*} ",
        "conclusions": "We have studied 31 young stars by means of spectro-astrometry, and find the following.   \\begin{enumerate}  \\item A binary frequency of 68$\\pm$11\\% was found in the sample of 28 Herbig Ae/Be stars.  No other technique has been able to find such a high binary frequency in Herbig Ae/Be stars, demonstrating the proficiency at which spectro-astrometry can detect binaries.  \\item There is a small hint that the Herbig Be stars are more likely to be found in binaries than the Herbig Ae objects (73$^{+12}_{-16}$\\% versus 31$^{+18}_{-14}$\\%, the latter rises to 62$^{+15}_{-17}$\\% when the possible detections are included). This could indicate that Herbig Be stars are more likely to be found in binaries than their lower mass Herbig Ae counterparts.   \\item We  detected asymmetric outflows from Z CMa and HD 87643.  We conclude that HD 87643 has a high velocity and low velocity outflow, towards the north. No evidence for jets was present in our data.  \\item We present evidence that the circumprimary discs are aligned with the binary systems. This favours the fragmentation scenario for the formation of Herbig Ae/Be binaries over stellar capture models which predict randomly oriented discs. The next step is to identify the mode of fragmentation, i.e. whether the clouds break up in 2 fragments (predicting roughly fixed mass ratios) or multiple fragments, predicting random mass functions.  \\item The strong alignment, even for early B type stars, provides evidence against the stellar merger scenario invoked to produce massive stars. Instead, as these objects have been found to have discs in the first place, disc accretion remains a viable possibility (cf. Norberg \\& Maeder 2000).   \\end{enumerate}"
    },
    "0512/astro-ph0512228_arXiv.txt": {
        "abstract": "{In order to study the circumnuclear star-forming activity along the Hubble sequence, we cross-correlate the Sloan Digital Sky Survey Data Release 2 (SDSS DR2) with the Third Reference Catalog of Bright Galaxies (RC3) to derive a large sample of 1015 galaxies with both morphological and spectral information. Among these, 385 sources are classified as star-forming galaxies and the SDSS fibre covered the circumnuclear regions (0.2 $-$ 2.0 kpc). By using the spectral synthesis method to remove the contribution from the underlying old stellar population, we measure the emission lines fluxes accurately which are then used to estimate the star formation rates(SFRs). Our main findings are that: (1) Early-type spirals show much larger H$\\alpha$ luminosities and hence higher SFRs, they also suffer more extinctions than late-type ones. The equivalent widths (EWs) of H$\\alpha$ emission lines show the similar trend, however, the very late types (Sdm $\\sim$ Irr) do have large fractions of high EWs. (2) We confirm that D$_n(4000)$ has strong correlation with the strengthes of metallic absorption lines (such as CN band, G band and Mg Ib). Both these lines and the Balmer absorption lines show interesting variations between Sbc to Sd type galaxies. (3) The bar structure tightly relates with the enhanced star formation activity, this effect is even more significant in the early-type spirals. But we should note that the bar structure is not a necessary or sufficient condition for galaxies to harbor circumnuclear star formations. ",
        "introduction": "The star formation history is one of the most important problems for understanding the evolutions of the galaxies and also the universe. Enormous progress have been made during the last two decades. More precise diagnostics of global star formation rates (SFRs) have been obtained in a broad range of techniques: e.g., integrated optical emission line fluxes (Kennicutt 1983; Kennicutt et al. 1994; Madau et al. 1998), near-ultraviolet continuum fluxes (Donas \\& Deharveng 1984; Deharveng et al. 1994; Madau et al. 1998; Bell \\& Kennicutt 2001), infrared continuum fluxes (Harper \\& Low 1973; Rieke \\& Lebofsky 1978, Telesco \\& Harper 1980; Kennicutt et al. 1987), and radio emissions (Condon 1992; Cram et al. 1998). Kennicutt (1998) has presented an excellent review on star formation along the Hubble sequence.  A galaxy's morphology carries an important information for the study of star formation history in galaxies (e.g., Kennicutt 1998). Hubble sequence (Hubble 1926), as one of the most widely used classification of galaxies, is closely connected with the global star formation activities of galaxies. The basic view of global SFRs along the Hubble sequence shows that the equivalent widths (EWs) of H$\\alpha$ emission lines increase from virtually zero in E/S0 galaxies to about a couple of hundreds \\AA\\ in late-type spiral and irregular galaxies (Kennicutt \\& Kent 1983; Kennicutt 1998). However, recent studies of H$\\alpha$ imaging indicate that a significant fraction of the early-type spirals shows remarkable star formation activities ( Young et al. 1996; James et al. 2004; Hameed \\& Devereux 2005).  Many galaxies are shown to harbor luminous star forming activities at the circumnuclear region with properties different from those of extended disks (Morgan 1958, S\\'ersic \\& Pastoriza 1967). Comprehensive surveys of the star formation properties of galactic nuclei have been carried out both in the optical band (Stauffer 1982, Keel 1983, Kennicutt et al. 1989, Ho et al. 1997a, b) and mid-IR photometry (Rieke \\& Lebofsky 1978, Scoville et al. 1983, Devereux et al. 1987, Giuricin et al. 1994). Ho et al. (1997a) found that the fraction of nuclear emission spectra with HII region-like line ratios increased from virtually zero in elliptical galaxies, 8\\% in S0 galaxies to 80\\% in Sc-Im galaxies. However, these could be contaminated by LINER or Seyfert nuclei.  In most galaxies, the nuclear SFRs are quite modest, while the highest SFRs are mainly seen in the IR observations because the very luminous starbursts are mostly associated with dense molecular gas and suffered substantial extinctions in the optical band The details of these IR luminous circumnuclear properties could be seen in Veilleux et al. (1995), Lutz et al. (1996), and also in Kennicutt (1998).  In this paper, we cross-correlate the largest spectral database available from the Sloan Digital Sky Survey (SDSS) and the Third Reference Catalog of Bright Galaxies (RC3) to derive a large sample of galaxies with both morphological and spectral information, we thus are able to statistically study the star forming activities along the Hubble sequence. This paper is organized as follows: we will describe the sample selection and also give a brief introduction of the stellar synthesis model in Section 2; the statistical results are presented in Section 3; in Section 4, we will discuss bar influence on the star formation activities; finally, we present our main conclusions in Section 5. ",
        "conclusions": "We have constructed a sample of galaxies by positionally matching up the SDSS and RC3 data. We then applied the stellar population synthesis model to statistically study the variations of the star formation histories and line features along the Hubble types. Our finding are summarized as follows: \\begin{enumerate} \\item A decreasing sequence is found for nebular extinction from early-type spirals to late-type ones, which is consistent with previous studies of Ho et al. (1997a) and Stasi\\'nska et al. (2004). We have also confirmed that nebular emissions do suffer more serious extinctions than stellar light and the later one shows almost no variations among the Hubble types. \\item We find that H$\\alpha$ luminosities of early-type spiral galaxies are much higher than the late-type ones, which causes the corresponding larger SFRs in early-type spirals. From the equivalent widths of H$\\alpha$, we find in general, early-type spirals show higher value of EWs, however, the dramatic change is found from Sdm type. We also plot the FIR-SFRs as a function of Hubble types and found that the early-type spirals show higher amount of FIR-SFRs and may partly due to the dust re-emission which will favor the early types. However, this difference of SFRs in IR band is comparably much smaller. Although large amount of galaxies show quite modest SFRs and may not actually contain bright HII regions, those who have strongly favor the higher star-forming activities in early-type galaxies. \\item We have measured the spectral absorption features and confirmed previous findings that tight correlations between D$_n$(4000) and equivalent widths of absorption lines: CN band, G-band, Mg Ib, H$\\delta$, H$\\gamma$, H$\\beta$ and H$\\alpha$. The dependence of mean EWs on the Hubble types could be seen in metallic absorption lines. We also find an increase sequence of mean Balmer absorption equivalent widths from Sbc to Sd type. Since all these EWs are measured automatically, they do not depend on the morphology types. \\item We have discussed the bar effects on the star forming activities on both early-type and late-type spirals, and confirmed previous findings that bar structures can enhance the star forming activities, especially in early type spiral galaxies. However, bars are not necessarily inducing strong star formation activities; and those without could still contain high star formation regions. We also find that, though in general, the barred early-type spirals show larger EW[H$\\alpha$], the largest difference of EW[H$\\alpha$] between barred and unbarred galaxies do exist in the very late-type spirals (Sd $\\sim$ Sm). Thus we believe that other effects like interactions and gas contents may also play important roles in the star formation activities. \\end{enumerate}"
    },
    "0512/astro-ph0512191_arXiv.txt": {
        "abstract": "{We present preliminary results from the first part of the LuckyCam late M-dwarf binarity survey. We survey a sample of 48 nearby ($<$40 pc) and red (M5-M9) stars with the novel high angular resolution visible light imaging technique Lucky Imaging, in only 8 hours of 2.5m telescope time. We discover 10 new binaries; although the survey is sensitive to brown dwarf companions none are detected. The orbital radius distribution of the newly discovered binaries broadly matches that of previous detections by other groups, although we do discover one wide binary at $\\sim$40AU. ",
        "introduction": "This short paper presents preliminary results from our \\mbox{LuckyCam} high angular resolution binarity survey of nearby ($<$40 pc) and red (M5-M9) stars. LuckyCam is a novel imaging system designed and built by our group that gives near diffraction limited images in the visible with 2.5m class telescopes. ",
        "conclusions": "Because the sample presented in this paper is inhomogeneous we feel it would be premature to calculate a binary fraction on the basis of our detections; this will be addressed in further work.  Although we are sensitive to very faint companions (i'=20 magn. at 1 arcsec, $\\Delta$m=5, in the brown dwarf regime), none of the detected companions are more than 2 magnitudes fainter than the primary. This suggests that, as has been found by previous authors, brown-dwarf companions to M-dwarfs are very rare at separations wider than a few AU.  The distribution of orbital radii (figure \\ref{FIG:sep_hist}) is consistent with that found in previous M-dwarf surveys, with a peak at around 4 AU and no companions found at larger radii than 40AU. However, we note that two of our new binaries fall near a previous single detection at about 35 AU. It is not yet clear  whether these wider systems are in the high separation tail of the  orbital radius distribution, or part of a distinct population.   \\begin{figure} \\centering \\resizebox{\\columnwidth}{!} { \\includegraphics{sep_hist.eps} } \\caption{\\footnotesize The distribution of orbital radii of the new binaries compared to that of previously known VLM binary systems collated in Siegler et al. 2005.} \\label{FIG:sep_hist} \\end{figure}   \\begin{figure} \\centering \\subfigure{\\includegraphics[width=0.40\\columnwidth]{1_z.eps}} \\subfigure{\\includegraphics[width=0.40\\columnwidth]{2_z.eps}} \\subfigure{\\includegraphics[width=0.40\\columnwidth]{3_z.eps}} \\subfigure{\\includegraphics[width=0.40\\columnwidth]{4_z.eps}} \\caption{\\footnotesize SDSS z' images of some of the discovered binaries; 10\\% selection (1\\% for top left) from 3000 frames.} \\label{FIG:binaries} \\end{figure}"
    },
    "0512/astro-ph0512472_arXiv.txt": {
        "abstract": "{}{ Following an earlier proposal by Choudhuri (2003) for the origin of twist in the magnetic fields of solar active regions, we model the penetration of a wrapped up background poloidal field into a  toroidal magnetic flux tube rising through the solar convective zone. }{ The rise of the straight, cylindrical flux tube is followed by numerically solving the induction equation in a comoving Lagrangian frame, while an external poloidal magnetic field is assumed to be radially advected onto the tube with a speed corresponding to the rise velocity. }{ One prediction of our model is the existence of a ring of reverse current helicity on the periphery of active regions. On the other hand, the amplitude of the resulting twist depends sensitively on the assumed structure (diffuse vs. concentrated/intermittent) of the active region magnetic field right before its emergence, and on the assumed vertical profile of the poloidal field. Nevertheless,  in the model with the most plausible choice of assumptions a mean twist comparable to the observations results. }{ Our results indicate that the contribution  of this mechanism to the twist can be quite significant, and under favourable circumstances it can potentially account for most of the current helicity observed in active regions. } ",
        "introduction": "The magnetic field of a typical sunspot is approximately vertical below the photosphere and then spreads radially at the photospheric level.  However, even visual observations show that the penumbral structures of sunspots are often twisted, as first pointed out by \\citet{Hale:vortices} and \\citet{Richardson:vortices}.  Vector magnetogram measurements in the past decade  have established that sunspot magnetic fields have helical structures, with a higher occurrence of negative helicity in the northern hemisphere (\\citen{Seehafer:helicity}; \\citen{Pevtsov1995}, \\citeyear{Pevtsov2001}; \\citen{Bao+Zhang:helicity}). One possible theoretical explanation of the observed helicity was proposed by \\citet{Choudhuri:tube.dynamo}, who suggested that the poloidal flux in the solar convection zone (SCZ) gets wrapped around a rising flux tube.   \\citet{Choudhuri+:helicity} carried out a numerical simulation demonstrating that the observed hemispheric handedness rule can indeed be reproduced by incorporating this mechanism in their dynamo model (\\citen{Nandy+Choudhuri2002}; \\citen{Chatterjee2004}; \\citen{Choudhuri+:reply}).  If the magnetic flux in the rising flux tube is nearly frozen, then we expect that the poloidal flux collected by it during its rise through the SCZ would be confined in a narrow sheath at its outer periphery. In order to produce a twist in the flux tube, the poloidal field needs to diffuse from the sheath into the tube by turbulent diffusion. However, turbulent diffusion is strongly suppressed by the magnetic field in the tube. This nonlinear diffusion process was studied in an untwisted flux tube by  \\citet{Petrovay+FMI:erosion} who concluded that a substantial amount of flux may be eroded away from a rising flux tube during the process of its rise through the SCZ.  The model was subsequently successfully applied for sunspot decay (\\citen{Petrovay+vDG:decay1}). In the present paper we extend this model by including the poloidal component of the magnetic field (i.e.\\ the field which gets wrapped around the flux tube) and study the evolution of the magnetic field in the rising flux tube, as it keeps collecting more poloidal flux during its rise and as turbulent diffusion keeps acting on it.  The conclusions drawn from our model hinge on some assumptions, especially concerning the subsurface magnetic field structure in the last phases of the rise of the tube.  Allowing the various parameters of the model to vary over reasonable ranges, we find that the poloidal field wrapped around the flux tube should be able to penetrate inside the flux tube if the magnetic field falls below the equipartition value in the top layers of the convection zone. On the other hand, if some physical effects keep the magnetic field well above the equipartition value, then the poloidal field may remain confined in a sheath instead of penetrating to the core of the flux tube.  We expect that much more high-quality magnetogram data will be available in the future and more will be known as to how the current helicity in a sunspot and in an active region varies with radial distance from the centre.  This may enable detailed comparisons between theory and observations in future, with the possibility of constraining various parameters in the theory.  Our aim at present is only to set up the basic theoretical framework and study some exploratory solutions.  Section~2 presents the mathematical formulation of our problem, while the numerical solutions are presented and discussed in Section~3. Section~4 concludes the paper. ",
        "conclusions": "Alternative mechanisms for the origin of twist in active region magnetic fields include buffeting of the rising flux tubes by helical turbulent motions (\\citen{Longcope+:Sigma.effect}); the effect of Coriolis force on flows in rising flux loops (\\citen{Fan+Gong:twist}); differential rotation (\\citen{DeVore:diff.rot.helicity}); and helicity generation by the solar dynamo (\\citen{Seehafer+:AdvSpRes}). The relative importance of these processes is currently a subject of debate {\\changes (\\citen{Holder:helicity}). } It is likely that the accretion of poloidal fields during the rise of a flux tube is just one contribution to the development of twist. Its importance may also be reduced by 3D effects: considering the rise of a finite flux loop instead of an infinite horizontal tube, the possibility exists for the poloidal field to ``open up'', giving way to the rising loop with less flux being wrapped around it. It is left for later multidimensional analyses of this problem to determine the importance of any such reduction. In any case, the results presented above indicate that the contribution of poloidal field accretion to the development of twist can be quite significant, and under favourable circumstances it can potentially account for most of the current helicity observed in active regions.  In our calculations we found that while the flux tube rises to a depth of about $0.85 \\Rs$, the effect of diffusion is small and the poloidal field remains confined in a narrow sheath at the periphery of the flux tube.  Afterwards, if the magnetic field is allowed to fall to very low values as suggested by simple flux tube rise simulations (case A), then the effect of diffusion is considerable and the poloidal field is able to penetrate into the interior of the flux tube. On the other hand, if various effects in the top layers of SCZ keep the magnetic field above $\\Beq $ (case B1) and $3\\Beq$ (case B3), then diffusion is less effective; in case B3 the poloidal field remains confined in the sheath.  For a poloidal field strength independent of depth, as assumed in these calculations, the best agreement with the observed current helicity values is found for case B3.  One rather strong prediction of our model is the existence of a ring of reverse current helicity on the periphery of active regions. On the other hand,  the amplitude of the resulting twist (as measured by the mean current helicity in the inner parts of the active region) depends sensitively on the assumed structure (diffuse vs. concentrated/intermittent) of the active region magnetic field right before its emergence, and on the assumed vertical profile of the poloidal field. Nevertheless, a mean twist comparable to the observations can result rather naturally in the model with perhaps the most plausible choice of assumptions (case B3).  Thus, by studying the distribution of the azimuthal magnetic field in sunspots and active regions, and by simultaneously studying the subsurface magnetic field structure in emerging active regions by means of local helioseismology, it may be possible to test the present model and to throw some light on the conditions prevailing during the last phase of flux tube rise through the top layers of SCZ.   We hope that magnetogram data will improve significantly in the next few years and it will be possible to draw meaningful inferences. Since at present we have very poor understanding of the nature of the magnetic field or the effect of turbulence when a flux tube rises through layers immediately below the solar surface, indirect inferences from such observations of sunspots are of great importance.  {\\changes Lastly, we have used a simplifying assumption that the poloidal field that gets wrapped around the flux tube has a constant value of 1 G inside the SCZ.  Dynamo models like the model of Choudhuri et al.\\ (2005) can be used to calculate the distribution of the poloidal field inside the SCZ during different phases of the solar cycle.  More detailed calculations of helicity can be done by using such poloidal field distributions instead of using the simple boundary condition (24).  Since some groups have started reporting on the possible cycle variation of helicity on the basis of the observational data (\\citen{Bao:helicity}; \\citen{Hagino:helicity}), such calculations may become relevant in future when more detailed observational data are available. }"
    },
    "0512/astro-ph0512644_arXiv.txt": {
        "abstract": "We study gravitational lensing with a multiple lens plane approach, proposing a simple analytical model for the probability distribution function (PDF) of the dark matter convergence, $\\kappa$, for the different lens planes in a given cosmology as a function of redshift and smoothing angle, $\\theta$. The model is fixed solely by the variance of $\\kappa$, which in turn is fixed by the amplitude of the power spectrum, $\\sigma_8$.  We test the PDF against a high resolution Tree-Particle-Mesh  simulation and find that it is far superior to the Gaussian or the lognormal, especially for small values of $\\theta \\ll 1'$ and at large values of $\\kappa$ relevant to strong lensing. With this model, we predict the probabilities of strong lensing by a single plane or by multiple planes. We find that for $\\theta \\sim 10''$, a single plane accounts for almost all ($\\sim 98\\%$) of the strong lensing cases for source redshift unity. However, for a more typical source redshift of $4$, about $12\\%$ of the strong lensing cases will result from the contribution of a secondary clump of matter along the line of sight, introducing a systematic error in the determination of the surface density of clusters, typically overestimating it by about $2-5\\%$.  We also find that matter inhomogenieties introduce a dispersion in the value of the angular diameter distance about its cosmological mean. The probable error relative to the mean increases with redshift to a value of about $8\\%$ for $z\\simeq 6$ and $\\theta \\sim 10''$. ",
        "introduction": "{\\rm }  Gravitational lensing provides a unique tool for studying the cosmological distribution of matter, because it directly probes the gravitational potential and is indifferent to the nature or the physical state of matter. It is, therefore, free from the assumptions that can plague other dynamical methods for studying gravitational fields.\\par Tidal deflections of light rays by structures along its path can weakly distort and magnify distant sources. These effects, although small on the individual level, can be statistically significant for an ensemble of sources. Thus, statistical properties of the intervening matter distribution can be inferred from such {\\it weak lensing} studies {\\citep[see, for e.g.][for reviews]{ blandford_narayan_1992, bartelmann_schneider_2001, refregier_2003}}. This field is extremely promising at the present time but is fraught with observational challenges, because these subtle effects have to be very delicately isolated from comparable systematic errors.\\par Occasionally, the distortion due to the intervening matter may be strong enough to produce multiple images of the source. {\\it Strong lensing} is extremely sensitive to the cosmology and hence, can be used to constrain cosmological parameters. Small changes in the cosmological parameters can significantly alter the nonlinear structure formation history of the universe and thereby result in huge differences in the abundance of compact objects which behave as potential strong lenses. This makes the probability of any specific type of strong lensing event a useful observable for constraining cosmology {\\citep[see, for example,][]{TOG,narayan_blandford_1991,kaiser_1992,CGOT,nakamura_suto_1997,cohn_etal_2001,keeton_2001,premadi_2001,oguri_etal_2002,li_ostriker_2002,li_ostriker_2003}}. For example, the statistics of lensed quasars or galaxies lensed by clusters into giant arcs can be used to study the population of lenses, their distribution in redshift space, mass, and hence the cosmological model. Strong lensing systems may also be studied individually, in order to obtain a detailed understanding of the structure and the distribution of matter in the lensing object and the source {\\citep[see, for example,][]{claeskens_etal_2000,broadhurst_etal_2000,shapiro_iliev_2000, cohn_etal_2001, belokurov_etal_2001, patnaik_narasimha_2001,inada_etal_2003, fassnacht_etal_2005, sluse_etal_2005, wayth_etal_2005, claeskens_etal_2006}}. In this paper, we shall be mostly interested in the former problem, namely the statistics of strong lensing events.\\par In the absence of satisfactory analytic tools, the most widely used method for studying strong lensing statistics has been the analysis of large numerical simulations. The method of tracing the path of light rays through a simulated universe evolving in accordance with some specific cosmological model is being extensively used to study the multitude of lensing phenomena. For example, giant arc statistics using ray tracing simulations of clusters in a $\\Lambda$CDM  cosmology are being thoroughly investigated \\citep[examples include][etc.]{ bartelmann_etal_1998, WBO_2004a,dalal_holder_hennawi_2004,ho_white_2004,li_etal_2005,hennawi_etal_2005a,hennawi_etal_2005b}.  However, these simulations are costly in that they require extensive computing power and time. This also restricts the range of cosmological models that can be investigated; in most cases only the currently favored cosmological model is used. One must not, however overlook the array of important analytical works from the so-called  halo model approach \\citep[see, for e.g.][]{li_ostriker_2003,oguri_etal_2003, oguri_keeton_2004}, which allow accurate computation of lensing probabilities for a single lens.  \\par The goal of this paper is to provide an analytic method to address some of the questions pertaining to gravitational lensing statistics that, at the cost of having much more limited information available than from detailed numerical studies, would have the benefit of allowing any definite cosmological model to be investigated. We test our results against numerical simulations performed by the Tree-Particle-Mesh (TPM) code \\citep{bode_ostriker_2003} for a $\\Lambda$CDM universe. The method lends itself easily to the treatment of lensing under other cosmologies.\\par We consider the fluctuations in the matter density along a light cone of a given  opening angle.  We propose an analytic method to compute the probability that  a light beam will encounter a given surface density while passing through a slice of matter centered around a given redshift. The model is completely defined by the variance of the projected density of the matter in the slice, smoothed on the angular scale corresponding to the beam. We show how the variance of surface mass density on the lens can be predicted from the power spectrum in Section~\\ref{sec:variances} and develop the analytic model in Section~\\ref{sec:model}.  Armed with this result, we adopt a multiple lens plane approach to study strong lensing. We divide up the line of sight distance to the source into contiguous slices of equal comoving depth, each of which is chosen to be larger than the correlation length so that each slice can be considered statistically independent of the others. We imagine the matter inside each of these slices to be projected onto a plane at its center. The probability density function (PDF) for the surface density on each plane is then constructed with our analytical method. The matter encountered by the light ray while propagating through the universe is then simulated by randomly picking a surface mass density from the PDF on each slice it passes through. In fact, the surface density scaled by the so called critical density required for multiple image formation, better known as the convergence, is computed on each of these planes. We identify different strong lensing scenarios according as whether the individual or the additive value of convergences so obtained is in excess of unity. We shoot a statistically significant number of rays to obtain the probabilities of different events. In particular, we study the probability of strong lensing as a function of source redshift and angular scale.\\par Another important question is the role of auxiliary planes aiding the main lens to produce multiple images. It is known that some multiply imaged quasars \\citep[e.g. the quadrupole system B1422+231, see][]{kormann_etal_1994} cannot be well modeled under the thin lens approximation, which assumes all mass responsible for lensing to be concentrated on a mathematical surface of zero thickness \\citep{SEF_1992}. Also, observations suggest cases where galaxies at disparate redshifts appear to have contributed to the lensing event \\citep{tonry_1998, augusto_etal_2001, chae_etal_2001,johnston_etal_2003,fassnacht_etal_2005}. In order to study the probabilities of such events, we identified, within the strong lensing cases, the ones that were caused by a single plane or by the concordance of two or more planes and studied their probabilities separately. Incidentally, a similar study was done in \\citet{WBO_2004b}, who studied the phenomenon at the pixel level ($\\sim 1.5''$) and reported a significant contribution of auxiliary planes in the strong lensing probabilities. With our method, we can conveniently include an element of angular smoothing depending on the strong lensing situation and study the probabilities as a function of smoothing radius. We treat this problem in Section~\\ref{sec:prob}. \\par We also note that our model works very well even in the weak lensing regime (angles $\\ga 1'$). The statistics of weak lensing convergence can be studied from the  magnification effects of clustered matter which leads to variations in the number density and image sizes of galaxies across the sky \\citep[see, e.g.,][]{jain_etal_2002}. Also, shear maps generated from weak lensing surveys yield statistical properties of the convergence. Thus a good analytic model for the convergence can be useful in constraining the properties of large scale structure \\citep{barber_etal_2004}, although the practicalities involved may seem rather daunting at the moment \\citep{bartelmann_schneider_2001}. The situation will improve with forthcoming surveys having high signal-to-noise ratios, enabling us to study the small scale clustering of matter.\\par In Section~\\ref{sec:mass}, we study the effect of the chance accumulation of secondary matter along the line-of-sight on the estimation of cluster masses and finally, in Section~\\ref{sec:angdia}, we examine the dispersion in the angular diameter distances due to inhomogeneous distribution of matter. ",
        "conclusions": "In the present paper, we have studied gravitational lensing with a multiple lens plane approach and have proposed a new analytic model for the dark matter convergence on the lens planes. Like the lognormal, this model is entirely defined by $\\av{\\k^2}_{\\theta}$, the variance of $\\k$ on a given lens plane and for a given smoothing angle $\\theta$. We showed that these variances can be theoretically predicted from the knowledge of the nonlinear power spectrum under a given cosmological model. Our tests indicate that the proposed model is far better than the lognormal in the ultra-nonlinear regime of small smoothing angles ($< 1'$) and describes the low-probability, high $\\k$ tail of the PDF of $\\k$ with quite high accuracy. We illustrated a number of possible applications of the model. With the PDF of kappa on each plane accurately defined by our model, we could mimic an expensive ray shooting simulation to find out the importance of multiple clumps of matter in producing strong lensing, by a hugely faster Monte Carlo technique. We showed that the probabilities that a single plane on the line-of-sight is supercritical can be approximately obtained by a simple analytical calculation. We further demonstrated that secondary matter clumps are unimportant at source redshifts $\\sim 1$ but may contribute up to about $15\\%$ of the single-plane cases at higher redshifts. The chance accumulation of significant amounts of secondary matter along the line-of-sight will therefore result in systematic inaccuracies in the determination of cluster masses, typically overestimating them by about $5\\%$ with in addition a $\\sim 5-6\\%$ standard deviation. We took a simplified model of the matter distribution inside the beam and were able to cast the problem of the determination of angular diameter distances into a multiple lensing problem. We could then use our analytic model to estimate the dispersions in angular diameter distances due to inhomogeneities in the matter field. As a cautionary note, we would like to add that the for lensing involving small angular scales, the role of baryons become important and the predictions from our model in this regime, should therefore be looked at as some kind of a lower limit. An important future improvement to our model would therefore be the inclusion of the effect of baryons. It should be noted that one important advantage of our method over the halo model approach is the ability to compute strong lensing by multiple lens planes, which is rather difficult in the halo approach. Therefore, our approach provides a complementary analytic tool to the halo model based calculations, in the sense that our model will be useful to examine the effect of secondary matter when combined with the halo model approach. \\par There are several ways in which the power of this model could be harnessed. As noted briefly at the end of Section~\\ref{sec:model}, this model may pave the way for a more accurate description of the {line-of-sight-integrated convergence} in the small angle regime where the lognormal approximation is sometimes inadequate. This issue is now in progress and will be presented elsewhere. It will also be of interest to study whether a similar model describes the 3-D PDF of dark matter density at very small scales. As the model is fixed entirely by the cosmology, one could test how efficiently it can be used to probe cosmological parameters. One of the several possible studies along these lines could be to determine how the predictions are altered by a different model of dark energy like the quintessence \\citep{caldwell_etal_1998}.  With the non-Gaussianities in the lensing magnification well described by this model, one can use the scatter in Type IA supernovae to better constrain parameters like $\\sigma_8$ \\citep{dodelson_vallinotto_2005, cooray_etal_2005} or use multiple lens systems to better estimate the Hubble parameter $H_0$ \\citep{fassnacht_etal_2005}. \\par The statistics of gravitational lensing encodes the statistical features of matter distribution and hence probes cosmology at both linear and nonlinear scales. We expect that the model proposed in this paper will be able to play an important role in  learning more about cosmology in a fashion complementary to that derivable from the analysis of large numerical simulations. The latter, of course, permit more detailed questions to be raised and answered than the analytical approach presented here. But the present approach is of sufficient accuracy and so much more computationally efficient than than based on direct simulation that we can now efficiently compare a suite of cosmological models. As an example of this, we could ask for the differences in the gravitational lensing predicted for models which give essentially identical, i.e. degenerate, cosmic microwave background signals \\citep[see, for e.g.][]{huey_etal_1999,maor_etal_2001,maor_etal_2002}."
    },
    "0512/hep-ph0512001_arXiv.txt": {
        "abstract": "We study the influence of nonlocality in the interaction on two spin one pairing patterns of two-flavor quark matter: the anisotropic blue color paring besides the usual two color superconducting matter (2SCb), in which red and green colors are paired, and the color spin locking phase (CSL). The effect of nonlocality on the gaps is rather large and the pairings exhibit a strong dependence on the form factor of the interaction, especially in the low density region. The application of these small spin-one condensates for compact stars is analyzed: the early onset of quark matter in the nonlocal models may help to stabilize hybrid star configurations. While the anisotropic blue quark pairing does not survive a big asymmetry in flavor space as imposed by the charge neutrality  condition, the CSL phase as a flavor independent pairing can be realized as neutral matter in compact star cores. However, smooth form factors and the missmatch between the flavor chemical potential  in neutral matter make the effective gaps of the order of magnitude $\\simeq 10$ keV, and a more systematic analysis is needed to decide whether such small gaps could be consistent with the cooling phenomenology.  \\noindent PACS number(s):  04.40.Dg, 12.38.Mh, 26.60.+c, 97.60.Jd ",
        "introduction": "The investigation of color superconducting phases in cold dense quark matter has attracted much interest in the last years since those phases could be relevant for the physics of compact star cores \\cite{Blaschke:2001uj}. Nevertheless, it became clear that large spin-0 condensates, like the usual two-flavor superconductor (2SC), although having large pairing gaps ($\\sim$ 100 MeV) and therefore a direct influence on the equation of state (EoS), may be disfavored by the charge neutrality condition if not unusually strong diquark coupling constants are considered \\cite{Aguilera:2004ag,Blaschke:2005uj,Ruster:2005jc}.  Alternatively, spin-1 condensates are being investigated \\cite{Schafer:2000tw,Alford:2002rz,Schmitt:2004et}. Due to their smallness, their influence on the EoS is negligible but they could strongly  affect the transport properties in quark matter and therefore have important consequences on observable phenomena like compact star cooling, see \\cite{Blaschke:2005dc}. A recent investigation of neutrino emission and cooling rates of spin-1 color superconductors constructed for conserved total angular momentum allows for color-spin locking, planar, polar, and {\\it A} phases \\cite{Schmitt:2005wg}. However,  none of these phases fulfills the requirements of cooling phenomenology that no ungapped quark mode should occur on which the direct Urca process could operate and lead to very fast cooling in disagreement with modern observational data \\cite{Grigorian:2004jq}.  In the present work, we consider spin-1 pairing patterns different from the above mentioned, like  the anisotropic third color (say blue) quark pairing besides the usual 2SC phase (2SCb) \\cite{Buballa:2002wy} and the s-wave color-spin locking phase (CSL) \\cite{Aguilera:2005tg}, which  have been introduced within the NJL model whereby small gaps in the region of  some fractions of MeV have been obtained. Such small gaps could help to suppress efficiently the direct Urca process in quark matter and thus to control the otherwise too rapid cooling of hybrid stars \\cite{Grigorian:2004jq,Popov:2005xa}.  One important feature is that the form of the regularization, i.e. via a sharp cut-off or a form factor function, is expected to have a strong impact on the resulting gaps due to the sensitive momentum dependence of the integrand in the gap equation \\cite{Buballa:2002wy}. Especially the behavior of quark matter in the density region of the suspected deconfinement transition plays a crucial role for determining the stability of compact star configurations. Models with a late onset of quark matter could  eventually lead to unstable hybrid star configurations.  For example, in \\cite{Aguilera:2005tg} it has been shown within the local NJL model how a possible pairing pattern for compact star matter could be described which fulfills the constraints from compact star cooling phenomenology. These require that all quark species be paired and the smallest pairing gap be of the order of 10 keV to 1 MeV with a decreasing density dependence. A caveat of the NJL model quark matter is, however, that a stable hybrid star can be realized only marginally, see \\cite{Baldo:2002ju}.  The advantage of nonlocal models is that they can describe the regularization of the quark interaction via form factor functions and therefore represent it in a smoother way, especially for low densities \\cite{Schmidt:1994di}. For  the case of the 2SC phase it has already been shown that the effect of nonlocality in the low density region is rather large and the pairing exhibits a strong dependence  with the form factor of the interaction \\cite{Aguilera:2004ag,Blaschke:2003yn}. Moreover, the early onset of quark matter for the dynamical chiral quark model, in contrast to the NJL model,  might help to stabilize hybrid star configurations.  As it has been shown in \\cite{Grigorian:2003vi} within a nonlocal generalization of the NJL model \\cite{Schmidt:1994di}, stable hybrid stars with large quark matter cores can be obtained. In order to describe the properties of these stars consistently, including their cooling phenomenology, the description of diquark pairing gaps as well as emissivities and transport properties for the above motivated CSL phase should be given also within a nonlocal quark model.  As a first step in this direction we will provide in the present paper the spin-1 pairing gaps for a nonlocal, instantaneous chiral quark model under neutron star constraints for later application in the cooling phenomenology. We give here a first discussion of the influence of nonlocality of the interaction in momentum space when compared to the NJL model case and discuss the possible role of spin-one condensates for compact star applications.  The paper is organized as follows: in Section 2 we briefly review how the NJL model is modified when nonlocality is introduced using a three-\\-dimensional nonlocal chiral quark model; in Sections 3  and 4 we present the nonlocal version of the anisotropic blue color paring (2SCb) besides the usual two-flavor color superconducting (2SC) phase and of the color-spin locking phase (CSL), respectively. In Section 5 we present preliminary results for neutral matter in compact stars for the Gaussian form factor and discuss whether requirements for hybrid star cooling phenomenology could be met. Finally, in Section 6,  we draw the Conclusions. ",
        "conclusions": "We have studied the effect of nonlocality on spin-1 condensates in two flavor quark matter: the 2SC+spin-1 pairing of the blue quarks (2SCb) and the color spin locking (CSL) phase. We found that the size of these small gaps is very sensitive to the form of the regularization. The nonlocality has a strong impact on the low density region and we obtain an earlier onset for the superconducting phases. This might be crucial to stabilize quark matter cores in hybrid stars.  On the other hand, due to the flavor asymmetry, we find that the 2SCb pairing phase can be ruled out for compact stars applications. The CSL phase, in contrast, is flavor independent and therefore inert against the constraint of electric neutrality. For electrically neutral quark matter in beta equilibrium we obtain effective CSL gaps which are of the order of magnitude of $10$ keV, which might help to suppress the direct Urca process, in accordance with recent results from compact star cooling phenomenology. Nevertheless, since our results are strongly dependent on the parameters used and on the form of the regularization, more systematic studies are needed in order to decide whether the CSL phase could be applied in the description of compact stars."
    },
    "0512/astro-ph0512122_arXiv.txt": {
        "abstract": "We present the first images of the 691.473 GHz CO J=6--5 line in a protoplanetary disk, obtained along with the 690 GHz dust continuum, toward the classical T Tauri star TW~Hya using the Submillimeter Array. Imaging in the CO J=6--5 line reveals a rotating disk, consistent with previous observations of CO J=3--2 and 2--1 lines. Using an irradiated accretion disk model and 2D Monte Carlo radiative transfer, we find that additional surface heating is needed to fit simultaneously the absolute and relative intensities of the CO J=6--5, 3--2 and 2--1 lines. In particular, the vertical gas temperature gradient in the disk must be steeper than that of the dust, mostly likely because the CO emission lines probe nearer to the surface of the disk. We have used an idealized  X-ray heating model to fit the line profiles of CO J=2--1 and 3--2 with $\\chi^2$ analysis, and the prediction of this model yields CO J=6--5 emission consistent with the observations. ",
        "introduction": "Disks around pre-main sequence stars are the likely sites of the formation of planetary systems (e.g. ~\\citealp{beckwith99}). Our quantitative understanding of the physical properties of such disks --- in particular their sizes, radial and vertical temperature and density structures, and gas survival timescales --- has improved dramatically in recent years as (sub)millimeter-wave interferometers have imaged the gas and dust surrounding over a dozen T Tauri and Herbig Ae stars at $\\sim$1$''$--5$''$ resolution (eg. \\citealp{koerner_s95, guilloteau_d98, dartois_d03, qi_k03, qi_h04}). Observations and modeling of molecular emission of trace species can be used to derive the density and temperature structure in disks. Although the structure inferred from observations of a single line is not unique, observations of a sufficiently large number of lines of various molecules can be used to constrain the temperature and density independently (van Zadelhoff et al. 2001).  TW~Hya is the closest known classical T Tauri star (at a distance of 56 pc), and its circumstellar disk has been imaged in the CO J=3--2 and J=2--1 lines with the SMA (Qi et al. 2004). A detailed comparison of the data with 2D Monte Carlo radiative transfer models was used to constrain the disk properties. One problem with these models was that the predicted CO J=3--2 intensities were always lower than the observed values, no matter what disk parameters were varied (including outer disk radius, inclination angle and molecular depletion factors). Since the disk is externally irradiated, the gaseous component can be detectable through emission lines produced in temperature inversions of the disk atmosphere where CO lines are optically thick. The disk thermal structure is mainly derived from the fitting of the spectral energy distribution (SED), which provides a good description of the dust temperature profile. However, the sampling of the disk temperature distribution by gas and dust is different, as the optically thin dust probes much deeper into the mid-plane and closer to the star in radius. Qi et al. (2004) suggested that the difference between the CO intensity and the canonical model may be due to the higher gas temperature probed by the CO emission than the dust temperature. For this reason, we searched for CO J=6--5 emission, which is expected to be sensitive to the slope of the gas temperature inversion at the surface of the disk.  We present here the first aperture synthesis mapping of CO J=6--5 in a protoplanetary disk, using the Submillimeter Array (SMA). We show that inclusion of additional disk surface gas heating in the disk model, in the form of an idealized X-ray heating model, provides a better fit to both our previous SMA observations of the CO J=3--2 and 2--1 lines, and a much better fit to the observations of the CO J=6--5 line, than the model without additional heating. ",
        "conclusions": "Both X-ray (\\citealp{glassgold_n04}) and ultraviolet (\\citealp{Jonkheid_f04}) radiation are effective in heating the gas in the surface layers of disks. Recent X-ray satellite observations confirm that essentially all young stellar objects are luminous X-ray emitters (\\citealp{feigelson_m99}). Observational constraints on the various aspects of X-ray irradiation of protoplanetary disks are rapidly improving, and it appears that the emitted X-rays can affect the environment of a young star out to hundreds of AU (\\citealp{glassgold_n01}). \\citet{igea_g99} showed that Compton processes increase X-ray penetration through column densities of order $10^{24}$-$10^{25}$ cm$^{-2}$ and so can be important in disk physics. By fitting both ROSAT and ASCA spectra, an X-ray luminosity of $L_x$ $\\approx 2\\times 10^{30}$~ergs~s$^{-1}$ was derived for TW~Hya by \\citet{kastner_h99}, with prominent line emission at energies $\\sim$ 1 keV. With these X-ray parameters, Glassgold et al. (2004) calculated the ionization rate at the top of the disk atmosphere at a radial distance of 1 AU to be $\\sim6\\times 10^{-9}$~s$^{-1}$.  The X-ray luminosity of TW Hya is known to be highly variable (Kastner et al. 1999, 2002) due to flaring events, but the X-ray ionization important for heating should be a time-averaged effect. The model calculations concerning X-ray ionization and its effects on the gas temperature are subject to uncertainty, since the hard X-ray flux (most important for penetrating the disk) assumed in the Glassgold et al. treatment is a simplified representation of the real X-ray emission, which when examined at low resolution is actually broadly distributed from a few tenths to several keV (Kastner et al. 1999; Stelzer \\& Schmitt 2004).  Motivated by the X-ray observations, we focus on X-ray effects on the disk surface to describe the additional surface heating. In order to model the effects of X-rays in heating the gas in the upper atmosphere of the TW~Hya disk, we follow the Glassgold et al. (2001, 2004) scheme to estimate the gas temperature for the region of the disk probed by CO rotational lines. The X-ray heating at the disk surface is expressed as \\begin{equation} \\Gamma _{X}=4.8\\times 10^{-11}erg \\zeta _{X}\\left( \\frac{\\Delta \\epsilon_{h}}{30eV}\\right) n_{H} \\end{equation} in terms of a mean heating energy per ionization, $\\Delta\\epsilon_h$ and the X-ray ionization rate $\\zeta _{X}=6\\times 10^{-9}s^{-1}\\left( \\frac{AU}{r}\\right) ^{2}$, which is adopted from Glassgold et al. 2004 for TW~Hya where r is the cylindrical radius from the disk axis.  The dust-gas cooling rate is taken to be \\begin{equation} \\Lambda _{dg}=2\\times 10^{33}erg cm^{3} s^{-1}\\left\\lceil \\frac{\\frac{\\rho _{d}}{\\rho _{g}}}{0.01}\\frac{500 {\\rm \\AA}}{a}\\frac{\\alpha }{0.5}\\right\\rceil T^{0.5}\\left(T-T_{d}\\right) n_{H}^{2} \\end{equation} where $\\alpha$ is the accommodation coefficient (from equation 3 of Glassgold et al. 2004), T is the gas-kinetic temperature, and T$_d$ is the mean dust temperature. When $T>T_d$ the dust cools the gas. The net volumetric gain rate of gas energy due to collisions with dust is then proportional to the local dust-to-gas ratio $\\frac{\\rho _{d}}{\\rho _{g}}$ and inversely proportional to the mean dust radius $a$. With the heating of the gas by X-rays balanced by the dust-gas cooling rate, we can calculate the difference of gas and dust temperature at each locale within the disk. Figure 3 shows the gas$-$dust temperature differences overlaid with the dust temperature contours from the TW~Hya model (Calvet et al. 2002). By incorporating these temperature difference into the disk model, we can fit the CO lines using the $\\chi^2$ analysis in the ($uv$) plane to estimate the goodness of fit for the various disk parameters, including the outer radius R$_{out}$ and the inclination angle $i$, $(vsini)_{100AU}$. A similar $\\chi^2$ analysis to determine best fit disk parameters was carried out on the disk of DM Tau by Guilloteau and Dutrey (1998). The $\\chi^2$ analysis in the visibility plane is essential to avoid the non-linear effects of deconvolution in the imaging process.  We adopt the disk physical density and temperature structure as derived by Calvet et al. (2002), where we take the disk temperature as the dust temperature, and we use the difference between the gas and dust temperature derived by balancing the relevant heating and cooling processes. We produce a grid of models with a range of various disk parameters and depletion schemes to simulate the disk as imaged by a telescope with the resolution constrained only by the grid sampling (typically of order 5-10 AU in the outer disk, or $0\\farcs1-0\\farcs2$). We use a 2D Monte Carlo model (\\citealp{hogerheijde_v00}) to calculate the radiative transfer and molecular excitation, and we produce simulated observations of the model disks with the MIRIAD software package using the UVMODEL routine to select synthetic visibility observations at the observed $(u,v)$ spacings. Then the $\\chi^2$ distance is calculated between the simulated visibility and the CO J=3--2 and J=2--1 data with sufficient signal-to-noise. The best fit parameters as derived by minimizing the $\\chi^2$ distance are shown in Table 2 (details to be presented in a future paper). The simulated spectra of the CO lines including X-ray heating are shown in Figure 2 with red solid lines. Including X-ray heating improves the agreement of the model with the data.  Note that the CO J=6--5 model prediction is based only on the lower transition lines.  The value $(\\frac{\\rho _{d}}{\\rho _{g}})/a$ is constrained by the $\\chi^2$ fit to be around $4\\times10^{-4}$ with $a$ in $\\micron$. If the local dust-to-gas ratio is around 0.01, then the mean dust radius is around 25$\\micron$, which seems large for the surface of the disk. If dust settles in the mid-plane, then the local dust-to-gas ratio might be lower, and if it were as low as as 0.001, then the mean dust radius would be around micron size or smaller, which seems more plausible.  Even though the disk model with X-ray heating can effectively match the various CO observations of TW~Hya, especially the line intensity at $\\sim$2$''$ resolution, the predicted width of CO J=6--5 line is nearly 50\\%  larger than observed. In the model, the CO J=6--5 emission originates mostly from the upper layers of the inner 60 AU of the disk, consistent with the double-peaked appearance in the model spectra, as the density of the outer disk is so low that CO J=6--5 becomes sub-thermally excited. The narrow observed linewidth of CO J=6--5 suggests that most of the emission is actually coming from the outer part of the disk (less contribution from the velocity projection from the regions with high Keplerian rotation). There are at least two scenarios that could account for this: (1) There is not much CO in the inner disk. Although it is possible to make the model CO J=6--5 spectra narrower by blanking out the CO emission from the inner 20 AU, it would then be hard to fit both the CO J=3--2 and 2--1 lines simultaneously, and also begs the question as to the origin of the CO M-band $v$=1--0 emission (\\citealp{rettig_h04}). (2) The gas temperature of CO in the inner 20 AU is actually lower than predicted by the models, possibly due to either additional cooling or perhaps shadowing induced by a puffed up inner rim of the disk, so that the emission from the higher excitation J=6--5 line is suppressed but with little effect on the lower excitation J=3--2 and 2--1 lines. Clearly the current model is still too simplified to fit all the CO lines perfectly. Further analysis of the thermal structure in the inner disk will hopefully help to resolve this puzzle."
    },
    "0512/astro-ph0512314_arXiv.txt": {
        "abstract": "{We discuss the effects of an enhanced interstellar radiation field (ISRF) on the observables of protostellar cores in the Orion cloud region. Dust radiative transfer is used to constrain the envelope physical structure by reproducing SCUBA 850~$\\mu$m emission. Previously reported $^{13}$CO, C$^{17}$O and H$_2$CO line observations are reproduced through detailed Monte Carlo line radiative transfer models. It is found that the $^{13}$CO line emission is marginally optically thick and sensitive to the physical conditions in the outer envelope. An increased temperature in this region is needed in order to reproduce the $^{13}$CO line strengths and it is suggested to be caused by a strong heating from the exterior, corresponding to an ISRF in Orion $10^3$ times stronger than the ``standard'' ISRF. The typical temperatures in the outer envelope are higher than the desorption temperature for CO. The C$^{17}$O emission is less sensitive to this increased temperature but rather traces the bulk envelope material. The data are only fit by a model where CO is depleted, except in the inner and outermost regions where the temperature increases above 30-40~K. The fact that the temperatures do not drop below $\\approx 25$~K in any of the envelopes whereas a significant fraction of CO is frozen-out suggest that the interstellar radiation field has changed through the evolution of the cores. The H$_2$CO lines are successfully reproduced in the model of an increased ISRF with constant abundances of 3--5$\\times 10^{-10}$. ",
        "introduction": "Giant molecular clouds are the formation sites of massive stars in our Galaxy with the nearby Orion molecular clouds being the prime candidates for detailed studies of the earliest protostellar stages. An interesting difference for the studies of these cores compared to isolated systems is the large number of OB stars in the immediate vicinity: the UV radiation from these early-type stars ionizes their surrounding material within a few parsecs, and also affects the thermal balance and chemistry in intermediate and low-mass protostellar cores distributed over much larger scales. As recent studies indicate that large numbers of solar-type stars may be formed in these regions, it is of great importance to address the feedback between high- and low-mass star formation, in particular, for comparison with low-mass protostars formed in relative isolation in clouds such as Taurus.  In this paper we present a radiative transfer study of three intermediate mass protostars in Orion from the sample of \\cite{johnstone03}. The physical properties of their envelopes are established from 1D dust radiative transfer modeling, and CO isotopic and H$_2$CO line observations are analyzed using Monte Carlo line radiative transfer. In particular, we discuss the importance of heating the protostellar envelope via an enhanced external radiation field and the constraints on this heating from optically thick CO isotopomers.  \\cite{li03} recently estimated the gas kinetic temperatures for a sample of pre-stellar cores in Orion using inversion transitions of NH$_3$. They found that these cores had lower temperatures than their surroundings, which they attributed to the impact of the strong interstellar radiation field (ISRF) in the Orion region. This is reminiscent of the situation for pre-stellar cores in more quiescent star forming regions where temperature gradients due to a standard interstellar radiation field are found in radiative transfer models reproducing submillimeter continuum maps \\citep[e.g.,][]{evans01}. \\cite{wilson99orion} suggested that the ratios of peak temperatures of optically thick lines with different critical densities could be used as a diagnostic of temperatures. Applied to cores in Orion their results also suggested strong temperature gradients in the region around Orion BN/KL with more dense gas traced by NH$_3$ inversion lines colder ($\\sim 24$~K) than less dense gas ($\\sim 40$~K) probed by CO lines. In contrast \\citeauthor{wilson99orion} found similar temperatures probed by the NH$_3$ and CO lines for cores further south from the Orion BN/KL region and they argued that these cores are dominated by internal rather than external heating. This illustrates that an important point for the thermal balance of protostellar objects is whether external heating can compete with the central source luminosity.  The strong interstellar radiation field may also be reflected in the chemistry, for example by increasing the photodissociation and photoionization. The strength of the UV field can for example be probed by the emission of C$^+$ and has been applied to the Orion clouds where UV fields enhanced by factors of $10^3-10^5$ have been suggested \\citep[e.g.,][]{tielens85b,tielens85a,burton90,stacey93,moorkerjea03}. The strong UV field will also lead to enhanced abundances of electrons and thus abundance decreases of molecular ions, such as HCO$^+$ and N$_2$H$^+$ that otherwise work as destroyers of common molecules. Support for this is found within the OMC1 cloud core located immediately behind the Orion Nebula Cluster where \\cite{ungerechts97} infer abundance gradients along the cloud, with increasing abundances of species such as HCN, CH$_3$OH, HC$_3$N, and SO toward the Orion BN/KL region.  The astrochemistry study by \\cite{johnstone03} considered a selection of cornerstone molecules in order to quantify the range of conditions for which individual molecular line tracers provide physical or chemical information. The morphological study compared a variety of locations along the Integral Shaped Filament (ISF) in Orion A \\citep{bally87,johnstone99}, chosen to represent a range of physical conditions including enshrouded protostars, a bright PDR knot, and a shock front. The main conclusion of the paper was that the abundances of many molecular species correlate with source energetics, likely a result of the importance of temperature dependent desorption in maintaining gas-phase molecules. A significant finding of the study was the need for a warm outer envelope for all of the protostellar sources, as expected given the proximity of numerous O and B stars. The study suffered, however, from a simplistic treatment of the molecular line abundances derived from single temperature and density statistical equilibrium calculations and thus was unable to fully consider the effect of an enhanced radiation field on the outer regions of the embedded protostars.  This paper presents a continuation of the \\cite{johnstone03} study, performing a detailed radiative transfer analysis of the physical and chemical properties of a subsample of the regions studied in that paper. We concentrate on the embedded intermediate mass protostellar sources, MMS6, MMS9, and FIR4 where the assumption of constant density and temperature throughout the envelope is most suspect. All three sources have large envelope masses, $M_{\\rm env} > 10 M_\\odot$.  The implied dust temperatures, $T_{\\rm d}$, for the sources are $>15$~K while the $^{13}$CO lines have peak brightness temperatures $\\gtrapprox 30\\,$K which is difficult to produce {\\it unless} the cores are bathed in an enhanced interstellar radiation field. The properties of the cores studied in this paper are summarized in Table~\\ref{sample} with further details presented in \\cite{johnstone03}. Even though the focus of this paper is on a few selected cores in Orion, a lot of the results brought up in this discussion are valid beyond these specific sources.  This paper is laid out as follows: sect.~\\ref{model} introduces the general problem, including the observational indications for a strong interstellar radiation field, and describes the modeling approach, including the possible constraints from continuum and line observations. Sect.~\\ref{discuss} then discusses possible refinements and implications of the models, in particular, the resulting CO abundance structures and the reproduction of H$_2$CO multi-transition observations. The evolutionary implications are discussed in Sect.~\\ref{s:implic}.  Sect.~\\ref{s:conclusion} concludes the paper suggesting possible further tests and future work.  \\begin{table*} \\caption{Sample of sources in Orion. For details see \\cite{johnstone03}. }\\label{sample} \\begin{center} \\begin{tabular}{lllllll}\\hline\\hline Source       &                                 & MMS6 & MMS9 & FIR4  \\\\ \\hline $L_{\\rm bol}$ &[$L_\\odot$]                     & 60   & 90   & 400       \\\\ Projected distance to the Trapezium$^a$  &    [pc] & 2.9  & 2.3    & 1.7    \\\\ $S_{850 \\mu{\\rm m}}$ (peak) & [Jy beam$^{-1}$] & 7.5  & 2.5  & 7.5      \\\\ $I$($^{13}$CO 3--2)$^b$ &[K~km~s$^{-1}$]              & 40.7 & 59.1 & 76.2 \\\\ $I$(C$^{17}$O 3--2)$^b$ & [K~km~s$^{-1}$]             & 4.78 & 2.88 & 3.97 \\\\ $I$(H$_2$CO $3_{03}-2_{02}$)$^b$ & [K~km~s$^{-1}$]    & 5.33 & 3.07 & 9.82  \\\\ $I$(H$_2$CO $3_{22}-2_{21}$)$^b$ & [K~km~s$^{-1}$]    & 0.53 & 0.28 & 2.96  \\\\ $I$(H$_2$CO $5_{05}-4_{04}$)$^b$ & [K~km~s$^{-1}$]    & 2.56 & 0.94 & 5.24  \\\\ \\hline \\end{tabular} \\end{center}  $^a$Assuming a distance to Orion of 450~pc. $^b$Line intensity, $I=\\int T_{\\rm MB}\\,{\\rm d}v$. \\end{table*} ",
        "conclusions": "\\label{s:conclusion} We have presented an analysis of submillimeter continuum and line observations toward three intermediate mass protostellar cores in the Orion molecular cloud using detailed radiative transfer models. The main conclusions of this study are:  \\begin{itemize} \\item From simple luminosity and temperature considerations these sources must be subject to a strong external heating to reproduce observed $^{13}$CO 3--2 line temperatures and profiles. \\item Detailed radiative transfer modeling confirms this, constraining the ISRF enhancement to $10^3-10^4$ times the standard ISRF from the studied sources. Different strengths of the ISRF between MMS6 and MMS9 in one case and FIR4 in the other can be explained in part due to the latter being closer to the Trapezium stars supplying the strong UV field. \\item The C$^{17}$O 3--2 lines are optically thin and thus not sensitive to the same temperature enhancement. Differences between the $^{13}$CO and C$^{17}$O line widths reflect this difference between the optically thick and thin lines, which is well reproduced by the models. To reproduce the C$^{17}$O line intensities, significant CO depletion must have occurred in the part of the envelope where the temperature is lower than 30-45~K. A difference in the CO desorption temperatures derived for the sources, however, suggest an evolutionary difference with the OB stars (and thus UV field) evolving over similar timescales as the protostellar cores themselves. \\item Multi-transition H$_2$CO observations indicate high temperatures in the cores and their line ratios can only be reproduced in models with a strong ISRF. Constant abundances of 3$\\times 10^{-10}$--5$\\times 10^{-10}$ provide a good fit to the line intensities. \\end{itemize}  This work illustrates the necessity of establishing the environmental impact for the evolution of pre- and protostellar cores - for example in regions such as Orion where the impact of the radiation stars have a large impact on newly formed low- and intermediate mass protostars. As the above discussions also illustrate, detailed radiative transfer models can be used to address some of these issues, in particular if the chemistry is taken into account.  Additional mid-infrared observations of the cores can further constrain the relative contributions of the internal and external radiation field to the observed spectral energy distributions as it has previously been done for cores near the galactic center \\citep{lis01}. Further continuum observations with (sub)millimeter interferometers such as the SMA, CARMA and eventually ALMA could further confirm the temperature structure of the envelopes at smaller scales and in particular whether desorption of molecules are seen toward their centers due to the heating by the central protostar. Likewise in an extreme interstellar radiation field such as that of the Orion further studies of the chemistry in the outer envelopes and intervening material in the ridge are interesting to fully understand the impact of the ionizing sources.  Understanding these processes are important as the thermal balance and pressure of the cores controls the conditions and outcome of the collapse and thus establishing the properties of the newly formed stars.  \\begin{acknowledgement} We thank the referee, Neal Evans, for insightful comments which greatly improved the paper. The research of JKJ was funded through a Ph.D. stipend from the Netherlands Research School for Astronomy (NOVA) and NASA Origins Grant NAG5-13050. This research was also supported by a grant from the Natural Science and Engineering Research Council of Canada to DJ and a grant for from The Research Corporation (SDD). Astrochemistry research in Leiden is supported by an NWO/Spinoza grant. DJ wishes to thank the Sterrewacht Leiden for its kind hospitality during the past four summers, during which much of this research was conducted. \\end{acknowledgement}"
    },
    "0512/astro-ph0512587_arXiv.txt": {
        "abstract": "We investigate the merging process in N-body self-gravitating system from the viewpoints of the local virial relation which is the relation between the  local kinetic energy and the  local potential. We compare both the density profile and the phase space density profile in cosmological simulations with the critical solutions of collisionless static state  satisfying the local virial (LV) relation. We got the results that the critical solution can explain the characteristic density profile with the appropriate value of  anisotropy parameter $\\beta \\sim 0.5$. It can also explain the power law of the phase space density profile in the outer part of a bound state. However, it fails in explaining the central low temperature part which is connected to the  scale invariant  phase space density. It can be well fitted to the critical solution with the higher value of $\\beta \\sim 0.75$. These results indicate that the LV  relation is not compatible with the  scale invariant phase space density in cosmological simulation. ",
        "introduction": "Introduction}  It is well known that astronomical objects are formed through the merging process of small clusters under the homogeneous expanding background. It is  also well known that the  bound state after the merging process  in  cosmological simulations takes rather universal form of a cuspy density profile, $\\rho \\sim r^{-\\gamma}$ in the inner part, where $\\gamma \\approx 1$ \\cite{Navarro96}. There is another universal character for the virialized state following merging process, that is, the scaling law of a phase space density, $\\rho/\\sigma^3 \\sim r^{-\\alpha}$ where $\\alpha \\approx 1.9$ \\cite{Taylor01}. Several models are proposed to explain both of these characters. For example, W.Dehnen and D.E.McLaughlin show that the critical solution for the static Vlasov equations under the scale invariant  phase space density is compatible with NFW density profile \\cite{Dehnen05}.   In our previous paper, we examined  cold collapse simulations and got the result that the bound state after a collapse has several remarkable characters \\cite{Iguchi04,Sota04,Osamu05}. The LV relation is one of such characters. The Plummer model which is the analytical solution for polytrope with index n=5 satisfies this condition. The Plummer model has the property that the velocity dispersion is isotropic. However,  the velocity dispersion is not always isotropic for the quasi-equilibrium state and  is characterized with non-zero anisotropy parameter $\\beta \\left( r \\right)$. The analytical solution exists even for the non-zero constant value of $\\beta $ under the condition of the LV relation \\cite{Evans05}. We showed by combining these analytical solutions that the bound state after the cold collapse can be described quite well by the solution with the LV condition \\cite{Osamu05,Sota05}.  We also showed that the analytical solution in \\cite{Evans05} can be characterized as  a critical solution connecting two fixed points as is shown in  \\cite{Dehnen05}. We got the results that the LV relation is rather special in that the critical solution connecting the two fixed points exist only in the case with the virial ratio $b=1$ among the models with general constant $b$ models. Here we investigate whether or not the model with the LV relation can explain several characters of the cosmological simulations by using the constant $\\beta$ critical solutions. ",
        "conclusions": ""
    },
    "0512/astro-ph0512064_arXiv.txt": {
        "abstract": "TIMMI2 diffraction--limited mid-infrared images of a multipolar proto-planetary nebula IRAS 16594$-$4656 and a young [WC] elliptical  planetary nebula IRAS 07027$-$7934 are presented. Their dust shells are for the first time resolved (only marginally in the case of IRAS 07027$-$7934) by applying the Lucy-Richardson deconvolution algorithm to the data, taken under exceptionally good seeing conditions ($\\leq$0.5\\arcsec). IRAS 16594$-$4656 exhibits a two-peaked morphology at 8.6, 11.5 and 11.7 $\\mu$m which is mainly attributed to emission from PAHs. Our observations suggest that the central star is surrounded by a toroidal structure observed edge-on with a radius of 0.4$\\arcsec$ ($\\sim$640 AU at an assumed distance of 1.6 kpc) with its polar axis at P.A.$\\sim$80\\degree, coincident with the orientation defined by only one of the bipolar outflows identified in the HST optical images. We suggest that the material expelled from the central source is currently being collimated in this direction and that the multiple outflow formation has not been coeval. IRAS 07027$-$7934 shows a bright, marginally extended emission  (FWHM=0.3$\\arcsec$) in the mid-infrared with a slightly elongated shape along the N-S direction, consistent with the morphology detected by HST in the near-infrared. The mid-infrared emission is interpreted as the result of the combined contribution of small,  highly ionized PAHs and relatively hot dust continuum. We propose that IRAS 07027$-$7934 may have recently experienced a thermal pulse (likely at the end of the AGB) which has produced a radical change in the chemistry of its central star. ",
        "introduction": "Planetary nebulae (PNe) are the result of the evolution of low- to intermediate-mass stars (0.8--8 M$_{\\odot}$). These stars experience a phase of extreme mass loss during the previous asymptotic giant branch (AGB) that causes the ejection of the stellar envelope. When this mass loss ceases the AGB phase ends and the star evolves into a short-lived evolutionary stage called the `post-AGB' or `proto-PN' (PPN) phase just before the star becomes a PN. At present, the formation of axisymmetric structures in PNe (ranging from elliptical to bipolar) is believed to be completed by the end of the AGB phase (Balick \\& Frank 2002; Van Winckel 2003). But unlike for PNe, the study of PPNe is more difficult since their central stars (CSs) are usually too cool to photoionize the gas. Therefore, we cannot study the formation of axisymmetric morphologies in PPNe by mapping the ionized gas. We must use alternative techniques based on the analysis of: (i) the light scattered by the surrounding dust at optical wavelenghts; (ii) the neutral molecular gas in the envelope in the near-infrared (H$_2$), submilimeter (e.g. CO) or radio domain (e.g OH, SiO, H$_{2}$O, CO); and (iii) the dust emission emerging at mid- to far-infrared wavelengths.  PPNe generally show a double-peaked spectral energy distribution (SED) (Kwok 1993; Volk \\& Kwok 1989; van der Veen, Habing \\& Geballe 1989) with the photospheric emission coming from the central  star dominating in the optical range and a strong infrared excess  indicating the presence of a cool detached envelope (T$_{d}$$\\sim$150--300 K). This strong infrared excess is produced by the thermal emission of the dust  present in their circumstellar shells previously expelled during the  AGB phase. The processes that lead to a wide variety of different morphologies observed in PNe  (e.g. Manchado et al. 2000) are, however,  still unknown. Several mechanisms have been proposed: the interaction of stellar winds (e.g. Mellema 1993), binary systems as central stars (e.g. Bond \\& Livio 1990; Morris 1987), non-radial pulsations (e.g. Soker \\& Harpar 1992) or the influence of magnetic fields (e.g. Pascoli 1992; Soker \\& Harpar 1992; Garc\\'{\\i}a-Segura et al. 1999). To establish  which one(s) of the above is the dominant process, it is essential to study  these morphologies as early as possible after the departure from the spherical symmetry takes place, that is, in the PPN phase. Only recently, with  the help of high spatial resolution observations it has been possible to study the intrinsic axisymmetric nature of the dust shells around a few compact PPNe at subarcsec level (Meixner et al. 1997; Meixner et al. 1999; Ueta et al. 2001).  In this paper, we present for the first time mid-infrared images (8--13 $\\mu$m) at subarcsec level of a PPN IRAS 16594$-$4656 (hereafter I16594) and of a very young [WC] PN IRAS 07027$-$7934 (hereafter I07027) with the aim of mapping the dust  emission originated in the innermost regions of their circumstellar dust shells. The observations made in the mid-infrared are presented in Sect. 2 while the data reduction process is described in Sect. 3.  We show  the results obtained in Sect. 4, which are later discussed in  Sect. 5. The main conclusions derived from our analysis are given in  Sect. 6. ",
        "conclusions": "We have presented diffraction limited mid-infrared images of the PPN I16594 and the [WC] PN I07027 at 8.6, 11.5 and 11.7 $\\mu$m taken under exceptionally good seeing conditions ($\\leq$0.5\\arcsec). By applying the Lucy-Richardson deconvolution algorithm, we have resolved, for the first time, the subarcsecond dust shell structures around both objects.  I16594 displays two emission peaks in the innermost region of the circumstellar dust shell at the three wavelengths observed. This two-peaked mid-infrared morphology is interpreted as an equatorial density enhancement revealing the presence of a dusty toroidal structure with a 0.4$\\arcsec$ radius size (or $\\sim$640 AU corresponding to a dynamical age of $\\sim$190 yr at the assumed distance of 1.6 kpc). The observed size is used to derive the dust temperature at the inner radius of the shell. This result has been combined with the information derived from the ISO observations of I16594 to conclude that the mid-infrared emission detected in our TIMMI2 images must be dominated by PAH molecules or clusters which must be mainly distributed along the torus, as suggested by the similar size and morphology observed in all filters. We have also found that the axis of symmetry observed in the mid-infrared is well aligned with only one  of the bipolar outflows (at P.A.$\\sim$84$\\degree$) seen as optical reflection nebulae in the optical HST images. We suggest that the multiple outflow formation has not been coeval and that, at present, the outflow material is being ejected in this direction. Consistently, the H$_2$ shocked-emission seen in the HST NICMOS image is mainly distributed along the same bipolar axis where the fast post-AGB wind is interacting with the slow moving material ejected during the previous AGB phase. The presence of several other bipolar outflows at a variety of position angles may be the result of past episodic mass loss events.  I07027 exhibits a slightly asymmetric mid-IR emission core which is only marginally extended along the north-south direction with FWHM=0.3$\\arcsec$ at 8.6 and 11.5 $\\mu$m. This is the same orientation observed in recent HST images of the source taken in the near-infrared. The mid-infrared emission is attributed to a combination of emission from highly ionized, small PAH molecules plus relatively warm dust continuum located very close to the central star. The characteristics of the PAH emission observed in the ISO spectrum are also consistent with this interpretation. Taking into account the spatial distribution of the C-rich material deduced from our observations and because the OH maser emission from I07027 is expected to be located in the external and cooler regions, we propose that the dual chemistry observed in I07027 must be interpreted as the consequence of a recent thermal pulse (probably at the end of the previous AGB phase) which has switched the chemistry of the central star from the original O-rich composition to a C-rich one within the last 500 yrs. This might be the commom mechanism which originates the dual chemistry and strong stellar winds usually observed in other [WC]-type CSPNe."
    },
    "0512/astro-ph0512252_arXiv.txt": {
        "abstract": "We have mapped the SiO $J$=5--4 line at 217\\,GHz from the HH\\.211 molecular outflow with the Submillimeter Array (SMA). The high resolution map (1.6$''{\\times}$0.9$''$) shows that the SiO $J$=5--4 emission comes from the central narrow jet along the outflow axis with a width of $\\sim$0.8$''$ ($\\sim$\\,250\\,AU) FWHM\\@. The SiO jet consists of a chain of knots separated by 3--4$''$ ($\\sim$\\,1000\\,AU) and most of the SiO knots have counterparts in shocked H$_2$ emission seen in a new, deep VLT near-infrared image of the outflow. A new, innermost pair of knots are discovered at just $\\pm$2$''$ from the central star. The line ratio between the SiO $J$=5--4 data and upper limits from the SiO $J$=1--0 data of \\citet{Cha01} suggests that these knots have a temperature in excess of 300--500\\,K and a density of (0.5--1)$\\times$10$^7$\\,cm$^{-3}$. The radial velocity measured for these knots is $\\sim$\\,30\\,km\\,s$^{-1}$, comparable to the maximum velocity seen in the entire jet. The high temperature, high density, and velocity structure observed in this pair of SiO knots suggest that they are closely related to the primary jet launched close to the protostar. ",
        "introduction": "Several highly-collimated molecular outflows driven by deeply-embedded young stellar objects are known to have an extremely-high velocity (EHV) flow component, with terminal velocities of 50--150\\,km\\,s$^{-1}$\\citep[e.g.,][]{Bac96}. These EHV components are closely confined to the axes of the lobes and have large momenta, comparable to those of the slowly-moving (20--30\\,km\\,s$^{-1}$) ``classical'' outflows. As a result, it is felt that the EHV flow is closely connected with the ``primary jet'' responsible for driving the broader molecular outflow.  The HH\\,211 outflow (D\\,$\\sim$\\,315\\,pc) was discovered via near-infrared H$_2$ imaging \\citep{McC94} and is an archetypal outflow with a highly-collimated EHV jet. The outflow is driven by a low-luminosity (3.6\\,$L_{\\odot}$) Class~0 protostar ($T_{\\rm bol}$\\,$\\sim$\\,33\\,K) and is thought to be extremely young ($\\tau_{\\rm dyn}$\\,$\\sim$\\,750\\,yr). The CO $J$=2--1 images of \\citet[][hereafter GG99]{Gue99} with a spatial resolution of 1.5$''$ show remarkable features: the low-velocity CO delineates a pair of cavities whose tips are associated with the near-infrared H$_2$ emission, while the high-velocity CO traces a narrow jet whose velocity increases linearly with distance from the star. In spite of a small opening angle of 22$^{\\circ}$, the CO $J$=2--1 emission from the cavity appears on both sides of the central star at velocities close to the systemic velocity of $V_{\\rm LSR}$\\,=\\,9.2\\,km\\,s$^{-1}$ \\defcitealias{Gue99}{GG99}\\citepalias{Gue99}, implying that the outflow axis lies within 10$^\\circ$ of the plane of the sky. The jet component (but not the cavity) is also traced by thermal SiO emission, in $J$=1--0 \\citep[][hereafter CR01]{Cha01}, $J$=5--4 \\citep{Gib04}, $J$=8--7 and $J$=11--10 \\citep{Nis02} transitions. Since these lines have critical densities larger than ${\\sim}10^6$\\,cm$^{-3}$ and the energy level of $J$=11 is higher than 100\\,K, the detection of these SiO lines means that the jet is much denser and warmer than the lower-velocity cavity component. ",
        "conclusions": "\\subsection{The highly-collimated SiO jet} SiO $J$=5--4 emission was detected from HH\\,211 in two velocity ranges from $-$24km\\,s$^{-1}$ to $-$4\\,km\\,s$^{-1}$ (blueshifted) and from +4\\,km\\,s$^{-1}$ to +32\\,km\\,s$^{-1}$ (redshifted) with respect to the systemic velocity of $V_{\\rm LSR}$\\,=\\,9.2\\,km\\,s$^{-1}$. Integrated intensity maps of the blueshifted and redshifted emission are shown in Figure~1 and velocity channel maps at 4\\,km\\,s$^{-1}$ intervals are presented in Figure~2. The SiO $J$=5--4 emission is seen to be concentrated exclusively in the jet-like narrow region along the outflow axis and there is no counterpart to the low-velocity cavity component seen in CO $J$=2--1 \\citepalias{Gue99}. We have smoothed the SMA map to a resolution of 22$''$ and compared it with the SiO $J$=5--4 spectrum observed with the JCMT \\citep{Gib04}, finding that 80--100\\% of the single-dish flux is recovered by the SMA, despite the missing short spacing information. This confirms that almost all of the SiO $J$=5--4 emission arises from a narrow jet. After deconvolution of the SMA beam, the width of the SiO $J$=5--4 jet is $\\sim$0.8$''$ ($\\sim$\\,250\\,AU) FWHM\\@. It comprises a chain of knots each separated by $\\sim$\\,3--4$''$ ($\\sim$\\,1000\\,AU), with five discrete knots (R1--R5 in Fig.~2) on the redshifted side and six (B1--B5) on the blueshifted side. Knots B2--B5, R3, and R4 all appear to have near-infrared H$_2$ counterparts, while the lack of H$_2$ emission coincident with B1, R1, and R2 is probably due to the very high extinction associated with the dense protostellar envelope as traced in H$^{13}$CO$^+$ \\citepalias{Gue99} and NH$_3$ \\citep{Wis01}. In addition, the innermost knots, B1 and R1, have no counterparts in SiO $J$=1--0 \\defcitealias{Cha01}{CR01}\\citepalias{Cha01} or CO $J$=2--1 \\citepalias{Gue99}.  The blueshifted side of the SiO $J$=5--4 jet ends near an H$_2$ knot roughly 23$''$ to the east of the protostar and the redshifted side $\\sim$\\,18$''$ to the west of the center. The SiO $J$=5--4 does not not extend as far as the CO $J$=2--1 jet observed by \\citetalias{Gue99}, terminating at features BI and RI (marked on Figure 1), respectively: no SiO $J$=5--4 is seen beyond these positions, even though our three-pointing mosaic covered almost the entire outflow. The lowest transition SiO $J$=1--0 also terminates at BI and RI \\citepalias{Cha01}, and the implication is that beyond these points, the jet is no longer dense enough to excite SiO emission. An alternative hypothesis is that beyond BI and RI, the SiO has been destroyed via chemical processing in shocks, as described by \\citetalias{Cha01}.  An SiO $J$=5--4 position-velocity diagram along the jet axis (Figure~3) shows a different velocity structure to that seen in CO $J$=2--1 \\citepalias{Gue99}. In particular, the SiO $J$=5--4 exhibits a large velocity dispersion at the innermost knots B1 and R1, suggesting that these knots are kinematically distinct from those in the outer part. At $\\sim$\\,2$''$ ($\\sim$\\,630\\,AU) from the protostar, the velocity centroid and the terminal velocity reach $\\pm$\\,18\\,km\\,s$^{-1}$ and ${\\sim}{\\pm}$\\,30\\,km\\,s$^{-1}$, respectively . The outer part of the SiO $J$=5--4 shows a Hubble-like velocity structure, previously seen for the CO $J$=2--1. However, the SiO jet is moving $\\sim$\\,5\\,km\\,s$^{-1}$ faster than the CO jet, reaching a maximum radial velocity of $\\sim$\\,35\\,km\\,s$^{-1}$. Assuming that the axis of the flow is inclined by $\\sim$\\,10$^{\\circ}$ out of the plane of the sky, the deprojected outflow velocity in the jet would correspond to $\\sim$\\,200\\,km\\,s$^{-1}$, typical of the velocity of the primary jet driven by low-mass stars.  In Figure~4, we compare line profiles of the SiO $J$=5--4 at knots B1, R1, B4, and R3, with those observed in $J$=1--0 with the VLA by \\citetalias{Cha01}. In order to obtain line intensity ratios (Table~1), we convolved the data cubes of the two transitions to an equal angular resolution of 1.6$''$. Then, to derive physical parameters in the jet as traced by SiO emission, we carried out large velocity gradient (LVG) statistical equilibrium calculations following the method of \\citet{Nis02}. \\citetalias{Cha01} reported peak brightness temperatures of the SiO $J$=1--0 of 54\\,K and 47\\,K in the redshifted and blueshifted gas, respectively, averaged over a 0.58$''{\\times}$0.43$''$ beam. These SiO $J$=1--0 emission peaks correspond to the positions of R3 and B4 in our data, where the 5--4/1--0 ratio is seen to be $\\sim$\\,2. The LVG calculations suggest that a density of (0.5--1)$\\times$10$^{7}$\\,cm$^{-3}$, a temperature of 60--120\\,K, and an SiO abundance $X$(SiO) of (0.5--1)$\\times$10$^{-6}$ are required to satisfy both the lower limit of the SiO $J$=1--0 brightness temperature and the 5--4/1--0 ratio. As no $J$=1--0 emission was detected at the innermost knots, B1 and R1, the 5--4/1--0 ratio must exceed 10: in order to reproduce such ratios, the same density range of (0.5--1)$\\times$10$^{7}$\\,cm$^{-3}$ is again required, but at a much higher gas kinetic temperature, in excess of 300--500\\,K\\@. These density and temperature ranges for B1 and R1 agree with those derived from the higher-$J$ transition SiO lines observed with a single-dish telescope \\citep{Nis02} and the high-$J$ CO lines observed with the ISO Long Wavelength Spectrometer \\citep{Gia01}.  \\subsection{The structure of the HH211 outflow} The rim-brightened shapes seen in the near-infrared and CO $J$=2--1 images suggest that the outflow lobes are the cavities filled with lower-density gas. Since the mean density of the molecular cloud core in which HH\\,211 lies is $\\sim$\\,4$\\times$10$^4$\\,cm$^{-3}$ \\citep{Bac87}, the density inside the cavities should be on the order of 10$^{4}$\\,cm$^{-3}$ or less. On the other hand, the SiO jet along the lobe axis has a much higher density of (0.5--1)$\\times$10$^{7}$\\,cm$^{-3}$. The strong H$_2$ emission observed in HH\\,211 suggests that the shocks in this outflow are nondissociative C-shocks \\citep[e.g.,][]{Hol97} and therefore, it is unlikely that this emission is coming from cavity gas which has been compressed by a factor of $>$500 from the ambient $\\sim$\\,10$^4$\\,cm$^{-3}$. Rather, it is much more likely that the primary jet itself is has an intrinsically high density or that dense gas in the protostellar disk or envelope is being entrained by the primary jet. While the FWHM of the dense SiO $J$=5--4 jet is $\\sim$\\,0.8$''$, the CO $J$=2--1 jet was seen to be somewhat broader at $\\sim$\\,1.5$''$ near the protostar and increasing up to $\\sim$\\,3$''$ further out, and with a density of $\\sim$\\,10$^5$\\,cm$^{-3}$, roughly 50 times lower than seen in the SiO \\citepalias{Gue99}. These results suggest that the jet has an axial structure, with the higher density ($>$10$^6$\\,cm$^{-3}$) gas close to the axis surrounded by a sheath of less dense gas 10$^5$\\,cm$^{-3}$.  The Hubble law velocity structure seen in Figure 3 at $> \\pm 5''$ from the source is predicted from shells driven by wide-angle winds \\citep[e.g.,][]{Shu91, Lee01}. On the other hand, turbulent entrainment models require an acceleration region near the star followed by a deceleration region \\citep[e.g.,][]{Rag93}, and bow shock entrainment models show a broad range of velocities near the bow tips \\citep[e.g.,][]{Lee01}; these do not explain the Hubble low shown in Figure 3. The highly-collimated morphology of the SiO $J$=5--4 jet can be explained as an on-axis density enhancement within the $X$-wind type of wide-opening angle wind \\citep{Li96, Sha02}. Since the wide-angle wind model has large radial velocity component, this will explain the large velocity dispersion observed near the base (at B1 and R1) combined with the jet-like density enhancement, which cannot be explained by any other entrainment models.  The density and velocity structures observed with the SMA suggest that the highly-collimated jet traced by the SiO $J$=5--4 is closely related to the primary jet. In particular, the innermost pair of knots, B1 and R1, have a high density of $>$10$^6$\\,cm$^{-3}$, a high temperature of $>$300--500\\,K, and a large velocity dispersion  of 30\\,km\\,s$^{-1}$, and quite plausibly represent the primary jet itself, immediately after being launched from the protostar/disk system."
    },
    "0512/hep-th0512209_arXiv.txt": {
        "abstract": "The string effective action at tree level contains, in its bosonic sector, the Einstein-Hilbert term, the dilaton, and the axion, besides scalar and gauge fields coming from the Ramond-Ramond sector. The reduction to four dimensions brings to scene moduli fields. We generalize this effective action by introducing two arbitrary parameters, $\\omega$ and $m$, connected with the dilaton and axion couplings. In this way, more general frameworks can be analyzed. Regular solutions with a bounce can be obtained for a range of (negative) values of the parameter $\\omega$ which, however, exclude the pure string configuration ($\\omega = - 1$). We study the evolution of scalar perturbations in such cosmological scenarios. The predicted primordial power spectrum decreases with the wavenumber with spectral index $n_s=-2$, in contradiction with the results of the $WMAP$. Hence, all such effective string motivated cosmological bouncing models seem to be ruled out, at least at the tree level approximation. ",
        "introduction": "The standard cosmological model, based on the general relativity theory, has been successfully tested until the nucleosynthesis era, around $t \\sim 1\\,s$, that is, at $T \\sim 1\\,MeV$\\cite{kolb,bellido}. At times before this era, it is only possible to speculate on how the universe has behaved. There are some claims that the spectrum of the anisotropy of the cosmic microwave background radiation (CMBR) has allowed to test the inflationary scenario, opening a window to energy levels of about $10^{15}\\,GeV$\\cite{riotto,dominik,kinney}. However, this last statement remains controversial\\cite{lue}: there is not, at this moment, a unique, complete, inflationary model, free of problems like transplanckian frequencies and fine tuning of fundamental parameters. In this sense the inflationary model remains a theoretical proposal asking for a full consistent formulation. Moreover, the fitting of the CMBR spectrum is made through a quite large number of parameters (up to ten), and more results, like the identification of the gravitational waves contribution to the CMBR anisotropy through the full detection of the polarization parameters of the CMB photons are necessary in order to surmount the degeneracy in the parameter space. Finally, neither the standard cosmological model nor the inflationary scenario address the question of the initial singularity, which is a major problem in the primordial cosmology. Hence, the very early universe remains an open area of research. \\par String theory\\cite{polchinski} is, at the moment, the most important candidate to a fundamental theory of nature unifying all interactions including gravity. At very low energy levels, string theory reduces to general relativity in its gravity sector. On the other hand, at higher energies, it predicts a deviation from the general relativity framework: the dilaton field couples non-minimally to the Einstein-Hilbert term; the Neveu-Schwarz sector exhibits also an axion field, which in four dimensions is equivalent to a scalar field which couples with the dilaton; in the Ramond-Ramond sector, scalar fields and gauge fields, minimally coupled to gravity and to the dilaton field, are present. Moreover, string theory is formulated in ten dimensions, and the reduction to four dimensions leads to the appearance of moduli fields associated with the dynamics of the extra dimensions. String theory itself may be incorporated in a more fundamental structure in $11$ dimensions ($M$ theory), $12$ dimensions ($F$ theory) and so on. All these complex structures lead, in principle, to a strong deviation from the standard cosmological scenario at very high energies. \\par The question we address here is how such a rich structure affects the evolution of the universe in its primordial phase. This problem has already been studied in many aspects in the literature\\cite{lidsey}. The pre-big bang scenario\\cite{gasperini}, the ekpyrotic model\\cite{ekp}, and the string gas cosmology \\cite{brand1} are some of them. However, all these attempts are plagued with some important difficulties: to construct a completely regular cosmological scenario, without any kind of singularity, which consequences can be successfully tested against observation. \\par General scalar-tensor systems which reduce to the particular string effective action for some values of the free parameters, have been considered in the literature \\cite{picco,patrick,kirill}. In reference \\cite{picco}, a $D$ dimensional lagrangian has been analyzed, with two free parameters $\\omega$ and $m$, which are connected with the coupling of the dilaton and the axion fields, respectively. The original configuration has been compactified considering flat, static extra dimensions. For a specific range of values of $\\omega$, scenarios with no curvature singularity have been obtained. However, the dilaton field is null initially, pointing out a singularity in the string expansion parameter, and thus rendering the effective model inadequate at that moment. In reference \\cite{patrick} fields from the Ramond-Ramond sector have been considered. This allows to avoid the singularity in the string expansion parameter, but only in the case of large negative values of $\\omega$, that is, $\\omega < - 3/2$. This implies the presence of negative energies when the lagrangian is re-expressed in the Einstein frame. The $D$ dimensional structure has been studied, in the vacuum case, in reference \\cite{kirill} and regular solutions have been obtained, again for large negative values $\\omega$. \\par In the present work, we perform a two-fold analysis. First, we explore another possibility with respect to some previous works (mainly references \\cite{picco,patrick}): we take into account the dynamics of the extra dimensions, and a radiative fluid whose presence is suggested by a maxwellian field in the Ramond-Ramond sector. Again, completely regular models can be obtained, this time for a larger range of values of the parameter $\\omega$, mainly with respect to the results of reference \\cite{patrick}: this parameter must still be negative, but in four dimensions it can be greater than $- 3/2$. The pure string case, which is characterized by $\\omega = - 1$, still leads to scenarios which are not completely regular; in fact this particular value can imply regularity only when the number of extra dimensions goes to infinity. On the other hand, even if $\\omega < - 1$, there are possible connections between $M$ and $F$ theories. We describe the features of these models, which display a bounce in the scale factor, while the dilaton and the moduli fields remain regular. We then test the new scenario obtained here against observations, as well as those obtained in \\cite{picco,patrick}, by computing the spectrum of scalar perturbations. The general result is that these models predict a decreasing power spectrum, which disagrees strongly with observations\\cite{teg}. As a general feature, all the scenarios lead to a spectral index for scalar perturbations that is around $-2$. This means that it is not possible to construct a realistic regular cosmological scenario based on the string (and more general frameworks) effective action at tree level unless, perhaps, some non trivial compactification scheme is introduced, or more complex configurations are considered like, for example, the (condensate) fermionic terms. \\par In the next section we describe the construction of the effective action. In section $3$, the background solutions are obtained. In section $4$ a perturbative analysis is carried out for the scalar modes and the power spectrum is computed. In section $5$ we present our conclusions. ",
        "conclusions": "We have studied cosmological models based on the string effective action at tree level. The dilaton, the axion, RR and the moduli fields were taken into account, besides gauge fields coming from the Ramond-Ramond sector. Two free parameters were introduced in the effective lagrangian, $\\omega$ and $m$, connected with the coupling of the dilaton and axion fields. These parameters were kept free in order to cover fundamental theories, like $M$ and $F$ theories, as well as $p$-branes configurations in the superstring theory. Regular solutions were found, but only for cases where $\\omega < - 1$, which excludes the pure string case. The new solutions found here, taking into account the moduli fields, exhibit compactification of the extra dimensions. The scale factor displays a bounce, exhibiting initially a contracting phase before entering the expansion phase. Asymptotically, a radiative phase is recovered allowing to match the solution with the standard cosmological model. At the same time, the string expansion parameter remains finite during the whole evolution of the universe. The effective gravitational coupling decreases with time what, at least, alleviate the hierarchial problem related to the characteristic energy scale of gravity. \\par The properties of the solutions found reveal that such \"string\" cosmological models can be candidates for describing the primordial phase of the universe. However, traces from this primordial phase can be compared with observations through the evaluation of the power spectrum of scalar fluctuations. This primordial power spectrum is inferred from the spectrum of the anisotropy of the CMBR. The observational results favor a flat spectrum. Here, we have computed the spectrum at the beginning of the radiative phase supposing that the initial fluctuations were formed in the beginning of the contraction phase, before the bounce, and that they were of quantum mechanical origin. The quantum fluctuations are compatible with the asymptotical behaviour in the beginning of the contracting phase. However, the final spectrum is strongly decreasing, in contradiction with the observational results, which favors at least a quasi-scale invariant spectrum. Hence, the regular solutions found, in spite of their nice features, are not candidate for a realistic primordial cosmological model, The same features are found for the regular \"anomalous\" solutions found in reference \\cite{patrick}.  In fact, the result $n_s=-2$ for the spectral index is not surprising. In reference \\cite{nelson}, the same spectrum has been obtained in the Einstein frame for a scalar field with negative kinetic energy together with radiation. Here the situation is very similar, see equations (\\ref{e14a},\\ref{e15a}), but with the presence of other fields. These other fields do not alter the bounce itself and, most important, the asymptotic behaviour. It seems that the intermediate phase they can be important is not relevant for the spectrum at large scales. Going to the Jordan's frame does not modify the result because the spectrum of $\\delta\\chi$ is negligible with respect to that of $\\Phi$ for small wave numbers. \\par These results seem to exclude such string motivated models at tree level. Nevertheless, it must be stressed that in the model developed here a quite simple compactification mechanism was considered: the internal space is flat, a $d$-dimensional torus. String theory admits many other kinds of compactifications. In particular, the Calabi-Yau manifolds are especially interesting since they can accommodate, in quite natural way, the gauge groups of the standard model of particle physics \\cite{green}. In this case, the effective model in four dimensions will be different to the one we have studied here. However, the negative results reported here indicate difficulties in constructing meaningful realistic cosmological models based on string motivated effective actions at tree level.    \\vspace{1.0cm}  {\\bf Acknowledgments:} We thank CNPq (Brazil) and CAPES (Brazil) for partial financial support. We thank also J\\'er\\^ome Martin and Patrick Peter for their criticisms and suggestions."
    },
    "0512/astro-ph0512584.txt": {
        "abstract": "\\noindent{\\it A brief history of the determination of the Hubble constant $H_{0}$ is given. Early attempts following \\citet{Lemaitre:27} gave much too high values due to errors of the magnitude scale, Malmquist bias and calibration problems. By 1962 most authors agreed that $75\\la H_{0} \\la 130$. After 1975 a dichotomy arose with values near 100 and others around 55. The former came from apparent-magnitude-limited samples and were affected by Malmquist bias. New distance indicators were introduced; they were sometimes claimed to yield high values of $H_{0}$, but the most recent data lead to $H_{0}$ in the 60's, yet with remaining difficulties as to the zero-point of the respective distance indicators. SNe\\,Ia with their large range and very small luminosity dispersion (avoiding Malmquist bias) offer a unique opportunity to determine the large-scale value of $H_{0}$. Their maximum luminosity can be well calibrated from 10 SNe\\,Ia in local parent galaxies whose Cepheids have been observed with HST. An unforeseen difficulty -- affecting {\\em all} Cepheid distances -- is that their P-L relation varies from galaxy to galaxy, presumably in function of metallicity. A proposed solution is summarized here. The conclusion is that $H_{0}=63.2\\pm1.3$ (random) $\\pm5.3$ (systematic) {\\em on all scales}. The expansion age becomes then (with $\\Omega_{\\rm m}=0.3, \\Omega_{\\Lambda}=0.7$) $15.1\\;$Gyr.} ",
        "introduction": "\\label{sec:01} % The present value of the Hubble parameter is generally called ``Hubble Constant'' ($H_{0}$). The {\\em present\\/} value requires minimum look-back-times; it is therefore to be determined at the smallest feasable distances and is adequately defined by \\begin{equation} H_{0}=\\frac{v}{r}\\;[\\mbox{km\\,s}^{-1}\\;\\mbox{Mpc}^{-1}], \\label{eq:01} \\end{equation} where $v=cz$, $z=\\Delta\\lambda/\\lambda_{0}$, and $r=$ distance in Mpc. As long as $z\\ll 1$, it is indicated to interprete $cz$ as a recession velocity because the observer measures the {\\em sum\\/} of the space expansion term $z_{\\rm cosmic}={\\cal R}_{0}/{\\cal R}_{\\rm emission} -1$ (${\\cal R}$ being the scale factor) and $z_{\\rm pec}$ caused by the density fluctuation-induced peculiar motions. At small $z_{\\rm cosmic}$ and in high-density regions $z_{\\rm pec}$ is not negligible. It is therefore mandatory to measure $H_{0}$(cosmic) at distances where $z_{\\rm cosmic}\\gg z_{\\rm pec}$ and outside of clusters. Any determination of $H_{0}$ must therefore compromise between two conditions: the smallest possible galaxy distances $r$ and a minimum influence of $z_{\\rm pec}$. The local Group is obviously useless for the determination of $H_{0}$ because it is probably gravitationally bound. The nearby Virgo cluster affects the local expansion field out to $\\sim\\!2000\\kms$ (see Section~\\ref{sec:06}). At $v\\sim3000\\kms$ the relative contribution of random velocities of field galaxies decreases  to less than $10\\%$, yet a volume of roughly similar radius has a bulk motion of $630\\kms$ with respect to the CMB. To be on the safe side it is therefore desirable to trace $H_{0}$ out to say $\\sim\\!20\\,000\\kms$. The expansion rate at this distance is for all practical purposes still undistinguishable from its present value.  The first spectra of galaxies and the measurement of their radial velocities by \\citet{Slipher:14} and later by M.~Humason and others was an epochal achievement. Today the observation of the redshifts needed for the calibration of $H_{0}$ is routine. The emphasis here lies therefore entirely on the determination of galaxy distances.  % ****************************************************************** %  2. The First Galaxy Distances % ****************************************************************** ",
        "conclusions": "\\label{sec:07} % In general astronomical distances depend on objects whose distances are already known and ultimately, with a few exceptions, on trigonometric parallaxes and hence on the AU. But methods of determining distances from the physics or geometry of some objects, without recourse to any other astronomical distance, are gaining increasing weight. Already the moving-atmosphere (BBW) method contributes to the calibration of the Galactic P-L relation of Cepheids. The single, intrinsicly accurate water maser distance of NGC\\,4258 \\citep{Herrnstein:etal:99} does not yet suffice for an independent calibration of the P-L relation \\citep[see][]{Saha:etal:05}. The recently improved expanding-atmosphere distance of SN\\,II 1999em \\citep{Baron:etal:04} agrees well with the Cepheid distance of its parent galaxy NGC\\,1637. \\citet{Nadyozhin:03} plateau-tail method for SNe\\,IIP yields $H_{0}=55\\pm5$ on the assumption that the $^{56}$Ni mass equals the explosion energy. Models of SNe\\,Ia yield $M_{\\rm bol}\\approx M_{V}=-19.5$ \\citep[][for a review]{Branch:98} in fortuitous agreement with the empirical value of $M_{V}=-19.46$.  Much promise to determine $H_{0}$ accurately lies in the Sunyaev-Zeldovich (SZ) effect and in gravitationally lensed quasars; extensive work has gone into both methods. The SZ effect yields typical values of $H_{0}=60\\pm3$, yet the systematic error is still $\\sim\\!\\pm18$ (\\citealt{Carlstrom:etal:02} for a review; see also e.g.\\ \\citealt{Udomprasert:etal:04}; \\citealt{Jones:etal:05}). Results from lensed quasars lie still in a wide range of $48<H_{0}<75$ \\citep[e.g.][]{Saha:03,Koopmans:etal:03,Kochanek:Schechter:04,York:etal:05}.  A strong driver to determine $H_{0}$ as accurately as possible comes from the CMB. The interpretation of its Fourier spectrum depends on at least twelve free parameters, several of which cannot be determined elsewhere. A simultaneous solution for all twelve parameters yields $H_{0}=66\\pm7$ \\citep{Rebolo:etal:04}. It would bring important progress for the understanding of the CMB spectrum if an independently determined value of $H_{0}$ could be used as a reliable prior.  The value of $H_{0}=62.3$ corresponds to an expansion age of $15.1\\;$Gyr in a flat $\\Lambda$CDM model with $\\Omega_{\\Lambda}=0.7$. This gives a sufficient time frame for the oldest ages in the Galaxy. Stellar-evolution models give for M\\,107 $14.0\\pm2.8$ \\citep{Chaboyer:etal:00} and for M\\,92 $13.5\\;$Gyr \\citep{VandenBerg:etal:02}. Models of the chemical evolution of the Galaxy yield ages of the actinides of $12.4-14.7\\;$Gyr \\citep{Thielemann:etal:87} or a U/Th age of $14.5\\pm2.5$ \\citep{Dauphas:05}. The emphasis has shifted over the last years to the Th/Eu dating of ultra-metal-poor giants. Typical results lie between $14.2\\pm3.0$ and $15.6\\pm4.0\\;$Gyr \\citep{Cowan:etal:99,Westin:etal:00,Truran:etal:01,Sneden:etal:03}. The ages are to be increased by the gestation time of the relevant objects; some galaxies may also have started their star formation before the Galaxy did. The present age determinations are not yet sufficiently accurate to set stringent limits on $H_{0}$, but the essential point is that none of these ages are significantly larger than allowed for by the expansion age.  \\newpage  \\noindent {\\em Acknowledgement.} My sincere thanks go to the Astronomische Gesellschaft which gave me the honorable and joyful opportunity for this report. I want to thank also my two honored teachers, Wilhelm Becker, himself a Karl-Schwarzschild Lecturer, and Allan Sandage, with whom I had the chance to collaborate for 43 years. Many other collaborators, too numerous to be named here, have helped me with their ideas and knowledge and friendship. Without the unbureaucratic support of the Swiss National Science Foundation during more than three decennia my scientific life would have been poorer.  % ****************************************************************** % Bibliography % ******************************************************************"
    },
    "0512/astro-ph0512209_arXiv.txt": {
        "abstract": "PSR B0355+54 is one of the handful of pulsars that has been observed with both \\cha\\ and \\xmm. The analysis of the archival data has revealed the pulsar and a $\\sim30$\\arcsec\\ compact nebula surrounding it. \\xmm\\ also has detected a trail which extends $\\sim$6\\arcmin\\ and similar to the compact nebula is also counter-aligned with the proper motion of the pulsar. The spectrum of both the pulsar and the extended emission are well described by an absorbed power-law model. The measured flux corresponds to an efficiency of converting the spin-down luminosity into X-rays in the 2-10 keV band of $\\sim$0.01\\% and $\\sim$1\\% for the pulsar and the extended emission, respectively. From the \\xmm\\ data we have detected pulsations at the expected radio period. The energetic and the extent of the extended emission can be explained by a Bow-shock formed by the motion of the pulsar through the interstellar medium. ",
        "introduction": "PSR B0355+54 is a middle age ($\\tau\\sim$0.56 Myr) pulsar with a spin-down energy loss rate of $\\dot{E}=4.5\\times 10^{34}$ \\ergsec, and a relatively close distance of $d=1.04^{+0.21}_{-0.16}$ kpc \\citep{chatterjee}. It was observed with \\ros\\ for a total of 19 ks and was detected with a non-thermal luminosity (2-10 keV) of $1.2\\times 10^{31}$ \\ergsec\\ \\citep{slane}. Taking into account the extended emission discussed in this paper, one should consider this as only an upper-limit to the pulsar's X-ray emission, considering \\ros\\ does not have a sufficient angular resolution to resolve these features.  Pulsar wind nebulae (PWNe) are useful probes of pulsars and their surroundings. In the case of older ($\\tau>10^{5}$ years) isolated neutron stars such as Geminga the PWN is thought to be the Bow-shock created by the pulsar proper motion through the interstellar medium (ISM) \\citep{deluca,caraveoa}. The distance, the energy loss rate, and the proper motion of some of the pulsars is well known through radio observations. The knowledge of such parameters make pulsar Bow-shocks particularly appealing to study. In this letter we report the discovery of an extended emission associated with the pulsar B0355+54, and investigate the possible implications. ",
        "conclusions": "The \\cha\\ and \\xmm\\ observations of PSR B0355+54 have proven the presence of extended emission associated with the pulsar. This emission has two components which are $\\sim$30\\arcsec\\ and $\\sim$6\\arcmin\\ long (see $\\S$~\\ref{observations}). The recently detected proper motion direction (see Table~\\ref{tab1}) being counter aligned with the trail together with the observed spectrum being well fitted with a power-law model suggests that the observed extended emission may be the result of synchrotron radiation from a wind-driven nebula supported by the ram pressure created by the proper motion through the ISM. This hypothesis have been found to be consistent with the size and energetics of the X-ray trails of Geminga \\citep{caraveoa} and PSR B1929+10 \\citep{beckera}.  There are two relevant time scales that need to be considered. The first is the $\\tau_{flow}$, the timescale for the passage of the pulsar over the length of its X-ray trail. The other one is $\\tau_{syn}$, the synchrotron lifetime of the radiation. This is given by the equation \\citep[see e.g.][]{gaensler} $$ \\tau_{syn} =39 B_{\\mu G}^{-3/2} \\left( \\frac{h\\nu}{\\mbox{keV}} \\right)^{-1/2} \\mbox{kyr}, $$ where $B_{\\mu G}$ is the magnetic field in the emission region in units of $\\mu$G. The first one of these can be estimated through the relation $\\tau_{flow}=r_{t}/v_{p}$, where $r_{t}$ is the extent of the trail which is measured to be $\\sim6$\\arcmin. This at a distance of 1.04 kpc corresponds to a linear distance of 1.8 pc.  The proper motion, $v_{p}$, is measured by \\citet{chatterjee} and is 61$^{+12}_{-9}$ km s$^{-1}$. This, together with the linear extent of the trail, gives the time it took PSR B0355+54 to travel across the trail as $\\sim 34$ kyr. Assuming these two time scales are comparable the estimate for the magnetic field around the emitting region is $B_{\\mu G}\\sim$1 for an X-ray photon of energy $h\\nu=1$ keV. We should note that this calculation ignores the possible inclination of the pulsar's proper motion with respect to the line of sight.  Assuming that the ISM in the vicinity of the pulsar carries a magnetic field that does not differ significantly from that of the average of our galaxy which is $\\sim 4$ $\\mu$G \\citep{beck}. This value should be the characteristic magnetic field for the trail where the X-ray emitting particles are no longer in the shocked region. For the CN, where the electrons are in the shocked region, we need to consider the change in density due to the shock and hence the change in the magnetic field. We can use standard hydrodynamical arguments to show that the density in the shock region is given by the relation $\\rho_{\\mbox{\\small{shock}}}=4\\rho_{\\mbox{\\small{ISM}}}$, which would imply a magnetic field of $B_{\\mbox{\\small{shock}}}=4B_{\\mbox{\\small{ISM}}}$ or $B_{\\mbox{\\small{shock}}}=16$ $\\mu$G, given the average ISM magnetic field. This value, given the large uncertainties, is in agreement with the value found here.  The termination shock radius is given by the balance of the ram pressure between the wind particles and the ISM, i.e. $$\\frac{\\dot{E}}{4\\pi c R_{s}^{2}}=\\frac{1}{2}\\rho v_{p}^{2},$$ where $\\rho$ is the density of the ISM and $v_{p}$ is the proper motion velocity. From here one can estimate the radius of the termination shock as $R_{s}\\sim3\\times 10^{6}\\dot{E}_{34}^{1/2}n^{-1/2}v_{p,100}^{-1}$ cm, where $n$ and $v_{p,100}$ are the number density in the ISM in units of cm$^{-3}$ and the pulsar space velocity in units of 100 km s$^{-1}$, respectively. Adopting $n=1$ cm$^{-3}$ the shock termination radius becomes 1.1$\\times 10^{17}$ cm. This corresponds to an angular size of $\\sim$7\\arcsec\\ at a distance of 1.04 kpc. This value is consistent with the data given the observational constraints.  In order to check the energy dependence of the extent of the CN we constructed two background-subtracted linear profiles along the proper motion direction, which contain photons from the intervals 0.2-2 keV and 2-10 keV (Figure~\\ref{radial}). For comparison also the point source contribution is plotted. Although the structure of the CN is not identical in both energy ranges the overall extent within the statistical uncertainties are equal. This is consistent with the identification of this region as the termination shock. This region is expected to show a constant X-ray spectrum across its extent. Due to the pressure difference between the regions ahead and behind the pulsar's motion the termination shock is not of uniform radius around the pulsar and is extended along the proper motion direction \\citep{bucciantini}. The larger structure seen by \\xmm\\ (trail) can be thought of as the emission originating from the wind particles that flow around the edge of the shocked region. The magnetic field in this region should be an order of magnitude less than the nebular magnetic field in the shocked region, if equipartition holds. The two components of the extended emission could be explained by the Bow-shock theory.   PSR B0355+54 is yet another pulsar with known proper motion and a collimated outflow in the opposite direction. Such flow, also observed in Geminga, could be interpreted as a pulsar jet \\citep{pavlov}. This is reminiscent of the argument proposed for the Vela pulsar \\citep{markwardt} where the outflow in opposite direction is proposed as the mechanism to boost pulsar velocities in contrast to birth-kicks."
    },
    "0512/astro-ph0512386_arXiv.txt": {
        "abstract": "Recent observations have discovered star formation activities  in the extreme outer regions of disk galaxies. However it remains unclear what physical mechanisms are responsible for triggering star formation in such low-density gaseous environments of galaxies. In order to understand the origin of these outer star-forming regions, we numerically investigate how the impact of dark matter subhalos orbiting a gas-rich disk galaxy embedded in a massive dark matter halo influences the dynamical evolution of outer HI gas disk of the galaxy. We find that if the masses of the subhalos ($M_{\\rm sb}$) in a galaxy with an extended HI gas disk  are as large as $10^{-3} \\times M_{\\rm h}$, where $M_{\\rm h}$ is the total mass of the galaxy's dark halo, local fine structures can be formed in the extended HI disk. We also find that the gas densities of some apparently filamentary structures can exceed a threshold gas density for star formation and thus be likely to be  converted into new stars in the outer part of the HI disk in some models with larger $M_{\\rm sb}$. These results  thus imply that the impact of dark matter subhalos (``dark impact'') can be important for better understanding the origin of recent star formation discovered in the extreme outer regions of disk galaxies. We also suggest that characteristic morphologies of local gaseous structures formed by the dark impact can indirectly prove the existence of dark matter subhalos in galaxies. We discuss the origin of giant HI holes observed in some gas-rich galaxies (e.g., NGC 6822) in the context of the dark impact. ",
        "introduction": "Star formation activities in the extreme outer regions of gas-rich disk galaxies have recently come to be discussed extensively not only  in the context of star formation laws in low-density environments  of galaxies but also  in the context of formation and evolution of disk galaxies (e.g., Ferguson et al. 1998a, b; Leli\\`evre \\& Roy 2000; Cuillandre et al. 2001; Martin \\& Kennicutt 2001; de Blok \\& Walter 2003; Gil de Paz et al. 2005; Thilker et al. 2005). Recent observational studies of M83 by the {\\it Galaxy Evolution Explore (GALEX)} have discovered UV-bright stellar complexes associated with filamentary HI structures in the extreme outer disk at $R \\sim 4R_{\\rm HII}$, where $R_{\\rm HII}$ corresponds to the radius where the majority of HII regions are detected (Thilker et al. 2005). Furthermore a number of small isolated HII regions have been recently discovered at projected distance up to 30 kpc from their nearest galaxy (e.g., NGC 1533) in the {\\it Survey for Ionization in Neutral Gas Galaxies (SINGG)} (e.g., Meurer 2004; Ryan-Weber et al. 2003).  It has been discussed whether the observed properties of these recent star formation activities  in the extreme outer HI regions of disk galaxies can be understood in terms of local gravitational instability within the HI disks (e.g., Ferguson et al. 1998a;  Martin \\& Kennicutt 2001). Ferguson et al. (1998a) suggested that a simple picture of the local gravitational instability can not explain self-consistently the observed radial distributions of HI gas and H$\\alpha$ flux of star-forming regions. Star formation (proven by H$\\alpha$ regions) in dwarf irregular galaxies (e.g., ESO 215-G?009) with extended HI disks is observed to occur in their very outer parts, where the gas densities are well below a critical density of star formation (e.g., Warren et al. 2004). Thus it remains unclear what can control the star formation activities in the extreme outer regions of disk galaxies.  Using numerical simulations of  galaxy formation  based on the  cold dark matter (CDM) model, Font et al. (2001) demonstrated that dark matter subhalos predicted in the CDM model can play a minor role in the heating of the disk owing to the very small number of subhalos approaching to the solar radius of 8.5 kpc. The sizes of HI disks of gas-rich galaxies are generally observed to be significantly larger than their optical disks with the sizes of $R_{\\rm s}$ (Broeils \\& van Woerden 1994) and some  fraction of low luminosity galaxies have HI gas envelopes extending out to 4--7 $R_{\\rm s}$ (e.g., Hunter 1997). No theoretical attempts have been made to investigate the dynamical impact of dark matter subhalos (hereafter referred to as ``dark impact'' for convenience) on {\\it the extended HI disks of galaxies}, though several numerical studies have already investigated the influences of triaxial halos with figure rotation and tidal galaxy interaction on the evolution of the extended HI disks (e.g., Theis 1999; Bekki \\& Freeman 2002; Masset \\&  Bureau 2003; Bekki et al. 2005a, b).  The purpose of this Letter is to propose that the dynamical interaction between dark matter subhalos and extended HI gas disks can be important for better understanding the origin of recent star formation observed in the extreme outer regions of disk galaxies. By using hydrodynamical simulations of the dark impact on galaxy-scale HI disks, we show that local fine structures (e.g., filaments and holes)  can be formed by the dark impact in the HI disks. We discuss whether star formation can occur in high-density regions of apparently  filamentary structures  formed by the dark impact. We  suggest that characteristic fine structures formed by the dark impact in a HI disk of a galaxy {\\it with apparently no interacting visible dwarfs close to the disk} can indirectly prove the presence of dark matter subhalos that frequently pass through the outer part of the HI disk. ",
        "conclusions": "Although the present study has suggested that the dark impact can be responsible for star formation in the outer HI gas disks of galaxies, it remains unclear what roles the dark impact has in the evolution of HI disks {\\it within stellar disks of galaxies.} Although the possibility of dark matter subhalos approaching galactic stellar disks is low (Font et al. 2001), we here suggest the following two possible roles of the dark impact. One is that the unique local structures of OB stars, young clusters, and super-associations, such as the galactic belt and the Gould belt in the Galaxy (e.g., Stothers \\& Frogel 1974), can result from the dark impact. The other is that giant HI holes without bright optical stellar counterparts (e.g., star clusters responsible for supernovae explosion that for giant HI holes) observed in some low-luminosity galaxies can be due to the dark impact. These speculative suggestions need to be investigated in a quantitatively way by our future high-resolution  simulations on the dark impact on inner HI gas disks of galaxies.   The present study has shown that the previous passages of dark matter subhalos through extended HI disks of galaxies can be imprinted on fine structures (e.g., filaments and holes) in the HI disks. The present model also predicts that if a galaxy-scale halo with an extended HI gas disk  has $\\sim 500$ subhalos,  $\\sim 8$ HI fine structures can be formed by dark impact for every 1 Myr. Recent HI observations have revealed that some fraction of galaxies have very extended HI disks (e.g., Hunter 1997; Warren et al. 2004), though the total number of galaxies whose  outer HI structures have been extensively investigated are  very small. Accordingly we suggest that if galaxy halos are composed of numerous subhalos, as the CDM model predicts, future high-resolution HI observations  on the fine structures of extended gas disks  for a statistically significant number of galaxies can {\\it indirectly} probe the existence of the subhalos and thereby provide some constraints on possible spatial distributions and kinematics of the subhalos. We also suggest that HI gas disks with apparently no interacting (visible) dwarf galaxies  close to the disks would be the best observational targets for proving subhalos in galaxies."
    },
    "0512/physics0512002_arXiv.txt": {
        "abstract": "Integral Field Spectroscopy (IFS) provides a spectrum simultaneously for each spatial sample of an extended, two-dimensional field. It consists of an Integral Field Unit (IFU) which slices and re-arranges the initial field along the entrance slit of a spectrograph. This article presents an original design of IFU based on the advanced image slicer concept~\\cite{Content1997}. To reduce optical aberrations, pupil and slit mirrors are disposed in a fan-shaped configuration that means that angles between incident and reflected beams on each elements are minimized. The fan-shaped image slicer improves image quality in terms of wavefront error by a factor~2 comparing with classical image slicer and, furthermore it guaranties a negligible level of differential aberration in the field. As an exemple, we are presenting the design LAM used for its proposal at the NIRSPEC/IFU invitation of tender. ",
        "introduction": "Introduction} Integral Field Spectroscopy (IFS) provides a spectrum simultaneously for each spatial sample of an extended, two-dimensional field. Basically, an IFS is located in the focal plane of a telescope and is composed by an Integral Field Unit (IFU) and a spectrograph. The IFU acts as a coupler between the telescope and the spectrograph by reformatting optically a rectangular field into a quasi-continuous pseudo-slit located at the entrance focal plane of the spectrograph. Therefore, the light from each pseudo-slit is dispersed to form spectra on the detector and a spectrum can be obtained simultaneously for each spatial sample within the IFU field.  The IFU contains two main optical sub-systems: the fore-optics and the image slicer. The fore-optics introduces an anamorphic magnification of the field with an aspect ratio of 1$\\times$2 onto the set of slicer mirrors optical surfaces. In such way each spatial element of resolution forms a 1$\\times$2~pixels image on the detector (i.e. the width of each slice corresponds to 2 pixels), which ensures correct sampling in the dispersion direction (perpendicular to the slices) and prevents under-sampling the spectra. This anamorphism can be avoided if under-sampling spectra is acceptable by science (for example, the SNAP~\\cite{Ealet2003} project) or if a spectral dithering mechanism is included in the spectrograph in order to recover for the under-sampling. The image slicer optically divides the anamorphic (or not) two-dimensional field into a large number of contiguous narrow sub-images which are re-arranged along a one-dimensional slit at the entrance focal plane of the spectrograph.  An image slicer is usually composed of a slicer mirror array located at the image plane of the telescope and associated with a row of pupil mirrors and a row of slit mirrors. The slicer mirror array is constituted of a stack of several thin spherical mirrors (called \"slices\") which \"slice\" the anamorphic field and form real images of the telescope pupil on the pupil mirrors. The pupil mirrors are disposed along a row parallel to the spatial direction. Each pupil mirror then re-images its corresponding slice of the anamorphic field on its corresponding slit mirror located at the spectrograph's focal plane (slit plane). The slit mirrors are also diposed along a row parallel to the spatial direction. Finally, each slit mirror which acts as a field lenses, re-images the telescope pupil (pupil mirrors) onto the entrance pupil of the spectrograph. The principle of an image slicer is presented in Fig.~\\ref{fig:principle}.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=7cm]{slicer_concept-schema-petit.ps} \\end{tabular} \\end{center} \\caption { \\label{fig:principle} The principle of an image slicer. The slicer mirror array, located at the image plane of the telescope, divides the entrance field of view (FOV) and re-mages the telescope exit pupils along a line on the pupil mirrors. Each pupil mirror then re-images its corresponding slice of the entrance FOV on its corresponding slit mirror located at the spectrograph's focal plane (slit plane). The reformatted FOV acts as the entrance slit in the spectrograph where all the slices are aligned as a pseudo long slit. } \\end{figure}  In order to improve image quality and/or reduce costs of image slicer, several adaptations have been recently developed: \\begin{itemize} \\item \\textbf{Catadioptric image slicer} where pupil and slit mirrors are replaced by dioptric elements~\\cite{Henault2004}. This allows to improve both image performances and costs while increasing the complexity of the opto-mechanical interface in cryogenic environment. Furthermore, dioptric elements present chromatic aberrations and result in a complex arrangement of pupil and slit mirrors since they are close together. \\item \\textbf{\"Staggered\" image slicer} where pupil mirrors are staggered in two rows instead of a single row~\\cite{Henault2004b}. This allows to place the pupil mirrors twice as far away while mainting the slit location so that the largest off-axis angles are reduced by a factor of two and then to improve image quality. \\item \\textbf{Image slicer using a flat facet slicer mirror array}. This image slicer~\\cite{Tecza2003} needs an additional spherical or cylindrical field lens located very close the slicer stack. A first look could conduct to think that this configuration optimally reduces cost manufacturing. But nevertheless, the fore-optics has to be more complex (re-imaging the pupil to a precise position after the slicer mirror array), and the progress in glass spherical slices manufacturing process permits to keep the cost differential small for a good system benefit. Furhtermore, such a configuration has the drawback of slightly decreasing the instrument throughput since an additional component is introduced in the optical layout. \\item \\textbf{Concentric image slicer} where the row of pupil mirrors, the row of slit mirrors and the collimator are disposed along concentric circles centered on the slicer mirror array~\\cite{Prieto2004,Dopita2004}. This configuration preserves aberrations in the field of view since angles are identical between each elements of each sub-slit channel. Thus this configuration is well adapted to diffraction limited instruments. \\end{itemize}  It is in the context of improving performances of such a complex system that we propose an original concept of image slicer called \"Fan-shaped\". As an exemple, we are presenting the design LAM used for its proposal at the NIRSPEC/IFU invitation of tender~\\cite{Prieto2003}. The fan-shaped image slicer is described in section~\\ref{sect:fanshaped}. Section~\\ref{sect:NIRSPEC} is devoted to the description of the whole IFU designed for the NIRSpec/JWST instrument and its performances. Section~\\ref{sect:comparison} compares performances of the fan-shaped image slicer with a classical image slicer design. ",
        "conclusions": "This article presents an original concept of image slicer called \"Fan-shaped\". Its design delivers good and homogeneous image quality over all IFU elements. We successfully apply its design to JWST/NIRSpec. Here we didn't discuss about manufacturing aspects since the performance aspects were preponderant however further investigations are under studying to drastically reduce costs and manufacturing aspects in such a design by preserving performances. Furthermore, a prototyping of the IFS (IFU and spectrograph) for the SNAP application is undergoing at LAM~\\cite{Aumeunier2005}."
    },
    "0512/astro-ph0512179_arXiv.txt": {
        "abstract": "If X-ray flashes (XRFs) and X-ray rich Gamma-ray Bursts(XRRGs) have the same origin with Gamma-ray Bursts (GRBs) but are viewed from larger angles of structured jets, their early afterglows may differ from those of GRBs. When the ultra-relativistic outflow interact with the surrounding medium, there are two shocks formed, one is a forward shock, the other is a reverse shock. In this paper we calculate numerically the early afterglow powered by uniform jet, Gaussian jet and power-law jet in the forward-reverse shock scenario. A set of differential equations are used to govern the dynamical evolution and synchrotron self-Compton effect has been taken into account to calculate the emission. In uniform jets, the very early afterglows of XRRGs and XRFs are significantly lower than GRBs and the observed peak times of RS emission are longer in interstellar medium environment. The RS components in XRRGs and XRFs are difficult to be detected. But in stellar wind, the reduce of very early flux and the delay of RS peak time are not so remarkable. In nonuniform jet(Gaussian jet and power-law jet), where there are emission materials on the line of sight, the very early light curve resembles isotropic-equivalent ejecta in general although the RS flux decay index shows notable deviation if the RS is relativistic(in stellar wind). ",
        "introduction": "Late time Gamma-ray Burst (GRB) afterglow emission has been well observed and studied since the first detection in 1997\\citep{p97nature, c97nature}. In contrast, afterglow shortly after or during the main GRB is still quite uncertain because of the lack of observation. This early afterglow may provide important information about initial parameters of the burst and shed light on explosion mechanism. SWIFT is able to observe multi-wavelength radiation rapidly after $\\gamma$ ray triggering. With new observations, the early stage should be investigated more profoundly. In the standard baryonic GRB fireball model, prompt GRB emission is produced by collisions of internal shocks while afterglow comes from interactions between burst-ejected materials and circum-burst medium. After the internal shock phase, as the fireball is decelerated by the circumburst medium, usually a pair of shocks develop \\citep{mr97apj, sp99apj}. One is a forward shock (FS) expands outward to heat the external medium, the other is a reverse shock (RS) propagates into the ejecta and heats this cold initial shell. The central energy is gradually transmitted to outer medium through shocked mediums and shocked ejecta. The early afterglow studied in this paper is during this transition stage.  \\citet{mr97apj} pointed out RS should emit detectable optical synchrotron photons. This prediction has been confirmed by the optical flash and the radio flare detected in GRB 990123\\citep{a99nature, k99apj, sp99apj, mr99mnras, np04mn}. The later more detailed investigation suggests that the RS region may be magnetized \\citep{f02chjaa, zkm03apj,pk04mnras}. The RS emission of GRBs in stellar wind medium has been discussed \\citep{cl00apj, wdhl03mnras, zwd05}. The RS emission powered by the magnetized outflow or neutron-fed outflow has been investigated in some detail \\citep{fww04aa, zk05apj, fzw05apj}. The pair-rich reverse shock emission has been discussed by \\citet{liz03apj} and \\citet{mkp05mnras}. Unfortunately, up to now, there are only a few additional candidates. They are GRB 021211 \\citep{f03apj, liw03apjl, w03aa, kp03mnras}, GRB 041219a \\citep{b05nature, fzw05apjl} and GRB 050525a \\citep{k05, sd05}. Whether the lack of RS emission events is due to observational limits or theoretical problems may be settled by further observation.  X-ray Flash (XRF) is an interesting phenomenon which resembles GRB. It is similar to GRB in many aspects except its lower peak energy ($\\sim$10keV) and lower isotropic energy ($\\sim10^{51}$ergs). Several models are proposed\\citep{dcb99apj, h01, hdl02mnras, b03aa, ldg05apj} among which the structured jet model is widely discussed. When a jet is ejected and propagating, observers on different viewing angles will see different phenomena. GRBs are observed near the jet center. XRFs may be detected at the edge of the jet. Between them are the X-ray rich GRBs (XRRGs) whose peak energies and isotropic energies are between those of GRBs and XRFs. Possible jet models for the XRFs and XRRGs include the off-beam uniform jet model \\citep{in01apj, yin02apj}, the Gaussian jet model \\citep{zm02apj, ldz04apj, zdl04apj}, and the power-law jet model\\citep{mrw98apj, jw04chjaa}. The very early afterglow powered by structured jets has been calculated by \\citet{fww04mnras} analytically. In this work, we present our numerical results.  The dynamical evolution is described in Section 2. The radiation calculation procedure is shown in Section 3. Afterglows from jets are prsented in Section 4. In section 5, our results are summarized with some discussions. ",
        "conclusions": "So far, there are several RS emission candidates reported in SWIFT era. One is GRB 041219a detected by INTEGRAL. \\citet{fzw05apjl} fitted its early IR afterglow with RS-FS emission model and found the RS region is magnetized. GRB 050525a may also be an possible candidate. \\citet{sd05} used RS-FS in standard scenario to fit the early optical bump and the parameters are reasonable.  However, among all the bursts targeted in optical band within few minutes after the prompt emission, only a small fraction of them have likely RS emission. This is conflicted with the theoretical estimation. It may be caused by the extinctions of the host galaxy. And as demonstrated in this paper, the overestimation of the $\\nu_m$, $\\nu_c$ may lead to a dimmer RS radiation. Up to now, the very early optical afterglows of XRFs have not been well detected, future observation may help us to modify the present theory.  In this paper, we calculate the early afterglow powered by various kinds of jets numerically. Dynamical evolution is solved from a set of differential equations. This is different from the analytical treatment of \\citet{fww04mnras} and gives similar but more exactly results. We find that the most unprecise estimation comes from $\\gamma_{34}-1$, which results in the overestimation of $\\nu_m$ and too rapid increasement of $\\nu_m$ before $t_\\times$ in previous treatment. This is especially significant in the ISM case since the RS is mildly relativistic in typical parameters taken here. At the same time, $\\nu_c$ was overestimated previously because of the ignoring of SSC effect. Considering these two factors, the peak flux of reverse shock should be dimmer, which may help us to explain the lack of detection of optical flashes in most GRBs. It also needs to be pointed out that electrons may be cooled not only by SSC process but also by external photons( \\eg, prompt emission) through inverse Compton scattering. Then $\\nu_c$ may be reduced further and lead to even fainter very early optical radiation.  Early afterglow from jets, both uniform and structured, are calculated. We find that the early afterglow varies significantly with different viewing angles and is dependent on the jet structure.  SWIFT XRT detection find that the early X-ray flare may be an common characteristic of GRB X-ray afterglow. \\citet{fw05} suggested that the RS synchrotron emission hardly can produce the very early X-ray flares. Our calculation confirms their results. We also tried to simulate it with RS SSC emission within a large parameters space but failed. The RS SSC X-ray emission is always lower and generally far lower than FS synchrotron emission."
    },
    "0512/astro-ph0512453_arXiv.txt": {
        "abstract": "The delay in light travel time along the line of sight generates an anisotropy in the power spectrum of 21cm brightness fluctuations from the epoch of reionization. We show that when the fluctuations in the neutral hydrogen fraction become non-linear at the later stages of reionization, the light-cone anisotropy becomes of order unity on scales $\\ga 50$ comoving Mpc. During this period the density fluctuations and the associated anisotropy generated by peculiar velocities are negligible in comparison. ",
        "introduction": "Fluctuations in the 21cm brightness from cosmic hydrogen at redshifts $z\\ga 6$ were sourced by the primordial density perturbations from inflation \\citep{loeb04,Bar05a,Bar05c} as well as by the radiation from galaxies \\citep{Scott,Madau,Fur04,Bar05b}. These two different components can be separated based on the angular dependence of the 21cm fluctuation power spectrum which is induced by peculiar velocities \\citep{Bha,Bar05a}. Uncertainties in the cosmological parameters lead to an additional apparent anisotropy due to the Alcock-Paczy\\'{n}ski effect; after accounting for constraints from the cosmic microwave background, this effect can still produce up to a $10\\%$ anisotropy which may be identifiable because of its particular angular structure \\citep{nusserAP,barkanaAP}. Several low-frequency arrays that could potentially detect the redshifted 21cm signal are being built around the globe, including the {\\it Primeval Structure Telescope} (web.phys.cmu.edu/$\\sim$past), the {\\it Mileura Widefield Array} (web.haystack.mit.edu/arrays/MWA), and the {\\it Low Frequency Array} (http://www.lofar.org).  In this work we examine an additional source of anisotropy in the 21cm power spectrum between the directions parallel and transverse to the line of sight. Previous discussions ignored the delay in light travel time (i.e., the ``light-cone'' constraint) along the line of sight [but see a rough estimate of the effect at the beginning of reionization in an appendix in \\citet{mcquinn}]. In particular, although the 21cm power spectrum is expected to be measured through averaging over three-dimensional volumes of finite radial extent, the power spectrum was previously evaluated at a fixed time slice of the universe, ignoring the fact that a fixed observing time implies a varying emission time as a function of distance from the observer. Here we evaluate the amplitude of the anisotropy which is sourced by this delay. This time-delay effect is not related to real causal effects or light-crossing times. Since two points at a given separation are in general seen at different redshifts, the correlation function, averaged over all such points, is affected by the change with time of the statistics of ionization (i.e., the distribution of H~II regions and their correlation with the underlying density field). ",
        "conclusions": "At the later stages of reionization, ionization fluctuations dominate the 21cm power spectrum (Figures~\\ref{fig:xiOfr1} and \\ref{fig:xiOfr2}) and the light travel-time delay generates a strong anisotropy. The signal is around a few mK on scales and redshifts where the time-delay anisotropy is large (Figure~\\ref{fig:xiOfz}), while noise and foreground-subtraction should still allow even first-generation 21-cm experiments to reach a sensitivity $\\sim 1$ mK on large scales \\citep{miguel,mcquinn}.  The strong line-of-sight anisotropy in the ionization fluctuations arises in our model since the cosmic mean neutral fraction changes rapidly with redshift and its fluctuations are highly non-linear. The scale and angular dependence of the light-cone anisotropy must be modelled carefully in order to separate the inflationary initial conditions from astrophysical effects \\citep{Bar05a}. Large-scale numerical simulations of reionization \\citep{Iliev05} could provide guidance for this difficult task, but fully self-consistent simulations with hydrodynamics and radiative transfer are required. Alternatively, the separation of the ``physics'' (i.e., the inflationary initial conditions) from the ``astrophysics'' may require observations at higher redshifts, when the neutral fraction is closer to unity and the 21cm power spectrum is dominated by density and peculiar velocity fluctuations."
    },
    "0512/hep-ph0512350_arXiv.txt": {
        "abstract": "\\noindent We study the effect on leptogenesis due to $B-L$ cosmic strings of a $U(1)_{B-L}$ extension of the Standard Model. The disappearance of closed loops of $B-L$ cosmic strings can produce heavy right handed neutrinos, $N_R$'s, whose $CP$-asymmetric decay in out-of-thermal equilibrium condition can give rise to a net lepton ($L$) asymmetry which is then converted, due to sphaleron transitions, to a Baryon ($B$) asymmetry. This is studied by using the relevant Boltzmann equations and including the effects of both thermal and string generated non-thermal $N_R$'s. We explore the parameter region spanned by the effective light neutrino mass parameter $\\tilde{m}_1$, the mass $M_1$ of the lightest of the heavy right-handed neutrinos (or equivalently the Yukawa coupling $h_1$) and the scale of $B-L$ symmetry breaking, $\\eta_{B-L}$, and show that there exist ranges of values of these parameters, in particular with $\\eta_{B-L} > 10^{11}\\gev$ and $h_1\\gsim 0.01$, for which the cosmic string generated non-thermal $N_R$'s can give the dominant contribution to, and indeed produce, the observed Baryon Asymmetry of the Universe when the purely thermal leptogenesis mechanism is not sufficient. We also discuss how, depending on the values of $\\eta_{B-L}$, $\\tilde{m}_1$ and $h_1$, our results lead to upper bounds on $\\sin\\delta$, where $\\delta$ is the the $CP$ violating phase that determines the $CP$ asymmetry in the decay of the heavy right handed neutrino responsible for generating the $L$-asymmetry. ",
        "introduction": "Present low energy neutrino oscillation data~\\cite{atmos_data, solar_data,kamland_data} are elegantly explained by the neutrino oscillation hypothesis with very small masses ($\\leq$ 1 eV) of the light neutrinos. Neutrinos can have either Dirac or Majorana masses. Small Majorana masses of the light neutrinos, however, can be generated in a natural way through the seesaw mechanism~\\cite{seesaw} without any fine tuning. This can be achieved by introducing right handed neutrinos, $N_R$'s, into the electroweak model which are neutral under the known gauge symmetries. The Majorana masses of these $N_R$'s are free parameters of the model and are expected to be either at TeV scale~\\cite{sahu&yajnik_prd.05} or at a higher scale~\\cite{type-I-bounds,type-II-bounds}. This indicates the existence of new physics beyond Standard Model (SM) at some predictable high energy scale.  The heavy right handed Majorana neutrinos are also an essential ingredient in currently one of the most favored scenarios of origin of the observed baryon asymmetry of the Universe (BAU), namely, the ``baryogenesis via leptogenesis\" scenario~\\cite{fukugita.86,luty.92,mohapatra.92,plumacher.96}. Majorana mass of the neutrino violates lepton number ($L$) and thus provides a natural mechanism of generating a lepton asymmetry in the Universe. Specifically, leptogenesis can occur via the $L$-violating, $CP$-asymmetric, out-of-equilibrium~\\cite{sakharov.67} decay of the $N_R$'s into SM leptons and Higgs. The resulting $L$-asymmetry is then partially converted to a baryon ($B$)-asymmetry via the non perturbative $B+L$ violating (but $B-L$ conserving) electroweak sphaleron transitions~\\cite{krs.85}. The attractive aspect of the leptogenesis mechanism is the link it implies between the physics of the heavy right handed neutrino sector and the experimental data on light neutrino flavor oscillation, thus making the scenario subject to experimental tests. Indeed, the magnitude of the $L$ (and thus $B$) asymmetry produced depends on, among other parameters, the masses of the heavy neutrinos, which in turn are related to the light neutrino masses via the seesaw mechanism. The mass-square differences amongst the three light neutrino species inferred from the results of neutrino experiments, therefore, place stringent constraints on the leptogenesis hypothesis.  A natural way to implement the leptogenesis scenario is to extend the SM to a gauge group which includes $B-L$ as a gauge charge. The heavy neutrino masses are then determined by the scale of spontaneous breaking of this gauge symmetry. Further, with $B-L$ a conserved gauge charge and $B+L$ anomalous, we can start with a net $B=L=0$ at a sufficiently early stage in the Universe. The observed $B$ asymmetry must then be generated only after the phase transition breaking the $B-L$ gauge symmetry.  It is well known that phase transitions associated with spontaneous breaking of gauge symmetries in the early Universe can, depending on the structure of the symmetry group and its breaking pattern, lead to formation of cosmic topological defects~\\cite{kibble.76,vilen&shell} of various types. In particular, the simplest choice for the $B-L$ gauge symmetry group being a $U(1)_{B-L}$, the phase transition associated with spontaneous breaking of this $U(1)_{B-L}$ in the early Universe would, under very general conditions, lead to formation of {\\it cosmic strings}~\\cite{kibble.76,vilen&shell} carrying quantized $B-L$ magnetic flux. These ``$B-L$\" cosmic strings can be a {\\it non-thermal} source of $N_R$'s whose decay can give an extra contribution to the $L$ and thereby $B$ asymmetry in addition to that from the decay of $N_R$'s of purely thermal origin (``thermal'' leptogenesis). This can happen in the following two ways:  First, since the Higgs field defining the $B-L$ cosmic string is the same Higgs that also gives mass to the $N_R$ through Yukawa coupling, the $N_R$'s can be trapped inside the $B-L$ cosmic strings as fermionic zero modes~\\cite{jeannerot.96}. Existence of zero energy solutions of fermions coupled to a Higgs field that is in a topological vortex string configuration is well-known~\\cite{jackiw.81,weinberg.81}, and has been studied in a variety of models allowing cosmic strings~\\cite{das.83,witten.85,Stern:1985bg,sdavis.00,vachaspati.01}. As closed loops of $B-L$ cosmic strings oscillate, they lose energy due to emission of gravitational radiation and shrink in size. Eventually, when the size of the loop becomes of the order of the width of the string, the string loop disappears into massive particles among which will be the $N_R$'s which were trapped inside the string as zero modes. Each closed loop would be expected to release at least one $N_R$ when it finally disappears, and the decay of these $N_R$'s would then give a contribution to the BAU through the leptogenesis route~\\cite{jeannerot.96}.  Second, collapsing, decaying or repeatedly self-intersecting closed loops of cosmic strings would in general produce heavy gauge and Higgs bosons of the underlying spontaneously broken gauge theory. In the context of cosmic strings in Grand Unified Theories (GUTs), the $CP$ asymmetric $B$-violating decay of the heavy gauge and Higgs bosons released from cosmic string loops would produce a net $B$ asymmetry~\\cite{pijush.82,brandenberger&co.91,riotto&lewis.94,pijush.98}. The sphaleron transitions would of course erase the $B$ asymmetry {\\it unless a net $B-L$ was generated}. If the strings under consideration are the $B-L$ cosmic strings, which can be formed at an intermediate stage of symmetry breaking in a GUT model based on $SO(10)$, for example, then the heavy gauge and Higgs bosons released from the decaying or collapsing $B-L$ cosmic string loops can themselves decay to $N_R$'s since the $N_R$'s have Yukawa and gauge couplings to the (string-forming) Higgs and gauge boson, respectively. The decay of these Higgs and gauge boson generated $N_R$'s can produce a net $B-L$ and thus contribute to the BAU through the leptogenesis route irrespective of the existence of zero modes of $N_R$'s on cosmic strings.  In a previous work~\\cite{bha_sahu_yaj.04}, we made a general analytical estimate of the contribution of the non-thermal $N_R$'s produced by $B-L$ cosmic string loops to BAU. It was shown there that, in order for the resulting $B$ asymmetry not to exceed the measured BAU inferred from the WMAP results~\\cite{spergel.03}, the mass $M_1$ of lightest right handed Majorana neutrino had to satisfy the constraint \\be M_1 \\lsim 2.4 \\times 10^{12}\\left(\\eta_{B-L}/10^{13} \\gev\\right)^{1/2}\\gev, \\label{M1-bound} \\ee where $\\eta_{B-L}$ is the $U(1)_{B-L}$ symmetry breaking scale. In the above mentioned analytical study we had (a) taken the $CP$ asymmetry parameter $\\epsilon_1$ to have its maximum value (see below), (b) not taken into account the contribution to BAU from the decay of the already existing thermal $N_R$'s, and (c) neglected all wash-out effects (see below) on the final $B-L$ asymmetry. Indeed, our aim there was to make a simple analytical estimate of the possible maximum contribution of the cosmic string generated non-thermal $N_R$'s to the measured BAU. It is, of course, clear that a complete analysis of leptogenesis in presence of cosmic strings can only be done by solving the full Boltzmann equations~\\cite{luty.92,plumacher.96} that include the non-thermal $N_R$'s of cosmic string origin as well as the already existing thermal $N_R$'s and take into account all the relevant interaction processes including the wash out effects. This is the study taken up in this paper.  The main results obtained in the present paper can be summarized as follows: First, we confirm the result, obtained earlier in our analytical study~\\cite{bha_sahu_yaj.04}, that $B-L$ cosmic string loops can give significant contribution to BAU only for $\\eta_{B-L}\\gsim 10^{11}\\gev$. Second, the numerical solution of the relevant Boltzmann equations in the present paper has enabled us to track the dynamical evolution of the contribution of the cosmic string generated non-thermal $N_R$'s to the final BAU. Specifically, we find that for sufficiently large values of $\\eta_{B-L}$ and for appropriate ranges of values of the other relevant parameters ${\\tilde m_1}$ and $h_1$ as detailed in sec.~\\ref{subsec:solving}, the effect of the cosmic string generated non-thermal $N_R$'s is to produce a late-time increase of the final value of BAU as compared to its value in absence of cosmic strings. This can be understood from the fact that while the thermal abundance of $N_R$'s decreases exponentially with decreasing temperature, the density of cosmic string generated $N_R$'s has a power law dependence (on temperature) inherited from the scaling behavior of the evolution of cosmic strings, leading to a domination of the string generated $N_R$'s over the thermal $N_R$'s at late times for sufficiently large values of $\\eta_{B-L}$. In such situations, we are required to place an {\\it upper} bound on the magnitude of $\\sin\\delta$, where $\\delta$ is the $CP$ violating phase that determines the $CP$ asymmetry in the heavy neutrino sector, in order not to overproduce the BAU.  The rest of our paper is organized as follows. In section \\ref{sec:brief}, we briefly review the standard thermal baryogenesis via leptogenesis scenario, and discuss the required lower bounds on the mass of the lightest right handed neutrino and the $B-L$ symmetry breaking scale. In section \\ref{sec:cosmic} we introduce closed loops of $B-L$ cosmic strings as non-thermal sources of $N_R$'s and write down the injection rates of these $N_R$'s due to the two main processes of disappearance of cosmic string loops. Section \\ref{sec:complete} is devoted to setting up and then solving the relevant Boltzmann equations for the evolution of the $B-L$ asymmetry, including the non-thermal $N_R$'s of cosmic string origin in addition to the usual thermal ones. The effects of the cosmic string generated non-thermal $N_R$'s on the evolution of the $B-L$ asymmetry are discussed. Finally, section \\ref{sec:conclusion} summarizes our main results. ",
        "conclusions": "} We have studied the effect of $B-L$ cosmic strings arising from the breaking of a $U(1)_{B-L}$ gauge symmetry, on the baryon asymmetry of the Universe. The disappearance of closed loops of $B-L$ cosmic strings can produce heavy right handed neutrinos, $N_R$'s, whose $CP$-asymmetric decay in out-of-thermal equilibrium condition can give rise to a net lepton ($L$) asymmetry which is then converted, due to sphaleron transitions, to a Baryon ($B$) asymmetry. We have solved the relevant Boltzmann equations that include the effects of both thermal and string generated non-thermal $N_R$'s. By exploring the parameter region spanned by the effective light neutrino mass parameter $\\tilde{m}_1$, the mass $M_1$ of the lightest of the heavy right-handed neutrinos (or equivalently the Yukawa coupling $h_1$) and the scale of the $B-L$ symmetry breaking, $\\eta_{B-L}$, we found that there exist ranges of values of these parameters, in particular with $\\eta_{B-L} > 10^{11}\\gev$ and $h_1\\gsim 0.01$, for which the cosmic string generated non-thermal $N_R$'s can give the dominant contribution to, and indeed produce, the observed Baryon Asymmetry of the Universe when the purely thermal leptogenesis mechanism is not sufficient. We have also discussed how, depending on the values of $\\eta_{B-L}$, $\\tilde{m}_1$ and $h_1$, our results lead to upper bounds on the $CP$ violating phase $\\delta$ that determines the relevant $CP$ asymmetry in the decay of the heavy right handed neutrino responsible for generating the $L$-asymmetry."
    },
    "0512/astro-ph0512335_arXiv.txt": {
        "abstract": "Angular differential imaging is a high-contrast imaging technique that reduces quasi-static speckle noise and facilitates the detection of nearby companions. A sequence of images is acquired with an altitude/azimuth telescope while the instrument field derotator is switched off. This keeps the instrument and telescope optics aligned and allows the field of view to rotate with respect to the instrument. For each image, a reference PSF is constructed from other appropriately-selected images of the same sequence and subtracted to remove quasi-static PSF structure.  All residual images are then rotated to align the field and are combined. Observed performances are reported for Gemini North data. It is shown that quasi-static PSF noise can be reduced by a factor $\\sim $5 for each image subtraction. The combination of all residuals then provides an additional gain of the order of the square root of the total number of acquired images. A total speckle noise attenuation of 20-50 is obtained for one-hour long observing sequences compared to a single 30s exposure. A PSF noise attenuation of 100 was achieved for two-hour long sequences of images of Vega, reaching a 5-sigma contrast of 20 magnitudes for separations greater than 8$^{\\prime \\prime}$. For a 30-minute long sequence, ADI achieves 30 times better signal-to-noise than a classical observation technique. The ADI technique can be used with currently available instruments to search for $\\sim 1$ M$_{\\rm{Jup}}$ exoplanets with orbits of radii between 50 and 300~AU around nearby young stars. The possibility of combining the technique with other high-contrast imaging methods is briefly discussed. ",
        "introduction": "Direct detections of very faint exoplanets and brown dwarfs near bright stars are essential to understand substellar formation and evolution around stars. This endeavor is now one of the major goals for next generation 10-m telescope instruments and future 30- to 100-m telescopes. The task is dauntingly difficult. The exoplanet or brown dwarf image is usually much fainter than the background from the brilliant stellar point spread function (PSF) image. Besides the Poisson noise, ground-based telescopes suffer from atmospheric turbulence that produces random short-lived speckles that mask faint companions. If these two limitations were the only ones, a simple solution would be to integrate longer to average these random noises and gain as the square root of the integration time. But observations have shown that, for integrations longer than a few minutes, the PSF noise converges to a quasi-static noise pattern, thus preventing a gain with increasing integration time \\citep{marois2003,marois2005,masciadri2005}. To achieve better detection limits, it is thus necessary to subtract the quasi-static noise using a reference PSF. Both ground- and space-based imaging are plagued with this stellar PSF calibration problem caused by imperfect optics and slowly evolving optical alignments. For ground-based imaging, subtraction of a reference PSF obtained from a star close to the target achieves a factor $\\sim 4$ of PSF noise attenuation, leaving residuals that are also quasi-static and thus severely limiting detection of fainter companions \\citep{marois2005}. For space telescopes that have a better PSF stability, like HST, a partial solution was found by subtracting two stellar images acquired during the same orbit with a different roll angle. This technique, called ``roll deconvolution'', successfully subtracts the stellar image by a factor 50 \\citep{schneider2003,fraquelli2004} but is also ultimately limited by PSF evolution. A similar technique, called angular differential imaging (ADI), can be used on ground-based altitude/azimuth telescopes to subtract a significant fraction of the stellar quasi-static noise and can potentially achieve detection limits that improve as the square root of the integration time.  In this paper, the ADI technique is described and its performance is analyzed using a simple analytical model and data from the Gemini Altair adaptive optics system and the NIRI near infrared camera. The PSF stability with Altair/NIRI is studied and its impact on ADI performances is discussed. Detection limits for three stars of our ongoing young nearby star survey are then shown. A comparison between ADI and normal imaging is also presented. Finally, the possibility of using ADI with other high-contrast imaging techniques is discussed. ",
        "conclusions": "The ADI observing technique was described and its performance using Altair/NIRI at Gemini was presented. It was shown that faint companions can be detected with better S/N when compared to classical observing techniques for a wide range of declinations. The ADI technique produces a reference PSF from the same target imaging sequence, removing the need to move to a nearby star for PSF calibration or to acquire sky exposures (for $H$-band imaging). Since the reference PSF is built using images acquired minutes apart, the reference PSF shows a good quasi-static speckle correlation.  The stability of the PSF plays a crucial role in ADI as it not only determines the speckle attenuation from the reference image subtraction but it also determines the regime in which the noise is attenuated with increasing observing time. It was reported that at Gemini with Altair/NIRI using 30s exposures, the PSF evolves on timescales of $\\sim 10-60$ minutes and the attenuation by subtraction of a reference image reaches $\\sim 2-6$ for short time intervals, achieving better speckle attenuation with better seeing conditions. The observations of HIP18859, for which a filter change during the sequence reduced significantly the speckle attenuation, underscore the importance of maintaining the optical setup fixed during the sequence. It was shown that the gain in S/N with increasing total observing time for separation greater than 2$^{\\prime \\prime}$ reaches more than 70\\% of the optimal case, indicating that the noise is mostly decorrelated between residual images for these separations. Typical residual speckle decorrelation time is of the order of a few minutes. The speckle noise residuals decorrelate faster for object having faster FOV rotation. In all cases, ADI guarantees a larger gain with longer observation sequences. To our knowledge, this is the first time that such behavior is clearly demonstrated for an acquisition and reduction technique designed for speckle attenuation. The wall raised by quasi-static speckles that prevents a gain with longer integration time for standard observing techniques \\citep{marois2003,marois2005,masciadri2005} can thus be removed by ADI. Comparison with a classical imaging technique shows that ADI achieves 30 times better speckle attenuation in 30 minutes integration time.  The noise attenuation obtained on Vega was 100, reaching a contrast of $\\sim$20 magnitudes at 8$^{\\prime \\prime}$ separation (63~AU). Observations of the young stars HD18803 and HD97334B yielded detection limits in difference of magnitude of 11.1-11.9 at 0.8$^{\\prime \\prime}$, similar to the SDI camera at VLT ($\\Delta m$ of $\\sim $11 at 0.8$^{\\prime \\prime}$), which is an optimized speckle suppression instrument. When combined to substellar models and estimated age for these stars, these observations show that ADI is well suited to search for jovian companions having a mass greater than 1-2 M$_{\\rm{Jup}}$ 50-300~AU away from nearby young stars. Finally, ADI could easily and advantageously be combined with SSDI, high-order AO and coronagraphy to improve the detection limits of exoplanets and brown dwarfs at all separations."
    },
    "0512/astro-ph0512103_arXiv.txt": {
        "abstract": "\\parbox{14cm} { The aim of this paper is to check if the models with realistic inhomogeneous matter distribution and without cosmological constant can explain the dimming of the supernovae in such a way that it can be interpreted as an acceleration of the Universe. Employing the simplest inhomogeneous model, i.e. Lema\\^{i}tre--Tolman model, this paper examines the impact of inhomogeneous matter distribution on  light propagation. These analyses show that realistic matter fluctuations on small scales  induce brightness fluctuations in the residual Hubble diagram of amplitude around $0.15$ mag, and thus can mimic  acceleration. However, it is different on large scales. All these brightness fluctuations decrease with distance and hence cannot explain the dimmining of supernovae for high redshift without without invoking the cosmological constant.This paper  concludes that models with realistic matter distribution (i.e. where variation of the density contrast is similar to what is observed in the local Universe) cannot  explain the observed dimming of supernovae without the cosmological constant. \\pb \\textbf{PACS Codes:} 98.65.Dx, 98.65.-r, 98.62.Ai } ",
        "introduction": "This paper examines  the supernova observations in order to  thoroughly estimate the influence of inhomogeneities on light propagation. Studies in this field proved that inhomogeneities can mimic the cosmological constant. However, this does not prove consistent with other astronomical observations. This paper  provides some  quantitative insight to matter fluctuations' influence in terms of the amplitude, $\\delta m$,  measured in the residual Hubble diagram.   The observations of supernovae are a powerful tool in   modern cosmology. Analyses  of the supernova brightness provide us with a reliable estimation of their distance from an observer. For this estimation to be satisfactory, all factors which might influence the observed supernova luminosity must be taken into account. In literature five factors are examined; namely, evolution of supernovae, dust absorption, selective bias, gravitational lensing, and cosmological models. Except for the last one they do not seem to be responsible for observed 'dimming' (for details see Refs. \\cite{Fil, Ris, Tor, Pr1, Pr2}). Analyses of supernovae in various homogeneous cosmological models imply a non-zero cosmological constant. However, similar analyses in inhomogeneous models have not been systematically studied.   The luminosity distance of supernovae without the homogeneity assumption is to be analysed as well. The luminosity distance in inhomogeneous models might differ from the FLRW results. To examine this issue  the Lema\\^{i}tre--Tolman model is employed. In this approach not only matter is inhomogeneously distributed but  the expansion of the space in not uniform as well. Results in the form of  the residual Hubble diagram provide us with the estimation of the impact of matter inhomogeneities.    The effect of inhomogeneous matter distribution on supernova observations was studied by many authors. For example employing the Lema\\^itre--Tolman model and the Taylor expansion of the luminosity distance in powers of the redshift C\\'el\\'erier \\cite{MNC1} showed that the inhomogeneities can mimic the cosmological constant. Iguchi, Nakamura and Nakao \\cite{INN} also used the Lema\\^itre--Tolman model to show that it is possible to fit supernova data without the cosmological constant. Similar results were obtained in the Stephani model by God\\l{}owski, Stelmach and Szyd\\l{}owski \\cite{GSS}. and in the Szafron model \\cite{Moffat}. Also models by Alnes, Amarzguioui and Gron \\cite{AAG} where the density is increasing with distance and models by Enqvist and Mattsson where expansion is decreasing with distance \\cite{EM} successfully fit supernova data without a need for the cosmological constant. There have also been other models proposed, in particular Swiss cheese models by Mansouri \\cite{Man} and Brouzakis, Tetradis and Tzavara \\cite{BTT1,BTT2}. For a review on explanation of the acceleration expansion without the cosmological constant the reader is referred to Ref. \\cite{MNC2}. The effect of inhomogeneities was also studied with aid of approximate methods \\cite{GCA,KMNR,KMR}. Recently Vanderveld, Flanagan and Wasserman \\cite{VFW} studied this issue using perturbation approach up to the second order in density fluctuations. Their results are similar to \\cite{GCA,KMNR,KMR} and indicate that the effect of inhomogeneity on the expansion of the Universe is small and thus cannot explain the apparent acceleration. However, because of the perturbation framework their results are valid only for small values of density fluctuations. Since the real density fluctuations in our Universe largely exceed $\\delta \\sim 1$ in order to draw reliable conclusion similar analyses should be conducted by employing exact solution of the Einstein equations.   These studies  have shown that matter inhomogeneities can explain the apparent acceleration of our Universe without employing the cosmological constant. This paper not only indicates that there are some specific conditions which enable explanation of the supernova dimming without $\\Lambda$ but also examines the influences of the realistic matter distribution on light propagation.  The structure of this paper is as follows:  Sec. \\ref{ltt} presents the Lema\\^itre--Tolman model; in  Sec. \\ref{obscon} presents observational constraints; Sec. \\ref{wlrfl} presents the residual Hubble diagram for models with  realistic density distribution but without the cosmological constant; Sec. \\ref{fita} presents results of fitting models to the supernova measurements without the cosmological constant; Sec. \\ref{ccl} presents the residual Hubble diagram for models with the realistic density distribution and with the cosmological constant. ",
        "conclusions": "This paper investigates the propagation of the light of supernovae  in the inhomogeneous Lema\\^itre--Tolman model. The inhomogeneous models are of  great flexibility and can fit the data without invoking the cosmological constant, which has been proved by Mustapha, Hellaby and Ellis \\cite{MHE}. Many authors before, like C\\'el\\'erier \\cite{MNC1} or recently Alnes, Amarzguioui and Gron \\cite{AAG} have proved that the matter inhomogeneities in the Lema\\^itre--Tolman model can mimic the cosmological constant and thus can be an alternative to dark energy. However, this paper indicates that the models which fit the supernova measurement without invoking the cosmological constant are very peculiar (see model 4 and 5, Sec. \\ref{fita}). These models have either a very peculiar expansion of the space (decreasing from the origin), or an unrealistic density distribution (increasing from the origin) or/and a  very large amplitude of the bang time function ($t_B(r)$). Introducing the {\\it Ockham's Razor} principle, it is more likely that the  Universe is accelerating rather than the conditions in our position in the Universe are so very special and extraordinary that they could be possibly responsible for the observed dimmining of the supernova brightness.  The results show that in the inhomogeneous Lema\\^itre--Tolman model the amplitude of brightness fluctuations observed in the residual Hubble diagram is significantly large for low redshifts of amplitude around  0.15 mag but it decreases for higher redshifts. Thus, for redshifts larger than $z \\approx 0.3$ these fluctuations are neglgible. All this may be the result of the evolution (as in the past the density fluctuations were smaller, and, consequently were of smaller influence on the brightness fluctuations). However, it is also possible that this fast decrease can be due to the symmetry restrictions. The Lema\\^itre--Tolman model assumes a spherical symmetry which puts too many constrains on the evolution and another parameters of the model. Therefore, it would be worth investigating the light propagation in the models which are both non-symmetrical and inhomogeneous. If in the inhomogeneous and non-symmetrical models the magnitude fluctuations do not decrease so fast, the observed scatter of supernova measurements might be partially possible to  explain.  The main conclusion of this paper is that matter inhomogeneities introduce the brightness fluctuations to the residual Hubble diagram of amplitude approximately  $0.15$ mag for low redshifts, and thus can mimic the acceleration on small scales. However, to explain the excess of faint supernovae without applying any special conditions (such as for instance peculiar expansion of the Universe) the cosmological constant has to be employed."
    },
    "0512/astro-ph0512429_arXiv.txt": {
        "abstract": "{We investigate the dependence of galaxy clustering on the galaxy intrinsic luminosity at high redshift, using the data from the First Epoch VIMOS-VLT Deep Survey (VVDS). The size (6530 galaxies) and depth ($I_{AB}<24$) of the survey allows us to measure the projected two-point correlation function of galaxies, $w_p(r_p)$, for a set of volume-limited samples up to an effective redshift $\\left<z\\right>=0.9$ and median absolute magnitude $-19.6< M_B < -21.3$. Fitting $w_p(r_p)$ with a single power-law model for the real-space correlation function $\\xi(r)=(r/r_0)^{-\\gamma}$, we measure the relationship of the correlation length $r_0$ and the slope $\\gamma$ with the sample median luminosity for the first time at such high redshift.  Values from our lower-redshift samples ($0.1<z<0.5$) are fully consistent with the trend observed by larger local surveys. In our high redshift sample ($0.5<z<1.2$), we find that the clustering strength suddenly rises around $M_B^*$, apparently with a sharper inflection than at low redshifts.  Galaxies in the faintest sample ($\\left<M_B\\right>=-19.6$) have a correlation length $r_0=2.7^{+0.3}_{-0.3}$ $h^{-1}$ Mpc, compared to $r_0=5.0^{+1.5}_{-1.6}$ $h^{-1}$ Mpc at $\\left<M_B\\right>=-21.3$. The slope of the correlation function is observed to correspondingly steepen significantly from $\\gamma=1.6^{+0.1}_{-0.1}$ to $\\gamma=2.4^{+0.4}_{-0.2}$. This is not observed either by large local surveys or in our lower-redshift samples and seems to imply a significant change in the way luminous galaxies trace dark-matter halos at $z\\sim 1$ with respect to $z\\sim 0$. At our effective median redshift $z \\simeq 0.9$ this corresponds to a strong difference of the relative bias, from $b/b* < 0.7$ for galaxies with $L < L*$ to $b/b* \\simeq 1.4$ for galaxies with $L > L*$. ",
        "introduction": "At the current epoch, luminous galaxies tend to be more clustered than faint ones (\\cite{davis}, \\cite{hamilton}, \\cite{iovino}, \\cite{maurogordato}, \\cite{benoist}, \\cite{willmer98}, Guzzo et al., 2000, \\cite{norberg}, \\cite{norberg02}, \\cite{zehavi}), with the difference becoming remarkable above the characteristic luminosity $L_*$ of the Schechter luminosity function. This effect is in general agreement with predictions from hierarchical models of galaxy formation (\\cite{white}, \\cite{valls}, Kauffmann et al. 1997, \\cite{benson}), in which bright galaxies are expected to occupy more massive dark matter haloes than fainter ones and these haloes are more strongly clustered than the overall distribution of dark matter (\\cite{kaiser84, mo-white, sheth-tormen}). If this is the case, the difference in clustering between faint and bright galaxies should become even more evident at high redshifts, where galaxy formation is supposed to be more confined to the highest peaks of the density field.  Understanding the relationship between galaxies and dark matter halos is one of the most difficult challenges of the theory in predicting the observed clustering of galaxies.  Over the last few years, {\\sl halo occupation models} have provided this connection in a phenomenological way, allowing one, e.g., to explain the detailed shape of the galaxy two-point correlation function (\\cite{zehavi04}, but see also \\cite{guzzo91}).  In these models, a statistically motivated recipe to describe galaxy formation determines the halo occupation distribution (HOD), specifying the probability $P(N|M)$ that a halo of virial mass $M$ contains $N$ galaxies of a given type, together with any spatial and velocity biases of galaxies (\\cite{kauff}, \\cite{benson}, \\cite{berlind03}, \\cite{kravtsov04}). This term (known as the {\\it one-halo component} of the correlation function) governs the behaviour of galaxy correlations on small ($<2$ $h^{-1}$ Mpc) scales, while at larger separations galaxy correlations are dominated by the gravitational clustering of virialized dark matter halos (the {\\it two-halo component}), with essentially no dependence on the more complex physics of the sub-dominant baryonic component.  Given cosmological parameters and a specified HOD, therefore, one can calculate any galaxy clustering statistic, on any scale (e.g., \\cite{abazajian}), either by populating the halos from an N-body simulation (e.g., \\cite{jing98}, \\cite{jing02}) or via analytical prescriptions (e.g., \\cite{peacock00}, \\cite{seljak00}, \\cite{marinoni02}, \\cite{cooray02}, \\cite{vdb03}). On the other hand, as it has been shown (\\cite{sheth2004}, \\cite{gao2005}, \\cite{harker2005}), there seems to exist a clear relationship between halo formation properties and halo clustering properties, which indicates that current HOD models may describe galaxy clustering at best approximately. Thus, observations of the relative clustering of galaxies with different intrinsic luminosities provide crucial constraints on HOD models.  The detailed luminosity dependence of clustering so far has been difficult to establish because of the limited dynamic range in luminosity for even the largest local galaxy redshift surveys (e.g. \\cite{norberg}). It is even more problematic to study this effect at redshifts significantly different to zero.  High redshift samples have been too small to allow sub-division into luminosity classes.  An additional complication relates to evolution of the overall luminosity function: galaxies become brighter on average going back in time, thus comparison of high-redshift measurements to local values requires accurate knowledge of the evolution of the global luminosity function.  The VIMOS-VLT Deep Survey (VVDS) provides us with unique information to address these issues in detail.  A first investigation of how the non-linear bias between galaxy and matter evolves with redshift for different luminosity classes has been presented in Marinoni et al. (2005). An analysis of the evolution of clustering of galaxies has been presented in Le F\\`evre et al. (2005a), and the evolution of the dependence of clustering on spectral types has been discussed by Meneux et al. (2005). In this paper we use the same VVDS first-epoch data to measure in more detail the dependence of galaxy clustering on luminosity at $\\left<z\\right>\\simeq 0.9$, and compare it to local values from 2dFGRS and SDSS. We describe the VVDS catalog and the construction of volume limited samples in Section 2. Section 3 presents the methods to estimate and retrieve the best-fit parameters for the real-space correlation function. We present our results on the projected correlation function in Section 4, while the comparison to existing local surveys, together with a discussion of the results is given in Section 5.  Throughout this paper we use a Concordance Cosmology with $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. The Hubble constant is normally parameterised via $h=H_0/100$ to ease comparison with previous works, while a value $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ is used when computing absolute magnitudes. All correlation length values are quoted in comoving coordinates. ",
        "conclusions": "The projected correlation functions that we have measured from our set of volume limited sub-samples of the VVDS are in general fairly well fitted by a single power-law in the range $0.1 \\le r/h^{-1} {\\rm Mpc} \\le 10$, both for the low-redshift and high-redshift samples.  This allows us to use variations in $r_0$ and $\\gamma$ to characterize the global dependence of clustering on luminosity at high redshift and compare it to similar low-redshift results.  The observed behaviour has strong implications for HOD models, as it directly impacts on any recipe for populating dark matter halos at high redshift. Deviations from the power-law shape, although extremely interesting for constraining  HOD models (e.g. \\cite{zehavi}) are not analyzed in this paper and will be the subject of future work.  We observe that for median redshifts $z\\simeq 0.9$ ($0.5<z<1.2$) the clustering length has a low, nearly constant value $r_0\\simeq 2.9$ $h^{-1}$ Mpc for luminosities $M_B<M_B^*$, rising suddenly for $M_B>M_B^*$ and approaching values $r_0\\simeq 5$ $h^{-1}$ Mpc similar to those of local galaxies with comparable luminosity (relatively to the characteristic value $M_B^*$, \\cite{norberg}). This behaviour is consistent with general predictions by hierarchical models of galaxy formation (\\cite{benson}), where luminous galaxies are more confined to the peaks of the large-scale density field, going back in redshift, simply due to the higher bias of the parent halos.  Another important result of this work is the clear detection of a systematic steepening of the high-redshift correlation function for absolute magnitudes brighter than $\\sim M_B^*+0.5$ (Figure \\ref{gamma}). This kind of behaviour is in general not seen either in our closer $z \\sim 0.4$ sample or in the large local surveys, although Zehavi et al. (2005) do detect an increase of $\\gamma$ in their most luminous volume-selected sample. A similar trend with increasing (UV) luminosity has been recently observed for a population of Ly-break galaxies at $z \\ge 4$ in the Subaru Deep Field (\\cite{kashikawa}).  These authors are able to reproduce their observed relationship with a HOD model in which they introduce multiple LBGs into massive dark matter haloes. This amplifies the clustering strength at small scales, steepening the correlation function.  A similar interpretation could be applied to our data. Finally, the observed relative bias of galaxies at high redshift provides evidence for a clear difference in the clustering properties of galaxies fainter or brighter than the characteristic luminosity: sub-samples with $M_B \\lesssim M_B^*$ behave in a way that is very similar to local samples while the relative bias of samples with $M_B \\gtrsim M_B^*$ remains significantly lower.  Results presented in this paper show that there is a significant redshift evolution of the luminosity dependence of both the normalization and slope parameter of the galaxy correlation function. This specific observation can provide an important test of galaxy formation models, constraining in particular the multiplicity of luminous galaxies within massive halos at $z=1$.  NOTE ADDED IN PROOF: In a parallel paper \\cite{coil05} perform a similar measurement at $z \\sim 1$ using the DEEP-2 survey. Although they explore a narrower range in median luminosities, they also detect a first hint of the steepening of $w_p(r_p)$ above $M_B^*$ and a rise of the correlation length with luminosity. Considering cosmic variance and the different selection function (unlike the VVDS, DEEP-2 is not a purely magnitude-limited survey), the overall results from these two data sets are thus in good agreement."
    },
    "0512/astro-ph0512273_arXiv.txt": {
        "abstract": "We present a short atlas illustrating the unusual Ca {\\sc ii} K-line profiles in upper main sequence stars with anomalous abundances. Slopes of the profiles for 10 cool, magnetic chemically peculiar (CP) stars change abruptly at the very core, forming a deep ``nib.\" The nibs show the same or nearly the same radial velocity as the other atomic lines.  The near wings are generally more shallow than in normal stars.  In three magnetic CP stars, the K-lines are too weak to show this shape, though the nibs themselves are arguably present.  The Ca {\\sc ii} H-lines also show deep nibs, but the profiles are complicated by the nearby, strong H$\\epsilon$ absorption. The K-line structure is nearly unchanged with phase in $\\beta$ CrB and $\\alpha$ Cir.   Calculations, including NLTE, show that other possibilities in addition to chemical stratification may yield nib-like cores. ",
        "introduction": "Abundance studies of chemically peculiar (CP) stars of the upper main sequence are usually based on classical, plane-parallel atmospheres.  Until recently, there was little evidence that this assumption would lead to significant errors, apart from the obvious case of spectrum variables.  Even in these cases, it was common to assume plane-parallel structure for localized regions (abundance patches) of the photosphere.  Recent work on the ionization equilibrium \\cite{ryab04} and the Balmer profiles \\citep{cowl01} of cooler CP stars in the magnetic sequence has drawn attention to significant departures from classical atmospheric structure.  Decades ago, spectroscopists had noted another indication of departures from a classical atmosphere in the Ca {\\sc ii} K-lines.  % \\cite{baba58} describes very complicated K-line profiles in a number of CP stars.  In a few cases, variable profiles suggested to him ``ejection of ionized clouds or streams such as the sun produces.\"   Analytical spectroscopists did little more than note these remarks at the time.  Most of the interest in CP star focused on their peculiar abundances and attempts, for example, by \\citet{conti65} and \\citet{vveer66} % to explain the abundance anomalies in Am stars by unorthodox models had been unsuccessful.  New, high-resolution observations make it appropriate to take a closer look at the K-lines of CP stars.  We confine ourselves to  spectra of cooler magnetic stars with minimal rotational broadening.  We find that the K-lines of these stars generally exhibit a ``wing-nib anomaly'' (WNA).     \\placefigure{fig:one}  Two forms for the  WNA are shown schematically in Fig.~\\ref{fig:one}. In the upper part of the figure, the wings show a gently changing slope, which appears to nearly level off prior the deep minimum.  In the lower part of the figure, the wings show an approximately constant slope beyond the deep nib.  In both cases, the region outside of the nib has a higher intensity than is observed for normal stars, or the Am star HR 1353. Individual examples are cited in the captions of figures to follow.    The Ca II H-line also shows a sharp, deep nib, but the structure is complicated by the presence of the strong Balmer H$\\epsilon$ line.   We shall not discuss the H-line in this paper. ",
        "conclusions": "\\placefigure{fig:bugger} It is our purpose to present an observational atlas and not a new interpretation of the wing-nib anomaly.  We do not dispute stratification models.  They are demonstrably capable of producing nibs. Nevertheless, a few remarks are appropriate.  The wing-nib and core-wing (Balmer lines) anomalies show that the atmospheres of at least some of the cool CP stars depart significantly from the classical models.  This was suggested decades ago as an alternate explanation to abundance anomalies---a way to explain the peculiar spectra with normal abundances. The objections were dismissed forcefully by \\citet{sarg66}.  % Briefly, he argued that the CP stars had normal colors, curves of growth, ionization temperatures, and Balmer profiles. His arguments were valid at that time though they are untenable today.  Nevertheless, his conclusion is widely accepted, that the photospheres of magnetic CP stars are chemically anomalous.  Traditional abundance studies were based on classical, plane-parallel atmospheres, with a uniform chemical composition. \\citet{bab92} showed that a stratified abundance structure would produce a Ca K-line nib while Ryabchikova et al. (2002) used a chemically stratified model similar to Babel's. They also obtained a K-line nib, though they did not fit it precisely.  They found certain other atomic lines that could be fit more satisfactorily with the assumption of abundance stratification than with a traditional model.  Much new work has been devoted to this phenomena. Indeed, \\citet{dwor05} % described the session of IAU Symposium 224 \\citep{zver05} on diffusion by saying ``the keword ...was `stratification.'\" \\citet{kochu02} showed that the core-wing anomaly of the Balmer lines might be explained by an ad hoc variation in the $T$--$\\tau$ relation. We have made LTE calculations to show that wing-nib structure of Ca {\\sc ii} K may also be obtained with a modified $T$--$\\tau$.  Most of the calculations relevant for the WNA have been made in LTE.  We have begun calculations with the Kiel non-LTE code (Steenbock \\& Holweger 1984) that will be reported in more detail elsewhere. We use the updated version of the Ca model atom by \\citet{wat85} \\citep{stur93}. The core of the Ca {\\sc ii} K-line forms in layers of the atmosphere, where both involved levels show only very moderate deviations from LTE ($<0.05$ dex). Preliminary results show that these small quantitative differences from LTE in the K-line core vary somewhat with effective temperature. As an example, we present here a cool atmosphere (T$_{\\rm eff}=6750$~K, $\\log g=4.0$) with an ad-hoc modified $T$--$\\tau$ relation (see inset of Fig.~\\ref{fig:bugger}). The atmosphere temperature was raised to $T\\sim 6000$~K over the optical depth range $-4.5 < \\log \\tau < -0.5$. Fig.~\\ref{fig:bugger} shows the resulting LTE (solid line) and NLTE (dashed line) Ca\\,{\\sc ii}\\,K profiles. NLTE leads to a $\\sim$ 5\\% weaker core, but does not affect the main profile and the wings; thus, it would not lead to significant abundance differences. At present it seems that a fully non-LTE approach could modify details of a workable model, but would not require qualitative changes.   Neither of the above modifications takes account of the influence of magnetic pressure on atmospheric structure. Thus, it is reasonable to explore a third plausible departure from a classical model atmosphere, based on a consideration of magnetohydrodynamical effects.  We lack the information necessary to take such effects explicitly into account.  One possibility is that they simulate a model with gravity that decreases outward \\citep{conti65, valy04}.  Since a model with reduced gravity in the upper layers would have lower pressured and therefore reduced atomic absorption, the effect should resemble that of stratified models.  We have made provisional calculations showing that in LTE, such a model is indeed capable of yielding a K-line nib.  Thus far, the modified-gravity models that give K-line nibs have not yielded core-wing anomalous Balmer profiles."
    },
    "0512/astro-ph0512051_arXiv.txt": {
        "abstract": "A homogeneous sample of $\\sim$2200 low redshift disk galaxies with both high sensitivity long-slit optical spectroscopy and detailed \\iband\\ photometry is used to construct average, or template, rotation curves in separate luminosity classes, spanning 6 magnitudes in \\iband\\ luminosity. The template rotation curves are expressed as functions both of exponential disk scale lengths \\rd\\ and of optical radii \\ropt, and extend out to 4.5--6.5 \\rd, depending on the luminosity bin. The two parameterizations yield slightly different results beyond \\ropt\\ because galaxies whose \\Ha\\ emission can be traced to larger extents in the disks are typically of higher optical surface brightness and are characterized by larger values of \\ropt/\\rd. By either parameterization, these template rotation curves show no convincing evidence of velocity decline within the spatial scales over which they are sampled, even in the case of the most luminous systems. In contrast to some previous expectations, the fastest rotators (most luminous galaxies) have, on average, rotation curves that are flat or mildly rising beyond the optical radius, implying that the dark matter halo makes an important contribution to the kinematics also in these systems. The template rotation curves and the derived functional fits provide quantitative constraints for studies of the structure and evolution of disk galaxies, which aim at reproducing the internal kinematics properties of disks at the present cosmological epoch. ",
        "introduction": "\\label{s_intr}  The determination of the rotational characteristics of galaxy disks is fundamental to understanding the role that dynamics plays in galaxy formation and evolution over cosmic time. Over the last decade, a large number of studies have addressed the issue of disk formation and evolution in current hierarchical cosmologies by means of analytical or semi-analytical theoretical models \\citep[e.g.,][]{dss97,mo98,mao98,mm04,afh98,fa00,vdb01,vdb02} and numerical smooth particle hydrodynamic simulations \\citep[e.g.,][]{vdb02b,ab03,gov04}. Other groups have investigated more specifically the characteristics of the density profiles of dark matter halos, which directly determine the observable rotation curves (RCs) of present-day disks, using $N$-body simulations \\citep[e.g.,][]{nfw96,nfw97,bul01a,bul01b}. A detailed characterization of the kinematic properties of observed disk galaxies at low redshift is of key importance to validate such models and place constraints on the properties of dark matter halos. In this context, special attention has been recently devoted to dwarf and low surface brightness galaxies, i.e. systems with an internal kinematics that is thought to be dominated by dark matter even at small radii, thus allowing to best probe the inner structure of their parent halos (e.g., \\citealt{sgh05} and references therein). For high surface brightness galaxies, the kinematics is easier to measure, but the presence of a more prominent stellar disk complicates the analysis of the RC, whose decomposition into a dark matter and a disk component is generally not unique \\citep[see, for instance,][and references therein]{jvo03,dcd05}.  The derivation of {\\em average} RCs representative of the kinematic properties of a large number of local spirals is also important on its own, to provide a standard reference against which similar, distant samples can be compared. Furthermore, statistical descriptors of the amplitude and radial variation of rotation velocity within disks would benefit all the applications that rely on the assumption of a RC model, such as determinations of rotational widths from spectroscopic data that lack velocity or spatial resolution. Examples of the latter include internal kinematic studies of intermediate redshift spirals through spatially resolved optical spectroscopy, where synthetic RCs are needed in order to estimate the full velocity width \\citep[e.g.,][]{vog97,sim99,zie03,boh04}.  Beginning perhaps with \\citet[]{rr73}, many previous works have explored the variation in rotation curve form among galaxies. As larger samples of rotation curves have been amassed, attempts have been made to produce more quantitative descriptions of empirically--derived average, or {\\em template} RCs. Notably, average RCs binned by luminosity and with radial distances expressed in units of the optical radius \\ropt\\ defined as the radius encompassing 83\\% of the total integrated light were previously presented by \\citet[hereafter PSS96]{pss96} for a sample of 616 spiral galaxies. This definition of \\ropt\\ is equivalent to 3.2 \\rd\\ for a pure exponential disk. PSS96 showed that their average RCs were well described by an analytical form, obtained as the combination in quadrature of a dark matter halo and an exponential disk, with fit coefficients depending on galaxy luminosity only. Based on these results they claimed the existence of a universal relation, the {\\em universal rotation curve} (URC), whereby the shape and amplitude of an {\\em observed} RC at any radius is completely determined by the galaxy luminosity, as originally proposed by \\citet{ps91}. According to the URC predictions, the most luminous systems have RCs that peak at $\\sim$0.8 \\ropt\\ and decline beyond that radius, implying a minor contribution of the dark matter halos to the internal kinematics of those galaxies. It should be also noticed, however, that PSS96's average RC tracings beyond \\ropt\\ are based on linear extrapolations of outer velocity gradients as functions of luminosity mostly obtained from \\hi\\ data. Several studies have discussed the inadequacy of the URC parameterization \\citep[e.g.,][]{cou97,ver97,wil99,gma04}.  The purpose of this paper is to provide a set of template RCs in bins covering a wide range of galaxy luminosity and based on an extensive set of high sensitivity long-slit optical RCs, for which a homogeneous body of high quality \\iband\\ photometry is also available. Our intent is not to propose a new universal relation, but rather to obtain a more reliable characterization of the average RCs of late type, high brightness disk galaxies by taking advantage of a much larger data set, thereby negating the need for extrapolation within the optical disk. In order to allow a direct comparison with PSS96's average RCs and URC models, we provide template RCs parameterized as functions of \\ropt. Furthermore, we also present a set of templates which parameterize the distance from the center of a galaxy by means of the exponential disk scale length \\rd, a quantity naturally adopted in models of galaxy disks.  This work makes use of a large spiral galaxy photometric and spectroscopic dataset, dubbed SFI++ (see \\S~\\ref{s_sample}), compiled to investigate the internal kinematics of disk galaxies at low redshifts, as well as to study the peculiar velocity field in the local universe via application of the Tully-Fisher \\citep[TF;][]{tf77} method. Details on the extraction of related quantities from long-slit spectra, including our RC folding technique and algorithm for the measurement of rotational velocities for TF applications, can be found in \\citet[hereafter CHG05]{chg05}.  This paper is organized as follows. In \\S~\\ref{s_sample} we describe our data sample and illustrate the procedure used to derive template RCs. The results obtained using the parameterizations in terms of exponential disk scale lengths or optical radii are presented and compared in \\S~\\ref{s_res}. In order to explain systematic differences between the outer slopes of the two sets of template RCs, the correlation between \\ropt/\\rd\\ ratio and RC extent observed for the galaxies in our sample is also investigated in some detail. The impact of internal extinction on the inner slopes of RCs is analyzed in \\S~\\ref{s_ext}, where template RCs are computed for separate intervals of galaxy inclination. The comparison with previously published results, most notably the URC models, is discussed in \\S~\\ref{s_urc}, and our conclusions summarized in \\S~\\ref{s_concl}. A value for the Hubble constant of $\\rm H_0=70$ \\kmsm\\ is assumed throughout this work. ",
        "conclusions": "Taking advantage of a large sample of long-slit optical RCs with available \\iband\\ photometry, we have reinvestigated the dependence of RC shape on galaxy luminosity. We determined {\\em template} relations by fitting a function, the Polyex model (eq. \\ref{eq_polyex}), to the average RCs calculated in 10 luminosity classes, spanning the \\iband\\ absolute magnitude range [$-$24.0,$-$18.0]. After ascertaining the effect of internal extinction on our results, we excluded the 472 most edge-on systems ($\\rm i\\geq$ 80\\deg) from the analysis. The template RCs are expressed as functions both of exponential disk scale lengths \\rd\\ and of optical radii \\ropt, the latter to allow a direct comparison with PSS96's average RCs and URC predictions. When scaled with \\rd, the template relations show a smooth transition of RC shapes with increasing \\iband\\ luminosity, with the most luminous systems being characterized by steeper inner velocity rises and flatter outer slopes. The \\ropt\\ templates show analogous variations of RC amplitudes and inner slopes with \\iband\\ luminosity, but their outer slopes are nearly constant. A direct comparison of the two sets of templates shows that the results are very similar within \\ropt; beyond that radius the \\rd\\ templates have steeper slopes. This difference is mostly attributed to the high brightness galaxies, whose RCs are typically traced further out in the disks, and which are characterized by \\ropt/\\rd\\ ratios larger than the average for the rest of the sample. We do not find convincing evidence for declining RCs {\\em within the spatial scales sampled by our data}, i.e. 4.5--6 \\rd, even for the most massive systems, in contradiction with the URC models (which, beyond \\ropt, are based on linear extrapolations). As argued in \\S \\ref{s_res}, there are a few examples of template RCs whose outermost points show a small velocity decrease, but there is no clear indication of declining outer slopes (as they would be measured from linear fits to the outer regions of the average RCs beyond 2 or 3 \\rd, or, equivalently, 0.6 or 0.9 \\ropt). One template RC (the one labeled $-$22.60 in Figure \\ref{avgrcs_ropt}) does appear to be declining beyond \\about 1.1 \\ropt, but the adjacent curves have rising or flat outer slopes. For less luminous objects or below \\ropt, our templates are qualitatively consistent with the URC tracings, when the larger scatter of PSS96's average RCs and different selection criteria (namely, the exclusion of edge-on galaxies from our analysis) are taken into account. Being significantly larger than the sample upon which the URC is based, our data set allows us to compute average RCs that extend beyond \\ropt\\ without reducing the number of luminosity bins or resorting to extrapolations. Our results are consistent with the declining RCs of \\citet{cvg91}, since the velocity decrease is observed in \\hi\\ data that span twice the distance sampled by our optical data, and with the hybrid \\Ha + \\hi\\ RCs of 8 fast rotators studied by \\citet{spe05}, which show remarkably flat or mildly declining outer slopes beyond 1.5 \\ropt.  Based on our results we cannot make quantitative inferences about the dynamic role of dark matter in galaxies, because optical RCs alone cannot constrain the properties of dark matter halos on large scales. In fact, a detailed mapping of the radial distribution of visible and dark matter in galaxies requires the use of luminosity profiles and extended \\hi\\ RCs in addition to optical RCs. Nonetheless, the comparison between Figure \\ref{avgrcs_ropt} and the URC models suggests that dark matter plays a more significant role in determining the internal kinematics {\\em within the optical disk} of the most luminous systems than what claimed by PSS96. At high luminosity, the profiles of PSS96's RCs show a strong velocity decrease between 1 and 2 \\ropt; their decomposition into a disk and a dark halo component in terms of the URC parameterization implies a modest dark matter content, whereas low luminosity systems are dark matter dominated. Since the templates RCs of the high luminosity systems do not appear to decline over those spatial scales, their dark matter content within the optical disk based on the URC predictions might be underestimated.  The template RCs, expressed by the Polyex parameters as a function of galaxy luminosity, can be used to constrain models of the circular velocity field of disks, with applications that range from studies of dark matter content and kinematic properties of galaxies, to numerical simulations of disk formation and evolution within the current cosmological framework, and in general studies that, for different reasons, must rely on RC models. Our parameterization in terms of disk scale lengths should prove especially useful in the theoretical and numerical modeling of disk structure and evolution.  The results presented here will be used to determine statistical corrections for aperture bias of emission line widths obtained with fiber spectroscopy, such as those that the on-going Sloan Digital Sky Survey \\citep{str02} is collecting for one million galaxies. Full rotational velocities of quiescent disk galaxies with apparent sizes larger than the fiber diameter can be statistically recovered by simulating the impact of the finite aperture on the observations, where the galaxy RCs are modeled using our template relations (B. Catinella et al., in preparation; \\citealt{cat05})."
    },
    "0512/astro-ph0512156.txt": {
        "abstract": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We present a detailed analysis of the velocity distribution and orientation of orbits of subhaloes in high resolution cosmological simulations of dark matter haloes. We find a trend for substructure to preferentially revolve in the same direction as the sense of rotation of the host halo: there is an excess of prograde satellite haloes. Throughout our suite of nine host haloes (eight cluster sized objects and one galactic halo) there are on average 59\\% of the satellites corotating with the host. Even when including subhaloes out to five virial radii of the host, the signal still remains pointing out the relation of the signal to the infall pattern of subhaloes. However, the fraction of prograde satellites weakens to about 53\\% when observing the data along a (random) line-of-sight and deriving the distributions in a way an observer would infer them. This decrease in the observed prograde fraction has its origin in the technique used by the  observer to determine the sense of rotation, which results in a possible misclassification of non-circular orbits. We conclude that the existence of subhaloes on corotating orbits is another prediction of the cold dark matter structure formation scenario, although there will be difficulties to verify it observationally. Since the galactic halo simulation gave the same result as the cluster-sized simulations, we assume that the fraction of prograde orbits is independent of the scale of the system, though more galactic simulations would be necessary to confirm this. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\label{sec:Introduction} There is mounting evidence that the Cold Dark Matter (CDM) structure formation scenario provides the most accurate description of our Universe. Observations point towards a ``standard'' \\LCDM\\ universe comprised of 28\\% dark matter, 68\\% dark energy, and luminous baryonic matter (i.e. galaxies, stars, gas, and dust) at a mere 4\\% \\citep[cf. ][]{Spergel03}. This so-called ``concordance model'' induces hierarchical structure formation whereby small objects form first and subsequently merge to form progressively larger objects \\citep{WhiteRees78, Davis85, Tormen97}. Whereas the large scale structure of our present Universe can be reconstructed very well by numerical simulations, the small scale structure still poses some problems. For example, there are many more subhaloes predicted by cosmological simulations than actually observed in nearby galaxies \\citep[cf. ][]{Klypin99,Moore99}. The lack of observational evidence for these satellites has led to the suggestion that they are completely (or almost completely) dark, with strongly suppressed star formation due to the removal of gas from the small protogalaxies by the ionizing radiation from the first stars and quasars \\citep{Bullock00, Tully02, Somerville02}. Others suggest that perhaps low-mass satellites never formed in the predicted numbers in the first place, indicating problems with the \\LCDM\\ model in general, replacing it with Warm Dark Matter instead \\citep{Colin00, Bode01, Knebe02}.  Recent results from (strong) lensing statistics suggest that the predicted excess of substructure is in fact required to reconcile some observations with theory \\citep{DalalKochanek02,Dahle03}, although this conclusion has not been universally accepted \\citep{SchechterWambsganss02, EvansWitt03, Sand04}. If, however, the lensing detection of halo substructure \\textit{is} correct and the overabundant satellite population really does exist, it is vital to understand the orbital evolution of these objects and their deviation from the background dark matter distribution. In order to test the underlying \\LCDM\\ model, more predictions are necessary which can be confirmed or disproved by observations.   Here, we are investigating cosmological simulations based on the standard $\\Lambda$-Cold Dark Matter model, concerning the sense of rotation of subhaloes. Host dark matter haloes usually carry a small internal angular momentum, which is established by the transfer of angular momentum from infalling matter via tidal torques \\citep{Peebles69, BarnesEfstathiou87}. However, \\citet{Gardner01} as well as \\citet{Vitvitska02} proposed another explanation for the origin of the angular momentum in galaxies and their dark matter haloes. They claim that haloes obtain their spin through the cumulative acquisition of angular momentum from satellite accretion. These two descriptions are certainly linked together and mutually dependent, respectively. A detailed analysis of the orbits of satellite haloes shows that they are directly connected to the infall pattern of satellites along the surrounding filaments \\citep[e.g.][]{Tormen97,Knebe04,Zentner05}. Those subhaloes falling into the host at early times establish the angular momentum of the inner regions of the primary halo \\citep[cf.][]{Vitvitska02} and are channelled into the host along the same direction as those merging at later times. This leads to the speculation that satellites are preferentially corotating with the host, which is the major assumption (and will be verified in) the current study.  A similar study was recently presented by \\citet{Azzaro05}. They performed corresponding investigations using both a cosmological simulation and observational data from the Sloan Digital Sky Survey. These observational data showed a signal of 61\\% corotating satellite galaxies, quite independent of magnitudes or distance to the host galaxy. When projecting their simulational data in order to ``observe'' the simulations, they found a fraction of 55--60\\% prograde satellite galaxies.  In another related study, \\cite{Shaw05} studied a sample of 2200 (low resolution) dark matter haloes and determined the sense of rotation of all substructure \\emph{particles} with respect to the host halo. They found a very strong signal for these particles to be corotating with the host. However, we are using here a different method for determining the fraction of prograde satellites, classifying each satellite individually and then counting the number of pro- and retrograde satellites rather than (dark matter) particles.  We present here the analysis of the sense of rotation for subhaloes in nine cosmological simulations: eight cluster-sized objects with varying merger histories and one galactic dark matter halo. These simulations are described in more detail in \\Sec{sec:simulations}, while \\Sec{sec:simus} deals with the results when analyzing the full six dimensional phase-space information at hand. There we show how the sense of rotation for satellites can be determined and discuss various influences on the fraction of prograde orbits. In \\Sec{sec:observer} we investigate how our results can possibly be validated observationally and conclude with a discussion and summary of our findings in \\Sec{sec:discussion} and \\ref{sec:conclusions}.   % TABLE % % data thrown together by hand \\begin{table*} \\begin{center} \\begin{tabular}{|l|l|c|c|c|c|c|c|c|c|c|c|c|c} \\hline host & $M_{\\rm vir}$ & $R_{\\rm vir}$ & $\\lambda$ & $V_{\\rm max}$ & $R_{\\rm max}$ & $\\sigma_{v, \\rm host}$ %& $e_1$ %& $e_2$ & $T$ & age & $z_{\\rm form}$\t%added in revision %& 1/age & $\\left< \\frac{\\Delta M}{\\Delta t M}\\right>$ & $\\sigma_{\\Delta M/M}$ \\\\ \\hline\\hline  C1 &   2.9 & 1355 & 0.0157 & 1141 &  346 & 1161 &   0.365 &   7.9  &   1.052& 0.128 &  0.125 \\\\ C2 &   1.4 & 1067 & 0.0091 &  909 &  338 &  933 &   0.388 &   6.9  &   0.805& 0.122 &  0.156 \\\\ C3 &   1.1 &  973 & 0.0125 &  828 &  236 &  831 &   0.265 &   6.9  &   0.805& 0.100 &  0.117 \\\\ C4 &   1.4 & 1061 & 0.0402 &  922 &  165 &  916 &   0.639 &   6.6  &   0.750& 0.127 &  0.207 \\\\ C5 &   1.2 & 1008 & 0.0093 &  841 &  187 &  848 &   0.909 &   6.0  &   0.643& 0.129 &  0.141 \\\\ C6 &   1.4 & 1065 & 0.0359 &  870 &  216 &  886 &   0.073 &   5.5  &   0.567& 0.147 &  0.153 \\\\ C7 &   2.9 & 1347 & 0.0317 & 1089 &  508 & 1182 &   0.531 &   4.6  &   0.443& 0.844 &  1.068 \\\\ C8 &   3.1 & 1379 & 0.0231 & 1053 &  859 & 1091 &   0.587 &   2.8  &   0.237& 0.250 &  0.225 \\\\[1ex] G1 &   1.2$\\times 10^{-2}$  &  214 & 0.0229 &  210 &   44 &  202 &  0.491 & 8.5 &   1.232&   0.073 &  0.107 \\\\ \\hline  % data from plot_orientsatcum.pro, which gives the data file satmass.dat  \\end{tabular} \\caption{Properties of the host haloes in our simulations. Masses are measured in 10$^{14}$\\hMsun, velocities in km/s, distances in \\hkpc, and the age in Gyr. ($M_{\\rm vir}$ = virial mass, $R_{\\rm vir}$ = virial radius, $\\lambda$ = spin parameter, $V_{\\rm max}$ = maximum of the rotation curve, $R_{\\rm max}$ = position of the maximum, $\\sigma_{v, \\rm host}$ = velocity dispersion of the host, $T$ = triaxiality parameter, age = time since half of the present day mass has formed, $z_{\\rm form}$ = corresponding formation redshift$, \\left< \\frac{\\Delta M}{\\Delta t M}\\right>$ = mean rate of relative mass change, $\\sigma_{\\Delta M/M}$ = dispersion of the rate of relative mass change) }  \\label{t:haloprop} \\end{center} \\end{table*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %Summarize everything and repeat the main ideas and results, put it all in a greater context. Maybe also a discussion?  We investigated the sense of orientation of satellite orbits in cosmological dark matter haloes at redshift $z=0$. Our set of simulations consisted of eight cluster-sized haloes and one simulation of a Milky Way-type galactic dark matter halo.  From the theory of the generation of angular momentum in dark matter haloes, we expected to get a trend towards satellites corotating with the host.  We found a mean majority of 59\\% prograde orbits averaged over all nine simulations. Three of our simulations expressed a rather isotropic distribution of orientations of orbits with respect to the host. Two of them could be assigned a rather small spin parameter of the host, while the third could be confirmed as a recent (triple) merger leading to random orientations of satellite orbits.  We validated our results by carrying out a Kolmogorov-Smirnov test for the cumulative distribution of angles between satellite orbit angular momentum and host halo spin. All simulations, except the already mentioned exceptions, are clearly not in agreement with an isotropic distribution. There is also no sign for a symmetry between prograde and retrograde orbits in these simulations, stressing that the excess of corotating satellites is a reliable prediction of the hierarchical \\LCDM\\ structure formation scenario.  We chose to use the inner host angular momentum, i.e. within 30\\% of the virial mass $M_{\\rm vir}$ (corresponding to $\\approx 0.2\\,R_{\\rm vir}$) for defining the rotation of the host halo. However, other choices (at 50\\%, 100\\% of $M_{\\rm vir}$) lead to the same mean prograde fraction of 59\\%. Thus, uncertainties in the angular momentum only have a minor effect.  We further checked the influence of the environment outside the virial radii $R_{\\rm vir}$ of the hosts by including satellites up to $5\\,R_{\\rm vir}$. Even though the prograde fraction decreased and the orientations became more isotropic, the signal yet remained -- in agreement with the picture that the environment of matter and satellites determines the sense of host halo rotation.  Thus, we conclude that there is a trend towards an overbalance of corotating subhaloes in $\\Lambda$CDM simulations at redshift $z=0$, regardless of the inner host boundary defining the host angular momentum and the outer boundary determining the sphere within which satellites are included. We could not confirm any relation between triaxiality and prograde fraction and, even more important, we find our signal in the cluster-sized simulations as well as in the galaxy-sized simulation.  The question now remained, how feasible it is to observationally test the prediction of having roughly 60\\% satellites on orbits corotating with their host halo. To obtain an answer, we chose 100 random lines of sight projecting our three-dimensional data into an observer's plane. The signal still remained, yet weakened to $53\\% \\pm 3\\%$, which is difficult to be detected.  The origin of this weakening lies in the observer's method for distinguishing between pro- and retrograde orbiting satellites by comparing their motion with the part of the host where they reside (cf. \\Sec{class-pro-retro}). This commonly used method works perfectly for circular orbits, yet satellite orbits are not necessarily circular but rather exhibit a distribution of circularities \\citep[cf.][]{Stuart2}. We showed that the recovery of the original signal can be improved by weighing the satellites with their relative line-of-sight velocities, yielding an observed prograde fraction of 54\\%.  Additionally, we studied the effect of interlopers, satellites which pass by the observed cluster or galaxy along the line of sight, and found it to be negligible (i.e. less than 1\\%).  Comparison with observational work yields different results: the study of \\citet{Zaritsky97} agrees with our findings, \\citet{Azzaro05}'s observational results give a much higher prograde fraction, whereas \\citet{Carignan-ea97} found for NGC 5084 a strong excess of retrograde satellite galaxies. More galaxy-sized simulations, maybe even including baryonic matter, would be necessary in order to study the prograde fractions on smaller scales in more detail, which would then enable a better comparison with observational results.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %  Acknowledgements, Bibliography and the Appendix-stuff  %  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0512/astro-ph0512267_arXiv.txt": {
        "abstract": "\\noindent Temperature maps of the Cosmic Microwave Background (CMB) radiation, as those obtained by the Wilkinson Microwave Anisotropy Probe (WMAP), provide one of the most precise data sets to test fundamental hypotheses of modern cosmology. One of these issues is related to the statistical properties of the CMB temperature f\\/luctuations, which would have been produced by Gaussian random density f\\/luctuations when matter and radiation were in thermal equilibrium in the early Universe. We analysed here the WMAP data and found that the distribution of the CMB temperature f\\/luctuations $P^{\\text{CMB}}(\\Delta T)$ can be quite well fitted by the anomalous temperature distribution emerging within nonextensive statistical mechanics. This theory is based on the nonextensive entropy $S_q \\equiv k \\{ 1-\\int dx \\, [P_q(x)]^q \\} /(q-1)$, with the Boltzmann-Gibbs expression as the limit case $q \\to 1$. For the frequencies investigated ($\\nu=$ 40.7, 60.8, and 93.5 GHz), we found that $P^{\\text{CMB}}(\\Delta T)$ is well described by $P_q(\\Delta T) \\propto 1/[1+(q-1) B(\\nu) \\Delta T\\:\\!^2]^{1/(q-1)}$, with $q = 1.055 \\pm 0.002$, which exclude, at the 99\\% confidence level, exact Gaussian temperature distributions $P^{\\text{Gauss}}(\\Delta T) \\propto e^{- B(\\nu) \\Delta T\\:\\!^2}$, corresponding to the $q \\to 1$ limit, to properly represent the CMB temperature f\\/luctuations measured by WMAP. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512321_arXiv.txt": {
        "abstract": "We present the latest results of the Clusters in the Zone of Avoidance (CIZA) survey, which is mapping the large-scale matter distribution behind the Milky Way by performing the first systematic search for X-ray luminous galaxy clusters at low Galactic latitudes.  The survey's approach, which uses X-ray emission to locate cluster candidates, overcomes the problems faced by optically selected cluster surveys which have traditionally avoided this region of the sky due to the severe extinction present along the Galactic plane.  We here present the second flux-limited CIZA cluster catalog containing 60 X-ray luminous galaxy clusters, 88\\% of which are new discoveries.  We also examine the degree to which known superclusters extend into the Zone of Avoidance and highlight newly discovered structures which have previously gone unnoticed.  We show that the survey has found far fewer rich clusters in the Great Attractor region than would be expected given the region's proposed mass.  Instead, we find a significant increase in the number of clusters \\emph{behind} the Great Attractor, with the most notable being an association of clusters near the Shapley supercluster.  We propose these clusters trace an extension of the large-scale filament network in which the Shapley concentration is embedded.  We also highlight an association of clusters near the Galactic anticenter, which is the first supercluster found to be completely hidden by the Milky Way.    Our finding of a less massive Great Attractor and the detection of significant structures behind the complex supports studies which suggest the motion of nearby galaxies, including that of the Local Group, is due, in part, to a large-scale bulk flow which is induced by overdensities beyond the Great Attractor region. ",
        "introduction": "Galaxy clusters are the nodes of the cosmic web which permeates the Universe, making them probes of the organization of matter on the grandest scales.  While galaxies have been the preferred tracer of the local distribution of matter (Geller \\& Huchra 1989; Saunders et al. 1995; Colless et al. 2001; York et al. 2000), clusters offer a complementary means to trace the large-scale structure of the Universe to distances beyond the effective range of galaxy redshift surveys (Tully et al. 1992; Mullis et al. 2001).  Although clusters more sparsely sample the cosmic web, the similarity in the shape of the galaxy and cluster power spectrums suggest both are effective in tracing the same underlying density field (Peacock \\& Dodds 1994; Tadros et al. 1998).  The advantage offered by clusters is that their significant X-ray luminosities and galaxy overdensities allow for the compilation of volume-limited samples out to extremely large distances ($r > 200 h^{-1}$ Mpc).  In addition to the study of large-scale cosmography, such samples provide a means to investigate a range of cosmological issues.  The location of clusters provides insight into the spatial distribution of the initial fluctuations which are thought to have seeded the current cluster population in early epochs, the mass function of clusters offers details on the amplitude spectrum of those primordial fluctuations, and the evolution of the number density of clusters with redshift can constrain fundamental cosmological parameters.  More locally, the deep gravitational wells of clusters act as accelerators of large-scale flows and knowledge of their distribution provides a means to reconstruct the peculiar velocity field (Branchini \\& Plionis 1996; Kocevski et al. 2004, 2005).     The largest cluster sample used for many such studies has been the optically selected Abell/ACO catalog (Abell 1958; Abell, Corwin \\& Olowin 1989, hereafter ACO).  The Abell/ACO sample has been widely employed due to its significant sky coverage (2/3), large characteristic depth ($\\sim240 h^{-1}$ Mpc) and relatively robust cluster identification criteria compared to other visually compiled samples such as the Zwicky catalog (Zwicky et al. 1961-68).  More recently, moderately deep X-ray cluster surveys over large parts of the sky have been carried out using data from the \\emph{ROSAT} All-Sky Survey (RASS, Tr\\\"{u}mper 1993), which is the first and only survey to provide X-ray imaging data over the entire sky.  While an array of cluster samples have been compiled from the RASS (Romer et al. 1994; Pierre et al. 1994; Burns et al. 1996; Ebeling et al. 1996, 1998, 2000, 2001, 2002; Scharf et al. 1997; De Grandi 1999; Henry et al. 1997; Vikhlinin et al. 1998; Ledlow et al. 1999; B\\\"{o}hringer et al. 2000; Mullis et al. 2003, 2004), the largest and most complete are currently the extended Brightest Cluster Sample (eBCS, Ebeling et al. 1998, 2000) and the \\emph{ROSAT}-ESO Flux Limited X-ray catalog (REFLEX, B\\\"{o}hringer et al. 2004). Covering the northern and southern hemispheres, respectively, the eBCS and REFLEX samples contain a combined 737 clusters with X-ray fluxes above $3\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ and redshifts of $z\\leq 0.3$.    \\subsection{The Zone of Avoidance}  Since finding overdensities in the galaxy distribution becomes increasingly difficult at low Galactic latitudes due to the severe extinction and stellar obscuration present in the direction of the Milky Way (MW), optically selected catalogs such as Abell/ACO are plagued with poor coverage in a $40^{\\circ}$ wide strip centered on the plane of the Galaxy known as the Zone of Avoidance (ZOA).  Unfortunately, despite the decreased extinction through the MW at X-ray wavelengths, both the eBCS and REFLEX surveys have followed the precedent set by optically selected surveys and have also avoided the Galactic plane.  With the eBCS and REFLEX surveys limited to Galactic latitudes of $|b|>20^{\\circ}$, the Galactic plane, which covers 1/3 of the sky, is the final region in which the cluster distribution remains to be mapped.  This is particularly troubling since there is considerable dynamical evidence that the ZOA harbors some of the most massive large-scale structures in the local Universe.  Both the Local Group's (LG) peculiar velocity (Kogut et al. 1993) and the motion of galaxies within $40 h^{-1}$ Mpc of the MW (Lynden-Bell et al. 1988, hereafter LB88) suggest galaxies in the local volume are flowing toward a vertex behind the Galactic plane, presumably due to the gravitational influence of overdensities in this region.  Assuming this motion was due to infall into a single ``Great Attractor'' (GA), LB88 estimated that the source of the flow was located roughly $43 h^{-1}$ Mpc away and near the Hydra-Centaurus cluster association.  This distance, coupled with the large peculiar velocities observed in nearby galaxies, implied the rather high mass of $\\sim5\\times10^{16} h^{-1}_{50}$ M$_{\\odot}$ for the GA complex.  LB88 found that such a relatively nearby, massive overdensity would be responsible for $\\sim70\\%$ of the LG's observed peculiar velocity.       The LB88 findings have been controversial since subsequent redshift surveys that have encompassed the GA region have failed to detect a mass overdensity as large as the one implied by the LB88 peculiar velocity data (Dressler 1988; Strauss et al. 1992; Hudson 1993, 1994), nor have they conclusively measured the backside infall into the GA one would expect if the region were best described as a single, stationary attractor (Mathewson et al. 1992, Courteau et al. 1993).  More recent work has suggested that a significant component of the LG's peculiar velocity is in the form of a large-scale bulk flow which continues past the GA region and is induced by even more distant structures in the ZOA (Zaroubi et al. 1999; Tonry et al. 2000; e.g. Hudson et al. 2003).  The difficulty in determining the source of the LG's motion and the possible bulk flow in which it participates is largely due to the incompleteness present in the direction of the MW.  With cluster surveys avoiding the Galactic plane, the distribution of large-scale structures which are the accelerators of such flows in this region has remained largely unknown.   A variety of techniques have been used to reconstruct the ZOA, ranging from uniform filling (Strauss \\& Davis 1988; Lahav 1987) to a spherical-harmonics approach which extends structures above and below the plane into the ZOA (Plionis \\& Valdarnini 1991, cf. Brunozzi et al. 1995), but the value of these reconstruction techniques is limited if the MW does indeed obscure dynamically significant structures, as has been suggested.  \\begin{figure*}[t] \\epsscale{1.1} \\plotone{figure1.nhmap2.eps} \\caption{Distribution of the current CIZA sample along the Galactic plane overlayed on contours of Galactic neutral Hydrogen column density in units of $10^{20}$ cm$^{-2}$.  Open circles denote the location of B1 clusters, while filled circles represent B2 clusters.  The current CIZA sample contains 133 clusters with X-ray detect fluxes greater than $3\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$.} % \\end{figure*}  In order to reveal the nature of regions such as the GA and other large-scale structures hidden behind the MW, we began the Clusters in the Zone of Avoidance (CIZA) project to map the cluster distribution behind the Galactic plane for the first time.  Our approach, which uses X-ray emission to locate cluster candidates, overcomes the problems faced by optically selected catalogs at low Galactic latitudes due to the decreased extinction at X-ray wavelengths.  The strategy of the CIZA survey is consistent with that of the eBCS and REFLEX surveys and therefore the completion of CIZA project will facilitate the construction of the first all-sky, X-ray selected cluster sample.  In this paper we present the latest results from the CIZA survey.  This is the second in a series of papers regarding the CIZA project; Ebeling, Mullis \\& Tully (2002, hereafter Paper I) recently introduced the survey and presented a subsample of the 73 brightest CIZA clusters (the B1 sample).  We have since extended the survey to lower fluxes and present here the next CIZA subsample (the B2 sample) as well as a progress report on the survey, which is nearing completion.  In what follows, we describe the CIZA survey in more detail in \\S2, present the B2 sample in \\S3, and discuss the large-scale cosmography of the cluster distribution behind the Galactic plane in \\S4.  Finally the future of the CIZA survey and our concluding remarks are presented in \\S5.  Throughout this paper we use $H_{0}=50$ $h$ km s$^{-1}$Mpc$^{-1}$ for the calculation of comoving distances and all listed X-ray fluxes are given in the 0.1--2.4 keV \\emph{ROSAT} band unless otherwise noted. ",
        "conclusions": "Although a handful of galaxy clusters have previously been found in the ZOA serendipitously, the CIZA survey is carrying out the first systematic search for clusters at low Galactic latitudes.  Thus far the survey has spectroscopically confirmed 205 galaxy clusters behind the Galactic plane, 60 of which are presented in this paper.  With the CIZA project proving effective in finding X-ray luminous clusters in the ZOA, the survey is providing the means to study the LSS behind the MW for the first time.  The current CIZA sample has already begun to shed light on previously hidden large-scale cluster associations and we may soon be able to settle the debate regarding the nature of the GA region.  Our preliminary findings suggest the GA contains far fewer rich cluster systems than would be expected given the regions proposed mass.  Having surveyed the entire GA complex, we find an order of magnitude discrepancy remains between the mass observed in the supercluster and that implied by the peculiar motion of nearby galaxies.  This inconsistency may be resolved by our detection a significant overdensity of clusters which appear near to the GA in projection, but are in fact four times as distant.  These include clusters between the GA and the SSC superclusters, as well as a large-scale association of clusters which may trace an extension of the SSC into the ZOA.  These findings support recent studies which suggest the peculiar velocity of nearby galaxies is only partly a result of infall into the GA, with the remaining motion due to a large-scale bulk flow induced by more distant overdensities in the ZOA.  To investigate these issues further we are currently extending the CIZA survey to fainter fluxes.  Our motivation is two-fold: (1) since the BSC detect fluxes are known to underestimate the total X-ray emission of low redshift clusters, it is possible that clusters in the nearby GA region have dropped below the flux limit of the B2 subsample, and have thus far been systematically overlooked, and (2) due to the bias in the BSC detect fluxes, the CIZA \\emph{detect} flux limit is not equivalent to that of the eBCS and REFLEX samples covering the extragalactic sky since they are both flux limited down to a \\emph{total} X-ray flux of $3\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$.  To address both of these issues we are currently completing the CIZA survey down to a BSC detect flux of $2\\times10^{-12}$ erg cm$^{-2}$ s$^{-1}$. In addition to containing intrinsically fainter clusters, we expect this B3 subsample to hold nearby, extended clusters whose total X-ray emission has been underestimated.  The B3 subsample currently contains 57 X-ray luminous clusters and we expect to complete the follow-up observation of our final cluster candidates shortly.  With the completion of this sample and ultimately the CIZA survey, we will be able to conclusively determine the extent of the GA region and combine the CIZA sample with the eBCS and REFLEX catalog to produce the first X-ray selected, all-sky cluster sample."
    },
    "0512/astro-ph0512447_arXiv.txt": {
        "abstract": "We present the spectral and temporal radiative signatures expected within the ``Supercritical Pile\" model of Gamma Ray Bursts (GRB). This model is motivated by the need for a process that provides the dissipation necessary in GRB and presents a well defined scheme for converting the energy stored in the relativistic protons of the Relativistic Blast Waves (RBW) associated with GRB into radiation; at the same time it leads to spectra which exhibit a peak in the burst $\\nu F_{\\nu}$ distribution at an energy $E_p \\simeq 1$ MeV in the observer's frame, in agreement with observation and largely {\\sl independent} of the Lorentz factor $\\Gamma$ of the associated relativistic outflow. Futhermore, this scheme does not require (but does not preclude) acceleration of particles at the shock other than that provided by the isotropization of the flow bulk kinetic energy on the RBW frame. In  the present paper we model in detail the evolution of protons, electrons and photons from a RBW % to produce detailed spectra of the prompt GRB phase  as a function of time from across a very broad range spanning roughly $4 \\, {\\rm log}_{10} \\Gamma$ decades in frequency. The model spectra are in general agreement with observations and provide a means for the delineating of the model parameters through direct comparison with trends observed in GRB properties. ",
        "introduction": "The discovery of GRB afterglows by BeppoSAX (Costa et al. 1997) and the ensuing determination of their redshifts (van Paradijs et al. 1997) by and large settled the issue of their distance and luminosity. This discovery, then, settled also the issue of their energetics in favor of emission by Relativistic Blast Waves (RBW) moving toward the observer with Lorentz factors $\\Gamma \\sim 100 - 1000$, as it had been discussed earlier by Rees \\& M\\'esz\\'aros (1992) and M\\'esz\\'aros \\& Rees (1992). The relativistic boosting of the radiation emitted at the rest frame of the RBW (by $\\simeq \\Gamma^4$ in luminosity) then resolved also the issue of their huge energy budget requirements (if at cosmological distances) and brought them to agreement with the energy release in stellar gravitational collapse. At the same time it provided consistency between the huge implied GRB luminosities and their apparently thin spectra (Krolik \\& Pier 1991; Fenimore, Epstein \\& Ho 1993; Barring \\& Harding 1995). The same considerations provided also (Rees \\& M\\'esz\\'aros 1992) an order of magnitude relation between the GRB duration $\\Delta t$ and the size of the radiating region, namely $R \\simeq 10^{16} (\\Delta t/30 \\; {\\rm sec})\\, (\\Gamma / 100)^2$ cm.  These estimates of the kinematic state of the GRB emitting plasma have been supplemented by certain dynamical considerations. For example, following the work of Shemi \\& Piran (1990) it has been generally accepted that a certain amount of baryons must be carried off with the blast waves responsible for the GRBs. This baryon contamination has even been deemed necessary for the efficient transport of the GRB energy away from the environs of its ``inner engine\", else the entire blast wave's internal energy would be converted into radiation on very short time scales, leading to events of very different temporal and spectral appearance (e.g. Paczy\\'nski 1986) than the observed GRB. In fact, the original models of Rees \\& M\\'esz\\'aros (1992) and M\\'esz\\'aros \\& Rees (1992) relied on the presence of a rather precise amount of baryons within the GRB fireball: enough to keep the fireball optically thick and thus allow the conversion of its internal energy to directed motion upon its expansion, but not too many as to render it only mildly (or even non-) relativistic. Even in the more recent (and perhaphs more plausible) variants of the same model that use Poynting  flux (rather than photon energy density) as the agent responsible for the RBW acceleration (see Vlahakis \\& K\\\"onigl 2002, 2004), the circumstellar matter swept up by the RBW to the radius of $R \\sim 10^{16}$ cm, contains roughly as much energy stored in protons as in magnetic field. In either case, the models are called to shed light to these two generic, still open, issues of the GRB dynamics: (a.) The acceleration of the associated Relativistic Blast Waves to Lorentz factors $\\Gamma \\sim 10^2-10^3$ (b.) The dissipation of the energy stored in protons and/or magnetic fields at the rates necessary to produce the prompt GRB emission with the proper attributes. Both these issues are still open in GRB physics.  An altogether different issue associated with the prompt GRB emission is that of their spectra. The differential photon GRB spectra can be fit very well by the so-called Band-spectrum (Band et al. 1993) that consists (asymptotically) of two power laws of indices $\\alpha$ and $\\beta$ joining smoothly at a break energy $E_b$, i.e. % \\begin{equation} N(E) \\propto \\left\\{ \\begin{array}{cc} E^{\\alpha} \\; e^{-(\\alpha - \\beta)E/E_b} & \\mbox {for $E <  \\, E_b$}\\\\ E^{\\beta} \\; E_b^{\\alpha - \\beta} \\; e^{-(\\alpha - \\beta)} & \\mbox {for $E >  E_b$} \\end{array} \\right. \\end{equation} It was pointed out by Malozzi et al. (1995), and confirmed by a larger sample of {\\sl BATSE} data (Preece et al. 2000), that the values of $E_b$ are narrowly distributed around $E_b = 200$ keV with a similarly narrow distribution for $\\alpha$ around the value $\\alpha =-1$, while the distribution of $\\beta$  has a maximum near $\\beta \\simeq -2.3$ and extends to values $\\beta \\lsim -4$, with only few bursts having $\\beta > -2.3$. These values imply that the GRB peak energy of their $\\nu F_{\\nu}$ spectra, (i.e. the peak energy of their emitted luminosity $E_{\\rm p} = (\\alpha + 2) \\, E_b /(\\alpha - \\beta)$) is equally narrowly distributed about the same energy which, when corrected for the GRB redshift ($z_{_{\\rm GRB}} \\sim 1-2$), shifts close to $E_{\\rm p}\\simeq 0.5$ MeV.  Both the presence of the narrowly distributed energy $E_{\\rm p}$ and the value of the low energy index $\\alpha \\simeq -1$ are hard to understand within the conventional wisdom models which suggest that the observed prompt GRB $\\gamma-$ray emission is due to synchrotron radiation by relativistic electrons. Under these assumptions, $E_p$ should be proportional to $\\Gamma^4$ (two powers of $\\Gamma$ come from the synchrotron emission, one from the magnetic field - assuming equipartition with the postshock matter, and one more from the transformation of this energy to the lab frame). This strong dependence on the value of $\\Gamma$ would imply a rather unique value for this latter parameter, else the $E_{\\rm p}$ range would be much broader than observed, in disagreement with observation. While there may indeed be a reason for a very narrow range in the values of $\\Gamma$ consistent with the observed range of $E_{\\rm p}$, as of today such a reason is unknown (at least to the authors). If the same emission is due to electrons accelerated to energies beyond those implied by the shock Rankine - Hugoniot conditions (i.e. $\\gamma_e \\simeq \\Gamma$), or to a low energy cut-off $E_c$ in the electron injection spectrum (with $E_c \\gg \\Gamma \\, m_ec^2$), the presence of a rather well defined peak in the GRB $\\nu F_{\\nu}$ spectra is even harder to understand, as in either case there is no apparent reason for such a well defined energy scale.  Considerations along similar lines do argue for very specific values for the low energy power law index $\\alpha$: If the observed peak in the $\\nu F_{\\nu}$ spectra is due to synchrotron emission by a $\\delta-$function--like electron distribution, $\\alpha$ should be that corresponding to the single electron synchrotron emissivity, $\\alpha = - 2/3$  (Katz 1994a), while if the observed peak is due to a low energy cuf-off in the electron injection spectrum, the associated extremely fast cooling should lead to $\\alpha = -3/2$ (Ghisellini 2002). The observed values extend outside the above range, while the most likely values are inconsistent with either of these assumptions in the context of synchrotron radiation for the prompt GRB emission.   The discovery of GRB afterglows,  following the prediction by Katz (1994b) shifted attention from the prompt GRB phase to that of the afterglow. Novel issues related to the physics of GRB ouflows have since emerged, e.g. the narrow range of the total GRB energy budget (Frail et al. 2001; Bloom et al. 2003) and the correlation of GRB luminosities with their spectral lags (Salmonson \\& Galama 2003). In general, the afterglows extended the frequency range of GRB study to the X-rays, optical, IR and radio thus greately expanding our means of probing of the RBW of GRB (see Zhang \\& M\\'esz\\'aros 2004; Piran 2004 for reviews). In the midst of this new flurry of activity, the issues of dissipation of the GRB proton kinetic energy and the narrow $E_{\\rm p}$ distribution, while quietly ignored, have remained largely unanswered. One of the few attempts to address these issues was that of Kazanas, Georganopoulos \\& Mastichiadis (2002; hereafter KGM): In search of a well defined and tractable dissipation mechansim, they proposed a process that utilises proton - photon collisions to convert the energy stored in protons behind the forward shock of the RBW to electrons (and then into radiation). KGM provided the kinematic and dynamic thresholds for this process to take place and showed that, in addition, it produced spectra that peaked at several well defined energies, namely at $ \\sim \\;2 m_ec^2/\\Gamma^2, \\; 2 m_ec^2, \\; m_ec^2 \\Gamma^2$, {\\sl in the observer's frame}. The interesting feature of this model is  that the combination of the threshold for pair production energy and the final Lorentz boost to the observer's frame leads to a spectral peak at $\\sim m_ec^2$ {\\sl largely independent} of the Lorentz factor $\\Gamma$ of the RBW, provided that the latter is larger than a threshold value which depends on the value of the magentic field of the RBW. Therefore, the process described by KGM provides a framework for resolving both the proton-to-electron-energy-transfer and the narrow $E_p$--range problems in a single fell swoop. Besides the above properties, this model differs from most in the present literature in several respects: (a) It does not require (but does not preclude either) the presence of accelerated particle populations, other than those produced by the isotropization of the RWB kinetic energy behind the shock. (b) It does not require equipartition between the proton and electron enery densities. (c) It posits that the observed radiation in the X - $\\gamma$ ray band is upscattered (and then blue-shifted) synchrotron radiation rather than simply blue-shifted synchrotron radiation.   The purpose of the present paper is to present detailed calculations based on the model outlined in KGM, which for reasons explained below (and in the original reference), is referred to as the ``Supercritical Pile\" model for GRB. In \\S 2 of this note we provide a qualitative description of the basic notions that underlie the ``Supercritical Pile\" model and how they apply to GRB. In \\S 3 we describe the details of the model and of the numerical method employed for the solution. In \\S 4 we describe the numerical tests used in making certain that our results make sense and the code works as planned in the case that contains no electrons but produces all necessary electrons from the proton radiative instability. In \\S 5 we repeat the calculations of \\S 4 in the more realistic case that includes also the effects of the presence of electrons. We also present light curves and spectra produced within the model for different values of the relevant parameters. Finally in \\S 6 the resuts are summarised and conclusions are drawn. ",
        "conclusions": ""
    },
    "0512/astro-ph0512392_arXiv.txt": {
        "abstract": "{We present \\Swift{} and \\XMM{} observations of the bright gamma-ray burst GRB\\,050326, detected by the \\Swift{} Burst Alert Telescope. The \\Swift{} X-Ray Telescope (XRT) and \\XMM{} discovered and the X-ray afterglow beginning 54~min and 8.5~hr after the burst, respectively. The prompt GRB\\,050326 fluence was $(7.7\\pm0.9) \\times 10^{-6}$~erg~cm$^{-2}$ (20--150~keV), and its spectrum was hard, with a power law photon index $\\Gamma = 1.25 \\pm 0.03$. The X-ray afterglow was quite bright, with a flux of $7 \\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$ (0.3--8~keV), 1~hr after the burst. Its light curve did not show any break nor flares between $\\sim 1$~hr and $\\sim 6$~d after the burst, and decayed with a slope $\\alpha = 1.70\\pm0.05$. The afterglow spectrum is well fitted by a power-law model, suffering absorption both in the Milky Way and in the host galaxy. The rest-frame Hydrogen column density is significant, $N_{{\\rm H},z} \\ga 4 \\times 10^{21}$~cm$^{-2}$, and the redshift of the absorber was constrained to be $z > 1.5$. There was good agreement between the spatial, temporal, and spectral parameters as derived by \\Swift-XRT and \\XMM. By comparing the prompt and afterglow fluxes, we found that an early break probably occurred before the beginning of the XRT observation, similarly to many other cases observed by \\Swift. However, the properties of the GRB\\,050326 afterglow are well described by a spherical fireball expanding in a uniform external medium, so a further { steepening} is expected at later times. The lack of such { a} break allowed us to constrain the jet half-opening angle $\\vartheta_{\\rm j} \\ga 7\\degr$. Using the redshift constraints provided by the X-ray analysis, we also estimated that the beaming-corrected gamma-ray energy was larger than $3 \\times 10^{51}$~erg, at the high end of GRB energies. Despite the brightness in X rays, only deep limits could be placed by \\Swift-UVOT at optical and ultraviolet wavelengths. Thus, this GRB was a ``truly dark'' event, with the optical-to-X-ray spectrum violating the synchrotron limit. The optical and X-ray observations are therefore consistent either with an absorbed event or with a high-redshift one. To obey the Ghirlanda relation, a moderate/large redshift $z \\ga 4.5$ is required. ",
        "introduction": "\\label{sec:intro}  The \\Swift{} satellite \\citep{Gehrels04} is a mission dedicated to the study of gamma-ray bursts (GRBs) and their afterglows. GRBs are detected and localized by the Burst Alert Telescope \\citep[BAT;][]{Barthelmy05}, and followed up at X-ray (0.2--10 keV) and optical/ultraviolet (1700--6000~\\AA) wavelengths by the X-Ray Telescope \\citep[XRT;][]{Burrows_XRT} and the Ultraviolet/Optical Telescope \\citep[UVOT;][]{Roming05}. During the first year of operation, \\Swift{} has observed some 75 GRB afterglows, already doubling the pre-\\Swift{} sample. This rich dataset has allowed to study in detail the X-ray light curves, both at early and late times, leading to the discovery of a complex behaviour \\citep[e.g][]{Taglia05,Nousek05,Chinca05,Cusu_050319}. Coupled with optical data, either from UVOT \\citep[e.g.][]{Blustin05} or ground-based observatories \\citep[e.g.][]{Berger05}, this has opened a new era in the afterglow modeling. \\Swift{} also provided the first detection of truly dark GRBs, that is, events with no optical emission up to very deep limits \\citep{Roming_dark}. The study of high-redshift GRBs has also started, with the discovery of the first burst at $z > 6$ \\citep{Watson05b,Cusu_050904,Haislip05,Price05,Taglia_050904,Kawai05}. Moreover, it was found that \\Swift{} GRBs have a larger average redshift than those discovered by earlier missions \\citep{Jakob05}.  During the performance verification and calibration phase (2004 Nov 20 through 2005 Apr 5), \\Swift{} observed sixteen GRB afterglows. Twelve of them were observed in automatic mode, and, among these, eight could be promptly (within 200~s since the trigger) observed by XRT and UVOT. In the remaining four cases, the beginning of the observation was delayed by approximately 50~min due to the  Earth occultation constraints. This is the case for the bright GRB\\,050326, which was discovered by BAT on 2005 Mar 26 at 9:53:55 UT \\citep{Markwardt05}. Its coordinates were $\\alpha_{\\rm J2000} = 00^{\\rm h}27^{\\rm m} 34^{\\rm s}$, $\\delta_{\\rm J2000} = -71\\degr22\\arcmin34\\arcsec$, with an uncertainty radius of 3\\arcmin{} \\citep[95\\% containment;][]{Cummings05}. This burst was also detected by the \\textit{Wind}-Konus experiment \\citep{Golenetskii05}, leading to the characterization of its broad-band gamma-ray spectrum.  The \\Swift{} narrow field instruments could begin observing only 54~min after the BAT trigger. A bright, uncatalogued X-ray source was detected by XRT inside the BAT error circle, and was proposed to be the X-ray afterglow \\citep{Moretti05}. However, no source was detected by UVOT at this location \\citep{Holland05}. XRT collected data up to 6.15~d after the burst. Subsequently, the decay of the light curve prevented any further detection of the afterglow. This object was also observed for 45.8~ks by \\XMM{} \\citep{EhlePM05,DeLuca_GCN}, starting 8.5~hr after the trigger.  Only limited ground-based follow-up was reported for this burst. This was likely due to its unfavorable location in the sky (very few telescopes can point at such low declination), as well as to the brightness of the Moon (which was 99\\% full at the time of the GRB explosion). No counterpart at wavelengths other than the X rays was reported.  In this work, we present a complete discussion of the \\Swift{} and \\XMM{} observations of GRB\\,050326. In Sect.~\\ref{sec:prompt} we describe the properties of the prompt emission. In Sect.~\\ref{sec:XRT} we describe in detail the XRT observations, the data reduction procedure, and the temporal and spectral analysis; in Sect.~\\ref{sec:XMM} we do the same for the \\XMM{} data. In Sect.~\\ref{sec:XRT_XMM} we compare the results of the two instruments. In Sect.~\\ref{sec:UVOT} we describe the UVOT optical observations. Finally, in Sect.~\\ref{sec:discussion} we present the physical implications of our observations in the framework of the standard GRB afterglow model. Our conclusions are summarized in Sect.~\\ref{sec:conclusion}.  Throughout this paper, all errors will be quoted at 90\\% confidence level for one parameter of interest, unless otherwise specified. The reduced $\\chi^2$ will be denoted as $\\chi^2_\\nu$, and the number of degrees of freedom with the abbreviation ``d.o.f.''. We follow the convention $F_\\nu(\\nu,t) \\propto t^{-\\alpha}\\nu^{-\\beta}$, where $\\alpha$ and $\\beta$ are the temporal decay slope and the spectral index, respectively. As time origin, we will adopt the BAT trigger \\citep{Markwardt05}. The photon index is $\\Gamma = 1 + \\beta$. Last, we adopt the standard ``concordance'' cosmology parameters, namely $\\Omega_{\\rm m} = 0.27$, $\\Omega_\\Lambda = 0.73$, $h_0 = 0.71$ \\citep[e.g.][]{Spergel03}. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a detailed analysis of the GRB\\,050326 prompt and afterglow emission. The combined capabilities of \\Swift{} (which sampled the light curve for a relatively long time span) and \\XMM{} (which ensured a large statistics) allowed to obtain a thorough characterization of the afterglow properties.  The prompt emission was relatively bright (with a 20--150~keV fluence of $\\sim 8 \\times 10^{-6}$~erg~cm$^{-2}$). The spectrum was hard (photon index $\\Gamma = 1.25 \\pm 0.03$), suggesting a peak energy at the high end of the BAT energy range or beyond. Indeed, thanks to the simultaneous detection of this burst by the \\textit{Wind}-Konus experiment \\citep{Golenetskii05}, the prompt spectrum could be fully characterized. The prompt bolometric fluence was $\\mathcal{F} \\sim 2.4 \\times 10^{-5}$~erg~cm$^{-2}$ (1--10\\,000~keV), and the observed peak energy was $E_{\\rm p,obs} = 200 \\pm 30$~keV.  Due to pointing constraints, XRT and UVOT observations could start only 54~min after the GRB. The X-ray afteglow was quite bright, with a flux of $7 \\times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$ (0.3--8~keV) 1~hr after the GRB. However, no optical counterpart could be detected. The X-ray light curve showed a steady decline, with no breaks or flares. The best-fit power-law decay index was $\\alpha = 1.70 \\pm 0.05$. Such regular behaviour is different from that usually observed by \\Swift, but this may be the result of the limited time coverage (observations could be carried out only between 54~min and 4.2~d after the burst). Indeed, extrapolation of the afterglow light curve to the time of the prompt emission overpredicts the burst flux, and may suggest a slower decay before the beginning of the XRT observation.  The analysis of the combined XRT and \\XMM{} data allowed to characterize in detail the afterglow spectrum. A fit with an absorbed power-law model provided a good description to the data, yielding a photon index $\\Gamma = 2.09 \\pm 0.08$ and a column density significantly in excess to the Galactic value. The best-fit model was thus computed adding an extra absorption component, leaving its redshift $z$ free to vary. Although both $N_{{\\rm H},z}$ and $z$ could not be effectively constrained, a firm lower limit $N_{{\\rm H},z} > 4 \\times 10^{21}$~cm$^{-2}$ could be set. Therefore, GRB\\,050326 adds to the growing set of afterglows with large rest-frame column density \\citep{GalamaWijers01,Stratta04,Campana_NH}. The limits measured in the optical and ultraviolet region by UVOT lie well below the extrapolation of the X-ray spectrum. In particular, they violate the synchrotron limit that the optical-to-X-ray spectral index should be larger than 0.5. This implies a large extinction and/or a high redshift.  The X-ray spectral analysis also allowed us to set the lower limit $z > 1.5$ to the redshift of the absorbing component (and therefore of the GRB). The isotropic-equivalent gamma-ray energy was then $E_{\\gamma,\\rm iso} > 1.4 \\times 10^{53}$~erg. The temporal and spectral properties of the afterglow were nicely consistent with a spherical fireball expanding in a uniform medium, with the cooling frequency above the X-ray range. We could therefore set a lower limit to the jet break time $t_{\\rm b} \\ga 4$~d. The jet opening angle could be constrained to be $\\vartheta_{\\rm j} \\ga 7\\degr$, with only a weak dependence on the (unknown) fireball energy. The beaming-corrected gamma-ray energy was $E_{\\gamma,\\rm j} = (3\\mbox{--}8) \\times 10^{51} (t_{\\rm b}/4~{\\rm d})^{3/4}$~erg, independently from the redshift. GRB\\,050326, thus, released a large amount of gamma rays \\citep[only GRB\\,990123 had a larger energy in the sample of][]{Ghirla04}.  To be consistent with the Ghirlanda relation \\citep{Ghirla04}, two redshift ranges are allowed, either at low ($z \\la 0.8$) or high ($z \\ga 4.5$) redshift. However, to simultaneously satisfy also the limits derived from the X-ray spectral analysis, only the high-redshift region is left. We note that the Ghirlanda relation is still based upon a small sample, so that any inference cannot yet be regarded as conclusive. However, the results from the X-ray spectra, the consistency of the GRB\\,050326 properties with the Ghirlanda relation, and the strong dearth of optical/ultraviolet afterglow flux, are overall consistent with a moderate/high redshift ($z \\ga 4$). A search for the host galaxy through deep infrared and optical imaging may conclusively settle this issue."
    },
    "0512/astro-ph0512501_arXiv.txt": {
        "abstract": "{ The {\\em all-sky survey of Faraday rotation}, a Key Science Project of the planned Square Kilometre Array, will accumulate tens of millions of rotation measure measurements toward background radio sources and will provide a unique database for characterizing the overall magnetic geometry of magnetic fields in galaxies and in the intergalactic medium. Deep imaging of the polarized synchrotron emission from a large number of nearby galaxies, combined with Faraday rotation data, will allow us to test {\\em primordial, gas flow, and dynamo models}\\ for field origin and amplification. The SKA will find the first magnetic fields in young galaxies and determine the timescale for building up small-scale turbulent and large-scale coherent fields. The spectrum of dynamo modes, if existing, will be resolved. The present-day coherent field may keep memory of the direction of the seed field which can be used for mapping the structure of the seed field before galaxy formation. ",
        "introduction": "\\label{intro}   Galactic magnetism may have evolved in subsequent stages:\\\\ (1) Field seeding by primordial fields embedded in the protogalaxy, or fields ejected into the protogalaxy by AGN jets, radio lobes, early supernova remnants, or gamma-ray bursts.\\\\ (2) Field amplification by compressing or shearing flows, turbulent flows, magneto-rotational instability, and dynamos.\\\\ (3) Field ordering by the large-scale dynamo.  Models referring to one or more of these stages can be tested by observations. Radio astronomy provides the best tools to measure galactic magnetic fields. The planned Square Kilometre Array (SKA) will allow fundamental advances in studying the origin and evolution of magnetic fields. ",
        "conclusions": ""
    },
    "0512/astro-ph0512411_arXiv.txt": {
        "abstract": "The intense 0.511 MeV gamma-ray line emission from the Galactic Center observed by INTEGRAL requires a large annihilation rate of nonrelativistic positrons.  If these positrons are injected at even mildly relativistic energies, higher-energy gamma rays will also be produced.   We calculate the gamma-ray spectrum due to inflight annihilation and compare to the observed diffuse Galactic gamma-ray data.  Even in a simplified but conservative treatment, we find that the positron injection energies must be $\\lesssim 3$ MeV, which strongly constrains models for Galactic  positron production. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512627_arXiv.txt": {
        "abstract": "{} {We report the detection of hard X-ray emission from the field of the star-forming region \\ngc with the the International Gamma-Ray Astrophysics Laboratory \\int.} {The \\jemx\\ monitor and \\isgri\\ imager aboard \\int\\ and \\chan\\ {\\it ACIS} imager were used to construct 3-80 keV images and spectra of \\ngc.} {The 3-10 keV and 10-35 keV images made with \\jemx\\ show a complex structure of extended emission from \\ngc\\s1. The \\isgri\\ source detected in the energy ranges 20-40 keV, 40-80 keV, and 20-60 keV coincides with the \\ngc ridge. The 20-60 keV flux from the source is (1.8$\\pm$0.37)$\\times$10$^{-11} \\enf$. Spectral analysis of the source revealed a hard power-law component with a photon index about 1. The observed X-ray fluxes are in agreement with extrapolations of X-ray imaging observations of \\ngc by \\chan {\\it ACIS} and \\asca {\\it GIS}.} {The X-ray data are consistent with two very different physical models. A probable scenario is emission from a heavily absorbed, compact and hard \\chan\\ source that is associated with the AGN candidate radio source NGC 6334B. Another possible model is the extended \\chan\\ source of nonthermal emission from \\ngc\\ that can also account for the hard X-ray emission observed by {\\it INTEGRAL}. The origin of the emission in this scenario is due to electron acceleration in energetic outflows from massive early type stars. The possibility of emission from a young supernova remnant, as suggested by earlier infrared observations of \\ngc\\s1, is constrained by the non-detection of $^{44}$Ti lines.} ",
        "introduction": "\\ngc is a star forming (SF) complex of a total bolometric luminosity $\\sim$1.9$\\times$10$^6 \\Lsun$ associated with a giant molecular cloud of mass $\\sim$1.6$\\times$10$^5 \\Msun$ (e.g. Loughran \\etal 1986). The complex structure of \\ngc with a number of clearly separated and localized SF sites was established by radio and infrared (IR) observations. Massive SF sites reside along a ridge of size about 20$\\arcmin \\times $3$\\arcmin$ ($\\sim$ 10 $\\times$ 1.5 pc) and are associated with the peaks in the main ridge of molecular gas. The ridge can be seen in many energy bands ranging from 843 MHz radio to several keV X-rays. The standard tracers of SF regions such as [C~II] 158$\\mu$m, [O~I] 145$\\mu$m, and  [O~I] 63$\\mu$m, as well as rich molecular emission spectra revealed structured emission of \\ngc with clumps of far-infrared (FIR) sources. Recently Kraemer and Jackson (1999) reviewed molecular gas in \\ngc and constructed detailed maps of the CO, CS, and NH$_3$ emission regions. They found a complex structure of the gas distribution with a number of molecular filaments and bubbles. High resolution VLA observations of \\ngc by Carral \\etal (2002) revealed shell-like structures of diameters from 0.12 to 3.5 pc probably tracing stellar winds. Nonthermal radio emission features, H$_2$O, OH, and methanol maser sources, as well as high magnetic fields (about 200 $\\mu$G) were observed in \\ngc (e.g. Sarma \\etal 2000 and references therein). We adopt the distance of $\\sim$~1.7~kpc to \\ngc (Neckel 1978).  McBreen \\etal (1979) discovered a bright ($\\sim$1.9$\\times$10$^5$ \\Lsun bolometric luminosity) source of IR emission in the southwestern part of the SF ridge which was not detected in previous IR observations of the region. The source position is associated with H$_2$O and OH masers. Bipolar structure of extended H$_2$ emission and an energetic outflow were found in the source (as well as in a nearby FIR source). The nature of the apparent variability was attributed by McBreen \\etal (1979) to a sudden luminosity onset during early SF phenomena or, alternatively, to a supernova (SN) hidden in a molecular cloud. Young SN remnants (SNRs) possibly hidden in dense molecular clouds can be identified by their gamma-ray emission lines from $^{44}$Ti with the \\isgri\\ detector (Lebrun \\etal 2003). Hard X-ray observations are a perfect tool to look for a SNR hidden in a dense molecular cloud. Being compact sites of powerful kinetic energy release and fast shock waves generated by stellar winds and SNRs, massive SF regions are expected to have non-thermal emission components (e.g. Bykov 2001). This possibility strongly motivates \\int\\ data analysis of \\ngc\\s1.  \\ngc was observed with \\asca\\ (Matsuzaki \\etal 1999; Sekimoto \\etal 2000). X-ray emission was detected from the SF ridge and from the source AXJ~1720.3-3544 that is about 10\\arcmin\\ to the north of the center of the ridge. The emission spectrum of the ridge sources was fitted with a thermal hot plasma of temperature $\\sim$ 9 keV. The source AXJ~1720.3-3544 was identified as the B0.5e star CD-35 11482. Recently Ezoe \\etal (2005) analyzed \\chan\\ data on \\ngc\\s1. In addition to 800 point sources in the field they found a diffuse X-ray emission region of 5$\\times$9 pc size and of luminosity 2$\\times 10^{33} \\ergs$. The authors suggested that thermal emission of several keV is due to plasma heated by stellar wind shocks, while a flat continuum is due to accelerated particles.  However, it is not easy to distinguish hot thermal gas from nonthermal emission generated by accelerated particles with an energy band that is limited to 10 keV. We present below the first hard X-ray (3-80 keV) images and spectra of \\ngc made with \\int.  \\begin{figure*} \\includegraphics[width=1.0\\textwidth]{fig1.eps} \\caption{a) {\\it ROSAT} RASS 0.4-2.4 keV map with {\\it MOST} MGPS 843 MHz contours. Massive stars, visible in the optical band are shown by boxcircles (O stars) and boxes (B stars). b) \\chan\\ 7-8 keV map smoothed with a 2 pixel Gaussian kernel and {\\it VLA} NVSS 1.4 GHz contours. The background region used for \\chan\\ analysis is marked as well as the main sources of \\chan\\ emission within the \\isgri\\ excess. c)~\\jemx\\ 3-10 keV map with \\asca {\\it GIS} 6-10 keV contours and {\\it VLA} GPS 8.35 GHz contours. d) \\isgri\\ 20-60 keV map with DSS-R contours. The P-shaped region with circles indicates the \\ngc\\ ridge with SF complexes as seen in 71$\\mu$m IR band by Loughran et al. (1986). Bright \\isgri\\ pixels are shown on panels a) -- c) as dashed rectangles.} \\label{fig:image} \\end{figure*} ",
        "conclusions": "1. We detected a hard X-ray emission source in the galactic star-forming region \\ngc with the \\jemx and {\\it IBIS/}\\isgri\\ telescopes aboard the International Gamma-Ray Astrophysics Laboratory \\int. The source has a nonthermal spectrum at least up to 100 keV.  2. From the multiwavelength analysis of the complex \\ngc region we concluded that the source may be associated both with the background, likely extragalactic, radio source NGC 6334B projected onto the \\ngc SFR, and with an extended HII region, associated with bright IR source and a radio shell NGC 6334A."
    },
    "0512/astro-ph0512557_arXiv.txt": {
        "abstract": "PPak is a new fiber-based Integral Field Unit (IFU), developed at the Astrophysical Institute Potsdam, implemented as a module into the existing PMAS spectrograph. The purpose of PPak is to provide both an extended field-of-view with a large light collecting power for each spatial element, as well as an adequate spectral resolution. The PPak system consists of a fiber bundle with 331 object, 36 sky and 15 calibration fibers. The object and sky fibers collect the light from the focal plane behind a focal reducer lens. The object fibers of PPak, each 2.7 arcseconds in diameter, provide a contiguous hexagonal field-of-view of 74 $\\times$ 64 arcseconds on the sky, with a filling factor of 60\\%. The operational wavelength range is from 400 to 900~nm. The PPak-IFU, together with the PMAS spectrograph, are intended for the study of extended, low surface brightness objects, offering an optimization of total light-collecting power and spectral resolution. This paper describes the instrument design, the assembly, integration and tests, the commissioning and operational procedures, and presents the measured performance at the telescope. ",
        "introduction": "\\label{sect:intro}  %  3D-spectrographs (3DS), or Integral Field Units (IFU) exist at many observatories, providing spectra for a large number of spatial elements (``spaxel'') within a 2-dimensional field-of-view, rather than only along a traditional 1-dimensional spectrograph slit. Depending on the instrument, up to hundreds or thousands of spectra are recorded simultaneously in any single exposure. While the instrumentation suite is diverse and based on various principles of operation (image slicers, lens-arrays, fiber-bundles or combinations of these), compromises with respect to field-of-view, spatial sampling, wavelength coverage, and spectral resolution have to be made, due to the limited detector space.  Since commissioning in May 2001, the Astrophysical Institute Potsdam (AIP) successfully operates PMAS, the Potsdam Multi-Aperture Spectrophotometer, at the Calar Alto 3.5~m Telescope in southern Spain (Roth et al.\\ 2004, Kelz et al.\\ 2003a).  An overall description of the PMAS instrument is given by Roth et al.\\ (2005), hereafter referred to as paper~I.  While PMAS is a unique spectrophotometer, covering a wide wavelength range from 350~nm to 900~nm, its standard~IFU, a fiber-coupled lens-array, provides 256 spectra and is limited to a maximum integral field-of-view of 16$\\times$16 arcseconds on the sky.  Driven by the ``Disk Mass'' project (Verheijen et al.\\ 2004), which requires imaging spectroscopy of nearby face-on galaxies (with typical sizes of 1 arcminute) at intermediate spectral resolution (of R$\\ge$8000), a science case was put forward to develop a larger IFU for PMAS. Based on the experience with the SparsePak bundle, which was constructed and commissioned for the 3.5m WIYN telescope at Kitt Peak (Bershady et al.\\ 2004, 2005), the PPak (PMAS fiber Package) fiber bundle was designed and built at the AIP in 2003 as part of the ULTROS project (ultra-deep optical spectroscopy with PMAS). This new IFU was produced on a short timescale of approximately six months and with a budget of less than 20.000 Euro for the hardware components (mainly lenses, filters and fibers). PPak was successfully integrated within PMAS in December 2003 and commissioned in spring 2004. The PPak-mode of PMAS is now fully operational and routinely employed for the Disk Mass project, as well as for a variety of other common user programmes, that require large integral-field spectroscopy.  \\begin{table*}% \\tabletypesize{\\small} \\caption{Selected IFU instrumental parameters with corresponding spectral capabilities. \\label{tab:IFUs}} \\begin{tabular}{lcccrccrcc} \\tableline \\tableline Instrument &     Telescope     &   FoV   &      Spaxel      & Spaxel & Filling & Range\\tablenotemark{4}& Resol.\\tablenotemark{4} &\\multicolumn{2}{c}{Grasp}          \\\\ &      diameter & (max.)              & size\\tablenotemark{1}                  & number\\tablenotemark{2} & factor\\tablenotemark{3}  & ($\\lambda_{cov}$)              & ($\\lambda/\\Delta \\lambda$)       &    specific\\tablenotemark{5}     &      total\\tablenotemark{6}      \\\\ &       [m] &   [arcsec]   &     [arcsec]     & &         &     [\\AA]     &        &[arcsec$^2$m$^2$]&[arcmin$^2$m$^2$]\\\\ \\tableline PMAS-PPak{$^7$}  & CA     \\hfill 3.5 & 74$\\times$64 &  2.68 $\\oslash$  & 331+36 &   0.60  &     400       &  8000  &  47\\phantom{.4} &  4.23 \\\\ SparsePak  & WIYN   \\hfill 3.5 & 72$\\times$71 &  4.69 $\\oslash$ &75+\\phantom{3}7&   0.25  &     260       & 12000  & 138\\phantom{.4} &  2.87 \\\\ DensePak   & WIYN   \\hfill 3.5 & 45$\\times$30 &  2.81 $\\oslash$ &91+\\phantom{3}4&   0.42  &     260       & 20000  &  49\\phantom{.4} &  1.25 \\\\ INTEGRAL   & WHT    \\hfill 4.2 & 34$\\times$29 &  2.70 $\\oslash$  & 115+20 &   0.67  &     360       &  4200  &  73\\phantom{.4} &  2.32 \\\\ \\tableline VIMOS\t   & VLT    \\hfill 8.2 & 54$\\times$54 & 0.67$\\times$0.67 &          6400 &   1.00  &     350       &   220  &  23\\phantom{.4} & 40.50\\phantom{7} \\\\ SAURON     & WHT    \\hfill 4.2 & 41$\\times$33 & 0.94$\\times$0.94 & 1577 &   1.00  &     540       &  1250  &  11\\phantom{.4} &  4.77           \\\\ SPIRAL{$^8$}\t   & AAT    \\hfill 3.9 & 22$\\times$11 & 0.7 $\\times$0.7  &           512 &   1.00  &     330       &  7600  &   5.4           &  0.77           \\\\ PMAS-LARR{$^9$}  & CA     \\hfill 3.5 & 16$\\times$16 & 1.0 $\\times$1.0  & 256 &   1.00  &     700       & 6000  &   8.2           &  0.58           \\\\ OASIS\t   & WHT    \\hfill 4.2 & 17$\\times$12 & 0.42 hex.        &    $\\sim$1100 &   1.00  &     370       &  2650  &   2.3           &  0.71           \\\\ GMOS\t   & Gemini~ \\hfill 8.1 &  7$\\times$5  & 0.2 hex.         &          1000+500 & 1.00  &     280       &  1700  &   1.8           &  0.49           \\\\ \\tableline \\\\ \\end{tabular} \\small{$^1$}{~corresponding value to the max. FoV, may depend on fore-optics magnification} \\\\ \\small{$^2$}{~dedicated sky-spaxels are listed separately}\\\\ \\small{$^3$}{~bare fiber bundles (upper part); lens-array-types (lower part)} \\\\ \\small{$^4$}{~selected values only; may span wide range depending on configuration and wavelength} \\\\ \\small{$^5$}{~specific grasp = telescope area [m$^2$] $\\times$ spaxel size [arcsec$^2$]} \\\\ \\small{$^6$}{~total grasp = telescope area [m$^2$] $\\times$ spaxel size [arcsec$^2$] $\\times$ number of spaxels } \\\\ \\small{$^7$}{~PPak-IFU and PMAS spectrograph with 2nd order gratings} \\\\ \\small{$^8$}{~SPIRAL with Littrow spectrograph (de-commissioned)} \\\\ \\small{$^9$}{~Lens-array IFU and PMAS spectrograph with 1st order gratings} \\\\ \\end{table*}  As a bare fiber-bundle IFU, PPak is based on earlier developments, such as DensePak (Barden \\& Wade 1988), INTEGRAL (Arribas et al.\\ 1998) and SparsePak (Bershady et al.\\ 2004). Similar to the above instruments, PPak opts for rather large fibers, that can not properly sample the (seeing-limited) image, but collect more flux and allow for wide fields. PPak and SparsePak span 74$\\times$64 and 72$\\times$71 arcseconds on the sky respectively and provide the largest fields-of-view of any IFU available worldwide (see Table~\\ref{tab:IFUs}). Additionally, a single PPak fiber with 5.7~arcsec$^2$ on the sky, collects twice the amount of light at the 3.5m telescope, than a single spatial element of the VIMOS-IFU (LeFevre et al.\\ 2003) at the 8.2m-VLT (see specific grasp in Table~\\ref{tab:IFUs}). PPak is attached to the efficient PMAS fiber-spectrograph, that provides resolutions R=$\\lambda/\\Delta \\lambda$ from 800 to 8000, or corresponding spectral powers (defined as the product of resolution times the number of spectral resolution elements N$_{\\Delta \\lambda}$ $\\lambda/\\Delta \\lambda$) between 10$^6$ to 10$^7$. The PPak fibers and optics are optimized for wavelengths between 400--900~nm.  \\begin{figure*}% \\epsscale{2.0} \\plotone{Figure1-07.eps} \\caption{The principle of operation for the two PMAS IFUs at the Cassegrain focal station. The LARR-IFU is on-axis and consists of fore-optics, a lens-array (LARR) and a fiber-bundle (dashed lines). The PPak-IFU is located off-axis and features a focal reducer lens (FORED) and a bare fiber-bundle. A dedicated PPak calibration unit can illuminate additional fibers (dashed-dotted lines). Both bundles connect to one spectrograph (solid outline). The acquisition and guiding system (dotted outline) can be used by both modes (see text for further explanation).} \\label{fig:ppak_principle} \\epsscale{1.0} \\end{figure*}  Therefore, PPak is ideally suited for spectroscopic studies of extended astronomical objects with low surface brightness, such as the outskirts of spiral galaxies, where sufficient signal is more important than detailed spatial resolution. The PMAS+PPak configuration offers a powerful combination of light-collecting power or grasp, wavelength range, and spectral resolution. \\\\  This paper is organized as follows: \\S~\\ref{sect:design} presents the instrument and its opto-mechanical design. \\S~\\ref{sect:mai} summarizes the manufacture, assembly, and integration of the PPak components.  \\S~\\ref{sect:operations} describes the operational procedure during observation, the data reduction, and visualization tools. The instrument performance at the telescope and test results are given in \\S~\\ref{sect:performance}. ",
        "conclusions": "\\label{sect:summary}  A new Integral-Field-Unit, based on the fiber-bundle technique, providing  high grasp and a large field was developed and successfully commissioned for the existing PMAS 3D-instrument. The central PPak-IFU features 331 object-fibers, which, projected by the 3.5~m Calar Alto telescope, span a hexagonal field-of-view of 74$\\times$64~ arcseconds with a filling factor of 60\\%. The individual spaxel (fiber) size is 2$^{\\prime\\prime}$.7 across, yielding a total grasp of 15200~arcsec$^2$m$^2$ at this telescope. An additional 36 fibers are distributed over six sky-IFUs, which surround the main IFU at a distance of 72$^{\\prime\\prime}$ from the field center, allowing a good coverage and subtraction of the sky background. For calibration purposes, 15 fibers can be illuminated independently with arc lamps during a science exposure, and can  keep track of spectral resolution and image shifts. A summary of the technical parameters is given in Table~\\ref{tab:pmas_param}. Further details regarding the PMAS spectrograph, available gratings and filters are given in paper~I, or can be found online\\footnote{\\url{http://www.caha.es/pmas}}.  The combination of spaxels with high grasp and the PMAS spectrograph with high efficiency and wide wavelength coverage, makes PPak a powerful tool for the study of extended low-surface brightness objects, which require a high light collecting power and a large field-of-view. Fig.~\\ref{fig:UGC463}, gives an example of the galaxy UGC~463, that was observed for the Disk Mass project. Despite the rather crude sampling of the fibers, the basic morphological structures of the galaxy seen in the POSS-II image (spiral arms, stars clumps...), are clearly visible in the PPAK reconstructed image. Apart from the ability to create mono- and polychromatic images from the resulting data, one exposure with PPak yields 331 spatially resolved spectra  of the target. The high number of fibers at the outer and fainter parts of the galaxy, offers the observer the option to adaptively bin spaxels as to increase the signal-to-noise further.  \\begin{figure*}[ht!] \\begin{center} \\includegraphics[width=0.3\\textwidth,angle=-90]{Figure21-03.ps} % \\caption{Comparison between the POSS-II R-band image (left panel) and the PPAK reconstructed image of the galaxy UGC 463 (right panel). The PPAK data were obtained using the V300 grating, centered at $\\approx$5300\\AA. Once reduced, the reconstructed image was created using E3D (S\\'anchez,\\ 2004): a 3D cube was created adopting a natural neighbor interpolation scheme to a common grid of 1$''$.35/pixel. After that, a 2D image was produced by co-adding the flux in the wavelength range between 4500 and 6000~\\AA.} \\label{fig:UGC463} \\end{center} \\end{figure*}  During 2004, PPak was available on a shared-risk basis, as its usage involved a complex change-over between the lens-array and PPak fibers, which needed to be done by AIP staff. With the integration of the two parallel fiber-slits, the PPak-IFU was permanently installed and is offered as common-user instrument from 2005 onwards. Two further upgrades are scheduled: a new mount to ease the exchange of order separating filters, and adjustments to the PPak data reduction software. Since commissioning, PPak attracted considerable interest from several observers, with the result, that 9 observing runs with a total of 26 nights, plus 9 `service buffer A' nights, were granted to PMAS-PPAK within its first year by the Calar Alto TAC. In 2005, approximately 50\\% of the allocated time for PMAS is used with the PPak-IFU. Throughout these runs, the PPak module worked without failure. \\begin{table}% \\caption{Summary of main PPak + PMAS instrumental parameters.} \\begin{center} \\begin{tabular}{lr} \\tableline \\tableline {\\bf PPak-IFU:} & \\\\ \\tableline principle design\t& focal reducer + fiber-bundle \\\\ focal reducer lens\t& F/10 to F/3.3, dia=50~mm \\\\ plate-scale \t\t& $17''.85~$/mm \\\\ fiber configuration\t& 331 object, 36 sky, 15 calibration \\\\ PPak - FoV \t\t& $74'' \\times 64''$ (hexagonal packed) \\\\ spatial sampling\t& $2''.68$ per fiber diameter \\\\ fiber pitch\t\t& $3''.5$ fiber-to-fiber \\\\ IFU-filter\t\t& 2~inch / 50~mm round filter \\\\ PPCU-filter\t\t& 1~inch / 25~mm round filter \\\\ wavelength range\t& 400--900~nm (`dry' fibers) \\\\ \\tableline % {\\bf PMAS spectrograph:} \t& \\\\ \\tableline PPak fiber slit \t& 0.15 $\\times$ 94 mm, 382 fibers \\\\ slit-filter\t\t& 140 $\\times$ 35~mm filter \\\\ collimator \t\t& fully refractive 450 mm, F/3  \\\\ camera \t\t\t& fully refractive 270 mm, F/1.5  \\\\ reflective gratings\t& 1200, 600 and 300 l/mm \\\\ detector \t\t& SITe 2k $\\times$ 4k, 15 $\\mu m$ pixels \\\\ linear dispersion\t& 0.53, 1.2, 2.6 \\AA /pixel (m=1)  \\\\ resolution\t\t& 0.3 \\AA /pixel, R$\\approx$8000 (m=2)  \\\\ wavelength coverage \t& 600-3400 \\AA /frame (1st order) \\\\ & 400 \\AA /frame (2nd order)\\\\ \\tableline % {\\bf PMAS A\\&G camera:} \t& \\\\ \\tableline A\\&G FoV \t\t& 3.4 $\\times$ 3.4 arcminutes  \\\\ A\\&G plate scale \t& 0.2 arcseconds/pixel \\\\ A\\&G-filter\t\t& 4 $\\times$ 2~inch/50~mm round filters \\\\ \\tableline \\end{tabular} \\end{center} \\label{tab:pmas_param} \\end{table}  \\vspace{5mm}"
    },
    "0512/astro-ph0512007_arXiv.txt": {
        "abstract": "We present mid-infrared observations of active galactic nuclei (AGN) in the Great Observatories Origins Deep Survey (GOODS) fields, performed with the Spitzer Space Telescope.  These are the deepest infrared and X-ray fields to date and cover a total area of $\\sim$0.1 square degrees. AGN are selected on the basis of their hard (2-8 keV) X-ray emission. The median AGN infrared luminosity is at least 10 times larger than the median for normal galaxies with the same redshift distribution, suggesting that the infrared emission is dominated by the central nucleus. The X-ray to infrared luminosity ratios of GOODS AGN, most of which are at 0.5$\\la$$z$$\\la$1.5, are similar to the values obtained for AGN in the local Universe. The observed infrared flux distribution has an integral slope of $\\sim$1.5 and there are 1000 sources per square degree brighter than $\\sim$50~$\\mu$Jy at $\\sim$3-6~microns. The counts approximately match the predictions of models based on AGN unification, in which the majority of AGN are obscured. This agreement confirms that the faintest X-ray sources, which are dominated by the host galaxy light in the optical, are obscured AGN. Using these Spitzer data, the AGN contribution to the extragalactic infrared background light is calculated by correlating the X-ray and infrared catalogues. This is likely to be a lower limit given that the most obscured AGN are missed in X-rays. We estimate the contribution of AGN missed in X-rays, using a population synthesis model, to be $\\sim$45\\% of the observed AGN contribution, making the AGN contribution to the infrared background at most $\\sim$2-10\\% in the 3-24 micron range, depending on wavelength, lower than most previous estimates. The AGN contribution to the infrared background remains roughly constant with source flux in the IRAC bands but decreases with decreasing flux in the MIPS 24~$\\mu$m band, where the galaxy population becomes more important. ",
        "introduction": "Hard X-ray surveys are not strongly affected by dust obscuration and thus provide a relatively unbiased view of the AGN phenomenon. Deep Chandra X-ray surveys have allowed us to test the basic prediction of population synthesis models of the X-ray background, namely, that a combination of obscured and unobscured AGN is needed to explain the X-ray background hard spectral shape \\citep{setti89,gilli01}. According to the AGN unification paradigm, obscuration comes from optically thick dust blocking the central engine along some lines of sight. The temperature in this structure, which can range up to 1000~K (dust sublimation temperature), and the roughly isotropic emission toward longer wavelengths should make both obscured and unobscured AGN very bright in the mid-far infrared bands.  This prediction was tested by IRAS (e.g., \\citealp{rush93}) and ISO (e.g., \\citealp{fadda02,matute02,franceschini02}) observations, which discovered a population of very luminous infrared sources coincident with X-ray emitters and therefore very good candidates to be obscured AGN. With relatively bright flux limits, however, these surveys sample the nearby Universe or extremely luminous, rare AGN. Very deep surveys were possible with ISO (e.g., \\citealp{fadda02}), but only over a very small area. With the Spitzer Space Telescope, it is now possible to perform very deep infrared surveys, $\\sim$100-1000 times deeper than typical IRAS observations, over appreciable areas, $\\sim$10-1,000 times more area than the deep ISO surveys, thus permitting us, for the first time, to study a representative sample of the AGN population at cosmologically-significant distances.  Since AGN comprise most of the X-ray background and are also very bright infrared sources, it is natural to ask what fraction of the extragalactic infrared background they may contribute. Based on previous observations, the predicted contribution of AGN to the infrared background ranged from $\\sim$20\\% \\citep{fabian99} to just a few percent \\citep{xu01}, depending mostly on the assumptions for the amount of energy re-radiated in the infrared and the spatial density of obscured AGN. The deeper infrared observations possible with Spitzer are important for understanding the true contribution of AGN to the extragalactic infrared background light.  Multiwavelength observations of a large, uniformly selected sample of moderate-luminosity AGN at cosmologically significant distances provide a complete picture of the AGN emission across the frequency spectrum and establish the total contribution of these sources to the energy budget of the Universe. This is one of the main goals of the Great Observatories Origins Deep Survey (GOODS), a 0.1-square-degree study divided in two fields that overlap with the deepest X-ray observations, the Chandra Deep Fields North \\citep{brandt01} and South \\citep{giacconi01}. It consists of deep imaging in the optical with the Hubble Space Telescope \\citep{giavalisco04} and in the infrared with the Spitzer Space Telescope \\citep{dickinson02}. These very deep multiwavelength observations allow a highly complete, relatively unbiased, view of the AGN phenomenon and a test of the AGN unification paradigm, well established in the local Universe, at higher redshifts.  In this paper we present Spitzer observations of the AGN in the GOODS fields, from 3.6 to 24 microns. In \\S~2 we describe the Spitzer IRAC and MIPS observations. Results and discussion are presented in \\S~3, addressing the infrared luminosities of AGN and normal galaxies, the X-ray to infrared ratios for GOODS and local AGN, the infrared number counts of AGN, and the AGN contribution to the extragalactic infrared background light. Our conclusions are presented in \\S~4. Throughout this paper we assume $H_0=70$ km s $^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda =0.7$. ",
        "conclusions": "We present here Spitzer observations of the moderate-luminosity, moderate- to high-redshift AGN in the GOODS fields, which include the deepest infrared and X-ray observations to date. In general these sources have similar X-ray to infrared luminosity ratios as AGN in the local Universe. AGN are $\\sim$10-15 times brighter in the infrared than normal galaxies, although there is some overlap between the luminosity distribution of AGN and normal galaxies, indicating that for low-luminosity AGN the contribution to the infrared light from the host galaxy can be very important and may well dominate the observed emission. This can explain the observed trend of the X-ray-to-infrared luminosity ratio to increase with X-ray luminosity.  We report the observed infrared flux distribution for GOODS AGN, which has a slope of $\\simeq$1.5 to 0.9 for fluxes higher than $\\sim$100~$\\mu$Jy, with a source density of $\\simeq$1000 AGN per square degree at a depth of $\\sim$50 $\\mu$Jy in the IRAC bands. Comparing to the expectations of a unification model for AGN, we find reasonable agreement, showing that the basic properties of the moderate to high redshift AGN population are well explained by the same unification paradigm that fits local sources very well, but with a factor of two less dust in the circumnuclear region.  From the Spitzer observations, we calculate the minimum AGN contribution to the extragalactic near-mid infrared background, obtaining a lower value than previously estimated, ranging from 2\\% to 10\\% in the 3-24 micron range. Accounting for heavily obscured AGN that are not detected in X-ray, the AGN contribution to the infrared background increases by $\\sim$45\\%, to $\\sim$3-15\\%."
    },
    "0512/astro-ph0512231_arXiv.txt": {
        "abstract": "Detection of ultra-high energy neutrinos will be useful for unraveling the dynamics of the most violent sources in the cosmos and for revealing the neutrino cross-section at extreme energy. If there exists a Greisen-Zatsepin-Kuz'min (GZK) suppression of cosmic-ray events above $\\egzk\\sim 5 \\times 10^{19}$~eV, as predicted by theory, then the only messengers of energies beyond $\\egzk$ are neutrinos. Cosmic neutrino fluxes can initiate air-showers through interaction in the atmosphere, or in the Earth.  Neutrino trajectories will be downgoing to nearly horizontal in the former case, and ``Earth-skimming'' in the latter case. Thus it is important to know the acceptances (event rate/flux) of proposed air-shower experiments for detecting both types of neutrino-initiated events. We calculate these acceptances for fluorescence detectors, both space-based as with the EUSO and OWL proposals, and ground-based, as with Auger, HiRes and Telescope Array. The neutrino cross-section $\\sig$ is unknown at energies above $5.2 \\times 10^{13}$~eV. Although the popular QCD extrapolation of lower-energy physics offers the cross-section value of $0.54 \\times 10^{-31}\\,(E_\\nu/10^{20}{\\rm eV})^{0.36}\\,{\\rm cm}^2$, new physics could raise or lower this value. Therefore, we present the acceptances of horizontal (HAS) and upgoing (UAS) air showers as a function of $\\sig$ over the range $10^{-34}$ to $10^{-30}$ cm$^2$. The dependences of acceptances on neutrino energy, shower-threshold energy, shower length, and shower column density are also studied. We introduce a cloud layer, and study its effect on rates as viewed from space and from the ground. For UAS, we present acceptances for events over land (rock), and over the ocean (water). Acceptances over water are larger by about an order of magnitude, thus favoring space-based detectors. We revisit the idea of Ref.~\\cite{KW} to infer $\\sig$ at $E_\\nu\\gsim 10^{20}$ from the ratio of HAS-to-UAS events, and obtain favorable results. Included in our UAS calculations are realistic energy-losses for taus, and Earth-curvature effects. Most of our calculation is analytic, allowing insight into the various subprocesses that collectively turn an incident neutrino into an observable shower. ",
        "introduction": "Detection of ultra-high energy ($E_\\nu > 10^{18}$~eV~$\\equiv$~EeV) neutrinos is important for several reasons. First of all, neutrino primaries are not deflected by magnetic fields and so should point back to their cosmic sources. This contrasts with cosmic-rays, which are charged and follow bent trajectories. Secondly, well above the Greisen-Zatsepin-Kuzmin (GZK) energy of $\\egzk\\sim 5\\times 10^{19}$~eV~\\cite{G66,ZK66}, they may be the only propagating primaries. As such, they may be the only messengers revealing the ultimate energy-reach of extreme cosmic accelerators, generally believed to be powered by black holes. Above $\\egzk$, the GZK suppression~\\cite{G66,ZK66,Zayyad1,Letessier1} of cosmic-rays results from the resonant process $N+\\gamma_{\\rm CMB}\\rarr \\Delta\\rarr N +\\pi$; $\\egzk$ is the lab-frame energy corresponding to the kinematic threshold $\\sqrt{s}=M_\\Delta$ for excitation of the intermediate $\\Delta$ resonance. A handful of cosmic-ray events have been detected with estimated energies exceeding $10^{20}$~eV. The record energy is the famous Fly's Eye event at $3\\times 10^{20}$~eV~\\cite{FlysEye}. The observable neutrino spectrum could extend to much higher energies. Thirdly, in contrast to cosmic-rays and photons, neutrinos are little affected by the ambient matter surrounding the central engines of Nature's extreme accelerators. Accordingly, neutrinos may carry information about the central engine itself, inaccessible with other primaries. In principle, neutrinos may be emitted from close to the black hole horizon, subject only to energy-loss due to gravitational redshifting. An analogy can be made to solar studies performed with photons versus neutrinos. The photons are emitted from the outer centimeter of the Sun's chromosphere, while the neutrinos are emitted from the central core where fusion powers the Sun. Fourthly, neutrinos carry a quantum number that cosmic-rays and photons do not have - flavor. Neutrinos come in electron, muon, and tau flavors. One may think of this ``extra'' flavor degree of information as the neutrino's superb analog to polarization for the photon, or nucleon number $A$ for the cosmic-ray. Each of these attributes, flavor, polarization, and nucleon number, carries information about the nature and dynamics of the source, and about the environment and pathlength of the inter-galactic journey. The flavor ratios of cosmic neutrinos are observable~\\cite{flavorID}. Several papers have recently analyzed the benefits that neutrino flavor identification offers for unraveling the dynamics of cosmic sources~\\cite{pimuchain}. The fifth reason why ultra-high energy neutrino primaries traveling over cosmic distances are interesting is that such travel allows studies of the fundamental properties of neutrinos themselves. For studying some properties of the neutrino, such as neutrino stability/lifetime~\\cite{nulifetime}, pseudo-Dirac mass patterns~\\cite{pDirac}, it is the cosmic distance that is essential; for other properties, it is the extreme energy that is essential. A clear example of the latter is any attempt to determine the neutrino cross-section at energies beyond the reach of our terrestrial accelerators.  In this paper we will examine the potential for cosmic-ray experiments designed to track ultra-high energy air-showers by monitoring their fluorescence yield~\\cite{fluoresence}, to detect horizontal air-showers (HAS) and upgoing air-showers (UAS) induced by a cosmic neutrino flux. We will also study the ability of these experiments to infer the neutrino-nucleon cross-section $\\sigma_{\\nu N}$ at energies above $10^{19}$~eV, from the ratio of their UAS and HAS events. Such energies are orders of magnitude beyond the energies accessible to man-made terrestrial accelerators. From the point of view of QCD, such a cross-section measurement would be an interesting microscope into the world of small-x parton evolution. The neutrino cross-section above $10^{19}$~eV could agree with any of the various QCD-motivated extrapolations that have been published~\\cite{GQRS,sigma20}, or not. The cross-section could also be quite different than the extrapolations. For example, if a new threshold is crossed between terrestrial neutrino energies $\\sim 100$~GeV, and the extreme energies reached by cosmic-rays, ~$\\sim 10^{11}$~GeV, then the cross-section could much exceed the QCD-extrapolations. On the other hand, saturation effects can significantly reduce the total cross-section at these very high energies~\\cite{saturation}. The nine orders of magnitude increase in lab energy reach corresponds to 4.5 orders of magnitude increase in center-of-momentum energy reach. Even the center-of-momentum energy at the $e-p$~HERA collider is more than three orders of magnitude below the cosmic-ray reach. This remarkable energy reach of cosmic-rays presents ample room for new physics beyond our Standard Model. Proposals for new physics thresholds in this energy region include low-scale unification with gravity, in which neutrino-nucleon scattering produces mini-black holes~\\cite{miniBH} and/or brane-wraps~\\cite{branewraps}, non-perturbative electroweak instanton effects~\\cite{EWinstanton}, compositeness models~\\cite{composite}, a low energy unification scale in string inspired models~\\cite{string}, and Kaluza-Klein modes from compactified extra dimensions~\\cite{extradimensions}. All of these models produce a strongly-interacting neutrino cross-section above the new threshold. Dispersion relations allow one to use low-energy elastic scattering to place constraints on the high-energy cross-section~\\cite{disprelation}, but the constraints are quite weak.  For HAS and UAS, we provide analytical calculations of the event-rate to flux ratio as a function of $\\sigma_{\\nu N}$. This ratio is known as the ``instantaneous experimental acceptance'', with units of {\\sl area$\\times$solid angle}. The time-averaged acceptance includes an experimental ``duty factor,'' the fraction of time that the experiment is functioning. We will not include the duty factor in our calculations of acceptances. We note that acceptances are also sometimes called ``apertures.''  Experimental acceptance offers a very meaningful figure of merit for statistical reach. One has merely to multiply an experiment's acceptance by Nature's flux to arrive at an event rate for the experiment. Multiplying again by the experiment's run time (including the duty factor), one obtains the total number of events.  Acceptance times run time is termed the experimental ``exposure''.  The acceptances we calculate are scalable to large area experiments such as HiRes, Auger, and in the near future Telescope Array, which are anchored to the ground, and to super-large area experiments such as EUSO and OWL, which are proposed to orbit the Earth from space. A horizontal shower, deeply initiated, is the classic signature for a neutrino primary. The weak nature of the neutrino cross-section means that horizontal events begin where the atmospheric target is most dense, low in the atmosphere.  In contrast, the ultra-high energy $pp$ cross-section exceeds 100~mb, so the air-nucleon cross-section exceeds a barn! Even the vertical atmospheric column density provides hundreds of interaction lengths for a nucleon, and so the cosmic-ray interacts high in the atmosphere. The weak nature of the neutrino cross-section also means that the event rate for neutrino-induced HAS is proportional to the neutrino-nucleon cross-section.  For an neutrino-induced UAS, the dependence on neutrino cross-section is more complicated, and more interesting. The Earth itself is opaque for neutrinos with energies exceeding about a PeV of energy.  However, ``Earth-skimming'' neutrinos, those with a short enough chord length through the Earth, will penetrate and exit, or penetrate and interact. In particular, there is much interest in the Earth-skimming process $\\nutau\\rarr\\tau$ in the shallow Earth, followed by $\\tau$ decay in the atmosphere to produce an observable shower. In Ref.~\\cite{KW} it was shown that the rate for the Earth-skimming process $\\nutau\\rarr\\tau$ is {\\sl inversely} proportional to $\\sigma_{\\nu N}$. There it was emphasized that $\\sigma_{\\nu N}$ could be inferred from a measurement of the ratio of HAS to UAS rates.~\\footnote{The prospects of inferring the neutrino-nucleon cross-section in the energy range of 100 TeV - 100 PeV at neutrino telescopes such as IceCube, were studied in Ref.~\\cite{Icecubecs}; prospects at higher energies were studied in Ref.~\\cite{Augercs} for the Auger observatory.} Of course, an implicit assumption is that there is enough neutrino flux at extreme energies to generate HAS and UAS event samples.  The inverse dependence of UAS rate on $\\sigma_{\\nu N}$ is broken by the $\\tau\\rarr$~{\\sl shower} process in the atmosphere. As the cross-section decreases, the allowed chord length in the Earth increases, and the tau emerges with a larger angle from the Earth's tangent plane. This in turn provides a smaller path-length in air in which the tau may decay and the resulting shower may evolve. This effect somewhat mitigates the inverse dependence of the UAS on $\\sigma_{\\nu N}$.  Ref.~\\cite{KW} provided an approximate calculation of the whole UAS process, and gave an approximate result for the dependence of the HAS/UAS ratio on $\\sigma_{\\nu N}$. In this work, we improve upon Ref.~\\cite{KW} in several ways. We include the energy dependences of the tau energy-losses in the Earth, and of the tau lifetime in the atmosphere. For the energy-losses, we distinguish between tau propagation in earth rock and propagation in ocean water. These calculations are carried out in Section~\\ref{sec:rates}. On the issue of shower development, we incorporate the dependence of atmospheric density on altitude.  We also impose requirements on the resulting shower such that a sufficiently long visible shower-length is projected onto the Earth's tangent plane, thus meeting experimental requirements for visibility. This is done in Section~\\ref{sec:visibleAS}. In the case of the upgoing showers, the pathlength of the pre-decayed tau may be so long that the Earth's curvature enters into the altitude dependence. We include the non-negligible correction from curvature in Section~\\ref{sec:visibleAS}. We include the partial loss of visibility due to high cirrus or low cumulus cloud layers in Section~\\ref{sec:clouds}. It is estimated that clouds will obscure the viewing area about 60-70\\% of the time. For ground-based observation, it is mainly  the low-lying cumulus clouds that limit visibility. For space-based observation, it is mainly the high cirrus clouds that limit visibility.~\\footnote{In fact, low-lying cumulus clouds may aid in HAS identification for space-based observing. When the HAS hits the cloud layer, diffuse reflection of the forward Cerenkov cone can be seen as a one-time ``\\v{C}erenkov flash''. The time of the flash and the measured height of the cloud then provide the absolute (t,z) coordinates of the shower.} And in Section~\\ref{sec:cloudsncurvature}, we combine the corrections from clouds with that from the Earth's curvature.  Our results are illustrated in a series of plots of acceptances, for ground-based and space-based experiments, versus neutrino-nucleon cross-section, in Section~\\ref{sec:results}. Situations with and without cloud layers are analyzed, as are events over solid earth and over the ocean. Incident neutrino energies, energy thresholds for experimental detection of the air-shower, and various shower-trigger parameters are varied.  Earth-curvature effects are included in our UAS calculations. These reduce the event rate. Next comes the discussion Section~\\ref{sec:discussion}. It presents several small issues, and includes a comparison of our work with prior work. A final Section recaps our conclusions. Some of the more tedious but necessary formulas are derived in an Appendix.  The reader who believes that a picture (or four) is a worth thousand words may wish to jump to Section~\\ref{sec:results}. Such a reader especially may find it useful to reference Tables~\\ref{table:variables} and \\ref{table:parameters}, where the variables and parameters are defined.  Among our conclusions, we find that the HAS/UAS ratio is or order of unity for cross-section values very near to the commonly extrapolated value of $0.5\\times 10^{-31}\\,{\\rm cm}^2$ at $E_\\nu\\sim 10^{20}$~eV. This is fortunate, for it offers the best possibility that both HAS and UAS rates can be measured, and a true cross-section inferred. We display our HAS and UAS acceptance plots for a cross-section range from superweak $10^{-34}\\,{\\rm cm}^2$ to a microbarn, $10^{-30}\\,{\\rm cm}^2$. This range includes the QCD-extrapolations of $\\sig$, and the region of the HAS/UAS cross-over. It also encompasses any effects of new neutrino physics, either increasing  or decreasing $\\sig$. The highest energy for which the neutrino cross-section has been measured is that at the HERA accelerator. The measurement is $\\sig\\sim 2 \\times 10^{-34} \\,\\rm{cm}^2$ at $\\sqrt{s}=314$ GeV~\\cite{HERAcs}, the latter corresponding to an energy on fixed target of $5.2 \\times 10^{13}$~eV (52~TeV). It is hard to imagine that $\\sig$ at $10^{20}$~eV would not have grown beyond the HERA value. Even so, the acceptances shown for superweak cross-sections may have some relevance to a possible WIMP flux~\\cite{WIMPaccept}. Modeling of a WIMP event rate requires modifications in the shower development for HAS, and in the chain WIMP$\\rarr$UAS, that we do not pursue here.  \\begin{table} \\begin{tabular}{|l|l|} \\hline $L$ & chord length of $\\nu$ trajectory through Earth \t\\\\ \\hline $\\zint$ & vertical height (depth) of HAS (UAS) $\\nu$ interaction \\\\ \\hline $\\zdk$ & altitude of upgoing $\\nutau$ decay (no Earth-curvature) \\\\ \\hline $\\zpdk$ & altitude of upgoing $\\nutau$ decay including Earth-curvature\\\\ \\hline $z_U$ & maximum visible shower altitude (HAS $\\neq$ UAS)\\\\ \\hline $z_L$ & minimum visible shower altitude (HAS $\\neq$ UAS)\\\\ \\hline $\\zb$ & altitude where shower first attains threshold brightness (HAS $\\neq$ UAS)\\\\ \\hline $\\ze$ & altitude where shower extinguishes (HAS $\\neq$ UAS)\\\\ \\hline $\\zcc$ & critical altitude for suppression from cloud layer\\\\ \\hline $\\zpbuas$ & altitude where UAS attains threshold brightness, including Earth-curvature \\\\ \\hline $\\zpeuas$ & altitude where UAS extinguishes, including Earth-curvature \\\\ \\hline $\\thz$ & zenith angle of HAS event   \t\t\\\\ \\hline $\\thn$ & nadir angle of UAS event (no Earth-curvature)\\\\ \\hline $\\thhor=\\frac{\\pi}{2}-\\theta_z$ & horizontal angle of UAS event\\\\ \\hline $\\thnp$ & nadir angle of UAS event including Earth-curvature \\\\ \\hline $\\thhorp$ & horizontal angle of UAS event including Earth-curvature\\\\ \\hline $\\dtot$& total column along chord of Earth \t\t\\\\ \\hline $\\dnu$ & column density of $\\nu$ in the Earth\t\t\\\\ \\hline $\\dtau$& column density of $\\tau$ in the Earth\t\t\\\\ \\hline $\\coshass$ & minimum shower angle, cloud-dependent, for space-observatory\t\t\t\t\\\\ \\hline $\\cosuasg$ & minimum shower angle, cloud-dependent, for ground-observatory\t\t\t\t\\\\ \\hline $\\zhath$ & maximum altitude from which initiated HAS can reach the ground\t\t\t\t\\\\ \\hline $\\zhatu$ & minimum altitude from which initiated UAS can reach $\\zthin$\t\t\t\t\\\\ \\hline $\\zphatu$ & $\\zhatu$ with Earth-curvature included      \\\\ \\hline $\\zhatugd$ & $\\zhatu$ modified for cloud layer above \t\\\\ \\hline $\\zhathsp$& $\\zhath$ modified for cloud layer below \t\\\\ \\hline \\end{tabular} \\caption{\\label{table:variables} List of variables and their meaning. (Conversion between variables $z$ and $w$ is given by $w\\cos\\theta =z$.)} \\end{table}  \\begin{table} \\begin{tabular}{|l|l|l|} \\hline $h$ & scale height of the atmosphere & 8~km \t\\\\  \\hline $\\zg$ & ground altitude, kept as a symbol for later substitutions & 0 \\\\ \\hline $\\zthin$ & altitude beyond which air is too thin to fluoresce significantly & $3h$ \\\\ \\hline $\\beta_{19}$& tau energy-attenuation constant at $E_\\tau=10^{19}$~eV &1.0 (0.55) $\\times 10^{-6}\\,{\\rm cm}^2/{\\rm g}$ for rock (water) \\\\ \\hline $\\alpha$ & exponent of the energy-dependence of $\\beta_\\tau$ & 0.2 \\\\ \\hline $\\dvert$& vertical atmospheric column density & 1,030~g/cm$^2$ \\\\ \\hline $\\dhor$ & horizontal atmospheric column density & 36,100~g/cm$^2$ \\\\ \\hline $\\dmin$ & minimum acceptable shower column density & 300, {\\bf 400}~g/cm$^2$ \\\\ \\hline $\\dmax$ & maximum shower column density at extinction & {\\bf 1200}, 1500~g/cm$^2$\\\\ \\hline $\\lmin$ & minimum acceptable shower length projected on the Earth's surface & {\\bf 10~km}, 5~km \t\\\\ \\hline $\\rfov$ & radius (or half-scale) of the experimental field of view & 230~km \t\t\t\t\t\\\\ \\hline $\\zw$ & depth of ocean &  3.5~km \\\\ \\hline $\\zc$ & altitude of cloud layer & {\\bf 2}, 4, 8, 12~km \\\\ \\hline $E_\\nu$ & incident neutrino energy & ${\\bf 10^{20}}$, $10^{21}$~eV; \\\\ \\hline $\\Esh$ & detector threshold energy & $10^{19}$, ${\\bf 5\\times 10^{19}}$ \\\\ \\hline $\\Eth$ & tau threshold energy & $\\frac{3}{2}\\,(3)\\times\\Esh$ for hadron (electron) mode \\\\ \\hline $\\sig$ & neutrino (or WIMP) cross-section & $10^{-30}$, ${\\bf 10^{-31}}$, $10^{-32}$, $10^{-33}{\\rm cm}^2$  \\\\ \\hline \\end{tabular} \\caption{\\label{table:parameters} List of parameters, their meaning, and their chosen value(s); the bold-faced value is the chosen ``canonical'' value.} \\end{table}  \\begin{table} \\begin{tabular}{|l|l|} \\hline HAS  & Horizontal Air-Shower \t\t\t\\\\ \\hline UAS  & Upgoing (``Earth-skimming'') Air-Shower \t\\\\ \\hline \\hsp & HAS seen from space-based observatory \t\\\\ \\hline \\usp & UAS seen from space-based observatory \t\\\\ \\hline \\hgd & HAS seen from ground-based observatory \t\\\\ \\hline \\ugd & UAS seen from ground-based observatory \t\\\\ \\hline \\end{tabular} \\caption{\\label{table:symbols} List of symbols and their meaning.} \\end{table} ",
        "conclusions": "\\\\ (i) Inference of the neutrino cross-section at and above $10^{20}$~eV from the ratio of UAS and HAS events appears feasible, assuming that a neutrino flux exists at these energies.\\\\ (ii) Space-based detectors enjoy advantages over ground-based detectors for enhancing the event rate. The advantages are a much higher UAS rate over water compared to land, and the obvious advantage that space-based FOV's greatly exceed ground-based FOV's.  Our hope is that space-based fluorescence-detection becomes a reality, so that the advantages of point (ii) can be used to discover/explore the extreme-energy cosmic neutrino flux. According to point (i), part of the discovery/exploration can be the inference of the neutrino cross-section at $E_\\nu\\sim 10^{20}$~eV."
    },
    "0512/astro-ph0512141_arXiv.txt": {
        "abstract": "{} {We observed a new and poorly studied cataclysmic variable (CV) \\object{SDSS~J123813.73-033933.0} to  determine its classification and binary  parameters.  }{  Simultaneous  time-resolved  photometric  and spectroscopic observations were carried out to conduct period analysis and Doppler tomography mapping.}   { From radial velocity measurements of  the  H$\\alpha$  line  we   determined  its  orbital  period  to  be $0.05592\\pm0.00035$ days  (80.53min).  This period is  longer than the first estimate  of 76 min  by Szkody et al.  (2003), but still at the very edge of the period limit for hydrogen-rich CVs.  The spectrum shows  double-peaked Balmer emission lines  flanked by strong broad  Balmer absorption,  indicating  a dominant  contribution by  the white dwarf primary star, and is similar to the spectra of short-period low-mass  transfer WZ Sge-like  systems.  The photometric  light curve shows complex variability.  The system undergoes cyclic brightening up to  0.4 mags, which  are of  semi-periodic nature  with periods  of the order  of 8-12  hours.   We also  detect  a 40.25  min variability  of $\\sim0.15$ mag corresponding to half of the orbital period.  Amplitude of the latter increases with the cyclic brightening of the system.  We discuss the variable accretion rate and  its impact on the hot spot as the   most   probable    reason   for   both   observed   processes.}{ \\object{SDSS~J123813.73-033933.0}  is   preliminary  classified  as  a WZ\\,Sge-like  short  period system  with  low  and unstable  accretion rate.}{}{ ",
        "introduction": "Cataclysmic variables  (CVs) are close  binaries that contain  a white dwarf (WD)  and a  late main-sequence (K-M  spectral type)  star.  The secondary  star fills  its Roche  lobe and  transfers mass  to  the WD through  the inner  Lagrangian point.   More than  1300 CVs  are known presently (Downes et al.  \\cite{Downes}) and orbital periods have been found for more  than 400 systems. The orbital  periods range from days to a minimum  period of about 70 min for systems  with a main-sequence secondary star.  An  important feature that stands out  in the orbital period distribution of  CVs is a sharp short-period  cut-off at 80min, the 'period minimum' (e.g. Barker \\& Kolb \\cite{Barker}).  Less than a dozen AM CVn-type CVs with even shorter periods are interpreted as CVs with helium star donors.  According to the population syntheses, there should be  a significant  number of systems  near the  minimum period, which  have not  been observed  (Kolb \\&  Baraffe  \\cite{Kolb}).  Some calculations predict that 99\\% of the entire CV population should have orbital  periods $<2$h (e.g.   Howell et  al. \\cite{Howell}),  but the number of CVs observed above and below the period gap are similar.  If the population models  are correct, then only a  small fraction of the existing CV population has been discovered so far.  About 400 new  CVs are expected from the  complete realisation of the Sloan  Digital Sky  Survey (SDSS).   The  four releases  of SDSS  have unveiled   132   new   CVs   ({Szkody   et   al.    \\cite{Szkody2002}, \\cite{Szkody2003}, \\cite{Szkody2004}, \\cite{Szkody2005})}.  Among them is  \\object{SDSS  J123813.73-033933.0},  which  was identified  as  an r=17.82  magnitude CV  (u=17.89, g=17.78,  i=17.97, z=18.07)  with the extremely short orbital period of 76 min determined from seven spectra obtained  by   Szkody  et  al.    (\\cite{Szkody2003}).   The  galactic coordinates of the object are $l=296.51$ and $b=+59.05$, which implies a  galactic extinction  factor of  only E(B-V)=0.03  (Schlegel  et al. \\cite{Schlegel}).    The   proper    motion   was   measured   to   be $p.m.=0\\farcs143$/year as presented in  the USNO B1.0 catalogue (Monet et al.  \\cite{Monet}).  The  optical spectrum of  this system  shows a blue  continuum with  broad absorption  features  around double-peaked Balmer emission  lines.  The spectral appearance  resembles the small, but intriguing, group of so-called \\object{WZ\\,Sge} objects remarkable for  their  large amplitude  outbursts,  long  recurrence cycles,  and peculiar  outburst lightcurves.  {They are  concentrated close  to the lower limit  of the  CV period distribution  and are believed  to have bounced back from the minimum-period limit; i.e. dwarf novae that have periods  that  are  lengthening  after  evolving  through  the  period minimum.}  These  are supposed to  be those very missing  systems with short  periods  that, according  to  theoretical calculations,  should produce a spike  in the number of CVs at  the period minimum. However, the  discovery  and study  of  them  is  hindered by  their  intrinsic faintness,  due  to  the   low-mass  transfer  rates,  and  infrequent outbursts.  The  subject   of  this   paper  is  a   time-resolved  simultaneous spectroscopic   and   photometric  study   of   the  CV   \\object{SDSS J123813.73-033933.0} (hereafter abbreviated as \\object{SDSS1238}).  In Sect.\\ref{Obs} we  describe our  observations and the  data reduction. The data  analysis and the  results are presented  in Sect.\\ref{DatAn} while a discussion  and a summary are given  in Sects.\\ref{Discus} and \\ref{Summ}, respectively.  \\begin{table*}[] \\begin{center} \\caption{Log of observations of SDSS J123813.73-033933.0.} \\begin{tabular}{llllccc} \\hline\\hline Date (2004y.)       & HJD Start+ & Telescope& Instrument/Grating  &Range/Band & Exp.Time/Num. of Integrations& Duration\\\\ Spectroscopy &  2453000       &   &        &     &     &        \\\\ 14 Apr & 109.868 &2.1m &B\\&Ch$^1$ 1200l/mm  & 6000-7100\\AA    &  600s$\\times$11   & 1.92h  \\\\ 15 Apr & 110.650 &2.1m &B\\&Ch\\ \\ \\ 1200l/mm  & 6000-7100\\AA    &  600s$\\times$22   & 3.84h  \\\\ 16 Apr & 111.724 &2.1m &B\\&Ch 400l/mm  &  4200-7300\\AA   & 600s$\\times$15    & 2.93h  \\\\ 17 Apr & 112.658 &2.1m &B\\&Ch 400l/mm   & 5200-8200\\AA    &  600s$\\times$20  & 3.27h  \\\\ 17 Apr & 112.852 &2.1m& B\\&Ch 400l/mm   & 3700-6800\\AA   &   900s$\\times$11  & 4.65h  \\\\ 19 May & 145.710 &2.1m &B\\&Ch 400l/mm   &  4200-7300\\AA    &   600s$\\times$23  & 4.32h  \\\\ Photometry &  &   &    & &                                               &                \\\\ 15 Apr  & 110.737 &1.5m & RUCA$^2$  & R    & 180s$\\times$101                     & 5.76h  \\\\ 16 Apr & 111.751  &1.5m & RUCA & R    &  180s$\\times$119                    & 5.18h  \\\\ 17 Apr & 112.653  &1.5m & RUCA & V    & 120s$\\times$129                        & 7.22h  \\\\ 16 May & 142.675  &1.5m & RUCA & V    & 120s$\\times$133                     & 5.38h  \\\\ 17 May & 143.663  &1.5m & RUCA & V    & 120s$\\times$128                     & 5.51h  \\\\ 18 May & 144.646  &1.5m & RUCA & V    & 120s$\\times$143                     & 5.80h  \\\\ 19 May & 145.652  &1.5m & RUCA & V    &120s$\\times$135                      & 5.64h  \\\\ \\hline \\end{tabular} \\label{tab1} \\end{center} \\begin{tabular}{l} $^1$ B\\&Ch - Boller \\& Chivens spectrograph (http://haro.astrospp.unam.mx/Instruments/bchivens/bchivens.htm) \\\\ $^2$ RUCA - CCD photometer (http://haro.astrospp.unam.mx/Instruments/laruca/laruca\\_intro.htm) \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{Discus}    \\object{SDSS1238}  is a  remarkable CV  on the  lower edge  of  the CV period distribution,  which exhibits a number  of interesting features that define the small class of CVs known as WZ Sge stars.  In addition to  these  features, it  undergoes  frequent  and cyclical  brightness changes, defined  here as LTV, that  sets it apart  from other similar CVs or any other CV as a matter of fact.  The observed period minimum  for CVs with hydrogen-rich secondaries is about 77min  (Kolb \\& Baraffe \\cite{Kolb}).  According  to our current understanding the number of CVs should peak close to the 80 min period limit,  as they evolve  to the  boundary where  they bounce  back (see Patterson \\cite{Patterson98} and references therein).  If this is true then only a tiny part  of that presumed stockpile is recovered.  There are only  about 50  systems known  above the period  minimum to  up to $\\sim$90 min.  Part of the  sample in this period range are \\object{SU Uma} (together with \\object{ER Uma})  objects (about 20 of them), some are magnetic  CVs (10 \\object{AM  Her} and 6 \\object{DQ  Her} objects) and  few are firmly  classified as  WZ Sge  stars based  on infrequent outbursts with distinct high  amplitude and echo re-brightening during fading from  the outburst (Patterson et  al. \\cite{Patterson02}).  The \\object{WZ  Sge} objects  are  probably the  only  ones identified  as evolved/bounced  systems.    Several  other  stars   are  proposed  as candidates  to the  \\object{WZ Sge}  class based  on  their quiescence properties.   \\object{SDSS1238} bears  certain resemblance  to  WZ Sge stars, particularly because of absorption features in the spectrum and the   double  humped   light  curve.    Besides   these  observational characteristics,  the relatively  low temperature  deduced for the  white dwarf (see Urban  et al. \\cite{Urban}) indicate  that \\object{SDSS1238} may belong to this enigmatic group of objects.  While the reason  for the double-humped light curve  is not understood well in these systems, it appears  to be more or less widespread among certain short-period  CVs. True,  in the majority  of them  it appears only during  the outburst.   In short period  CVs, the  extremely late type  secondary star  is  so  dim that  it  usually passed  undetected whether      spectroscopically       or      photometrically.       If SDSS013701.6-091234.9  may  be  anomalous  in that  sense  (Szkody  et al.  \\cite{Szkody2003},   Pretorius  et  al.    \\cite{Pretorius}),  in \\object{SDSS1238} the secondary  is not detected spectroscopically and the  double  hump  light curve  (detected  in  V  band with  the  same amplitude as  in the R) can by  no means be caused  by the ellipsoidal shape  of the  secondary.  As  already mentioned,  in the  majority of other  systems  the  double-humped  light  curve  is  observed  during outbursts.  According to  Patterson et  al.   (\\cite{Patterson02}) the double-hump wave  appears in WZ Sge  within 1 day  of outburst maximum and is then  replaced by a common superhump that  develops in a normal manner.  The  physical processes  producing double humps  in outbursts are not clear either.  In Patterson et al.  (\\cite{Patterson02}) there is  a discussion  of possible  causes and  references to  the proposed models without firm conclusions. Authors favour the model in which the development of a  strong two spiral-arm structure at  the beginning of an outburst, when  the 2 : 1 eccentric  resonance is reached, produces the desired waveform in  the lightcurve (Osaki \\& Meyer \\cite{Osaki}). Within  the evidence  supporting  this  model is  the  detection of  a two-armed spiral  in the Doppler  tomograms of \\object{WZ  Sge} during the first few days of outburst (Steeghs et al. \\cite{Steeghs}; Baba et al. \\cite{Baba}).   Another possible explanation  of the double-humped light  curve  (Patterson  et  al.   \\cite{Patterson02})  is  a  sudden increase in  transfered matter that creates  a hot spot  at the disc's outer edge.  A hot spot is  a common feature  in CVs and in  many high inclination systems  it produces a  distinctive shoulder in  the light curve.  The maximum in the optical light curve origining in such a hot spot is  around $\\phi  = 0.85$, when  the disc  is seen face  on.  The weaker maximum at opposite $\\phi  \\sim 0.3$ may be naturally explained by viewing the  hot spot, partially obstructed by  the accretion disc, from the opposite  side (Silber et al.  \\cite{Silber}).   If we assume that in low $\\dot{M}$ systems in quiescence the disc is optically thin and tiny, then  the hot spot can be visible through  it and create the possibility of a double-humped light curve with maxima at phases $\\phi \\sim 0.3$ and $\\phi \\sim 0.8$.  Incidentally   we  detect   both  phenomena:   (i)   an  unhomogeneous spiral-arm-like structure is  detected at the bottom of  LTV, and (ii) the  accretion  stream/disc  impact  hot  spot  }  dominating  Doppler tomogram during the bright part of  LTV.  As in the case of \\object{WZ Sge} (Lasota  et al. \\cite{Lasota}; Hameury  et al.  \\cite{Hameury}), the  hot spot model  suggests a  sudden burst  of mass  transfer.  The authors argue that the irradiation of the secondary provides the basis for  the  mass transfer  rate  change.   Of  course both  models  were developed  to explain  the observed  features of  \\object{WZ  Sge} and similar  in outbursts.   However the  phenomenon of  double  humps was observed in several occasions in the  quiescent state in WZ Sge and AL Com  (Patterson et  al.   \\cite{Patterson96}) and  a  few other  short period WZ Sge-like systems IY  UMa (Rolfe et al. \\cite{Rolfe}), WX Cet (Rogoziecky      \\&      Schwarzenberg-Czerny      \\cite{Rogoziecky}), \\object{SDSS013701.06-091234.9}  (Szkody  et  al.   \\cite{Szkody2003}, Pretorius  et al.  \\cite{Pretorius}), \\object{BW  Scl/RX J2353.0-3852} (Augusteijn      \\&     Wisotzki      \\cite{Augusteijn};     Abbott et al. \\cite{Abbott}).  In quiescence  there is no  chance that  the disc extends  as far  as the  2:1 resonance  radius, hence the  bright spot  is a better  explanation for  the double  humps.  Besides,  in most  of the other systems, the second hump  is often smaller in amplitude. This is good evidence  of a bright spot that is  weakened as it is  observed on the far side of the disc  and through it.  It is convincingly demonstrated by  Rolfe et  al. (\\cite{Rolfe})  using high  S/N spectra  and Doppler tomography  that the double  hump lightcurve  originates from  the hot spot.   Our data  on \\object{SDSS1238} were obtained certainly  outside of the outburst regime.   Therefore we simply  assume, by analogy  with other similar  systems, that  the  double-humped light  curve  profile is  a result of  increasing hot spot  brightness during the bright  phase of LTV.  Then,  naturally, the increased  brightness of the spot  and the system  as  a  whole can  only  be  explained  in terms  of  increased $\\dot{M}$. Why  the mass transfer  rate and hence the  brightness vary quasi-periodically  on  such  short  time  scales (7-12  hours)  is  a separate and  very interesting  question. One possible  explanation is the  irradiation   of  the  secondary   as  favoured  by   Hameury  et al. (\\cite{Hameury}).  Although this  idea was originally suggested to explain  the triggering of  the superoutburst  in the  \\object{WZ Sge} systems, it  is quite  possible that the  very late type  secondary in \\object{SDSS1238} barely fills its Roche lobe and maintains a delicate balance.  Then,  the slightest irradiation  from the bright,  hot spot will expand  the secondary  to fill its  Roche lobe  thereby expelling increased amounts  of matter toward  the primary.  The  secondary will then shrink, decreasing the mass  transfer only to be irradiated again and  to undergo  another cycle  once the  matter reaches  the  edge of accretion disc and re-brightens it.  In no other system than \\object{SDSS1238} is  cyclical brightening or LTV detected.  According to our proposed scenario this can be due to the  particular mass ratio established  in \\object{SDSS1238}. For example, in  \\object{WZ Sge} the estimated WD  mass is substantially higher, which will lead to a much smaller secondary Roche lobe."
    },
    "0512/astro-ph0512377_arXiv.txt": {
        "abstract": "The apparent shapes of spiral galaxies in the 2-Micron All Sky Survey Large Galaxy Atlas are used to constrain the intrinsic shapes of their disks. When the distribution of apparent axis ratios is estimated using a nonparametric kernel method, the shape distribution is inconsistent with axisymmetry at the 90\\% confidence level in the $B$ band and at the 99\\% confidence level in the $K_s$ band. If spirals are subdivided by Hubble type, the late-type spirals (Sc and later) are consistent with axisymmetry, while the earlier spirals are strongly inconsistent with axisymmetry. The distribution of disk ellipticity can be fitted adequately with either a Gaussian or a lognormal distribution. The best fits for the late spirals imply a median ellipticity of $\\epsilon \\approx 0.07$ in the $B$ band and $\\epsilon \\approx 0.02$ in the $K_s$ band. For the earlier spirals, the best fits imply a median ellipticity of $\\epsilon \\approx 0.18$ in the $B$ band and $\\epsilon \\approx 0.30$ in the $K_s$ band. The observed scatter in the Tully-Fisher relation, for both late and early spirals, is consistent with the disk ellipticity measured in the $B$ band. This indicates that excluding spirals of Hubble type earlier than Sc will minimize the intrinsic scatter in the Tully-Fisher relation used as a distance indicator. ",
        "introduction": "\\label{intro}  The intrinsic shapes of disks in spiral galaxies, although difficult to determine, are useful to know. In particular, the ellipticity of a disk in its outer region reflects the nonaxisymmetry of the potential in which it exists, and thus provides information about the shape of dark matter halos. Non-circular disks produce scatter in the Tully-Fisher relation between rotation velocity and absolute magnitude \\citep{tu77,fr92}; therefore, knowledge of the disk ellipticity helps in understanding the origin of scatter in the Tully-Fisher relation. In particular, the usefulness of the Tully-Fisher relation as a distance indicator would be enhanced for a population of disk galaxies with small intrinsic ellipticity. If the shape of a disk is approximated as an ellipsoid with principal axes of length $a \\geq b \\geq c$, the disk shape can be described by two parameters: the dimensionless disk thickness, $\\gamma \\equiv c/a$, and the disk ellipticity, $\\epsilon \\equiv 1 - b/a$. The simplest approximation, that disks are infinitesimally thin and perfectly circular ($\\gamma = 0$, $\\epsilon = 0$), is occasionally useful; for many purposes, however, it is useful to have a more accurate estimate of the shapes of disks.  The thickness of disks can be determined from catalogs of edge-on spiral galaxies, such as the Flat Galaxy Catalogue \\citep{ka93} and the Revised Flat Galaxy Catalogue \\citep{ka99}. Some superthin galaxies have axis ratios as small as $\\gamma \\approx 0.05$ \\citep{go81}. The study of \\citet{gu92}, using images of edge-on spirals from blue Palomar Sky Survey plates, found that $\\log q$ had a mean value of $-0.95$, where $q$ is the apparent axis ratio. The mean disk thickness depends on the spirals' morphological type \\citep{bo83}; \\citet{gu92} found that the mean value of $\\log q$ ranges from $-0.7$ for edge-on Sa galaxies to $-1.1$ for edge-on Sd galaxies. The disk thickness of an individual spiral galaxy also depends on the wavelength at which it is observed. Edge-on spirals have strong vertical color gradients, becoming much redder with increasing height above the midplane \\citep{da02}. This gradient reflects the transition from the thin disk, with its young blue stellar population, to the thick disk, with its older redder population. As a consequence, $q$ of an edge-on galaxy becomes larger at longer wavelengths. For example, the galaxies in the \\citet{da00} sample were chosen to have $q \\leq 0.125$ at blue wavelengths; in deep $R$-band images, the typical axis ratio increases to $q \\approx 0.23$ \\citep{da02}. Similarly, \\citet{mi03} found that edge-on spirals that have $q \\approx 0.17$ in blue images have $q \\approx 0.35$ in two-micron images.  Estimates of the distribution of disk ellipticity $\\epsilon$ can be made from the distribution of apparent axis ratios $q$ for a large population of randomly oriented spiral galaxies. The signature of non-zero ellipticity, in this case, is a scarcity of nearly circular ($q \\approx 1$) spiral galaxies. \\citet{bi81}, using spiral galaxies from the Second Reference Catalogue of Bright Galaxies \\citep{de76}, concluded that late-type spirals were better fitted by slightly elliptical disks ($\\epsilon \\approx 0.1$) than by perfectly circular disks. More recent studies \\citep{gr85, la92, fa93, al02, ry04} have confirmed the ellipticity of disks, and have permitted estimates of the distribution function for intrinsic ellipticity. The distribution of $\\epsilon$ is commonly fitted with a half-Gaussian peaking at $\\epsilon = 0$. Values of $\\sigma_\\epsilon$, in this case, range from $\\sigma_\\epsilon = 0.12$ \\citep{la92}, through $\\sigma_\\epsilon = 0.13$ \\citep{fa93}, to $\\sigma_\\epsilon = 0.21$ \\citep{ry04}.  Kinematic studies can also provide an estimate of the ellipticity of the gravitational potential in the disk plane. In an elliptical potential with small ellipticity $\\epsilon_\\Phi$, the closed stellar orbits will themselves be ellipses \\citep{bi78}. When the resulting elliptical stellar disk is seen in projection, the isophotal principal axes will be misaligned with the kinematic principal axes. Because of the misalignment, there will be a velocity gradient along the isophotal minor axis proportional to $\\epsilon_\\Phi \\sin 2 \\phi$, where $\\phi$ is the azimuthal viewing angle relative to the long axis of the potential \\citep{fr92}. Measurement of the kinematic misalignment for a sample of 9 disk galaxies yielded an average ellipticity of $\\epsilon_\\Phi \\approx 0.044$ \\citep{sc97,sc98}. A similar analysis of gas kinematics in 7 disk galaxies by \\citet{be99} found an average ellipticity of $\\epsilon \\approx 0.059$ for the gas orbits. \\citet{an02}, using the method of \\citet{an01}, combined two-dimensional kinematic and photometric data for a sample of 28 nearly face-on disk galaxies; they found an average halo ellipticity of $\\epsilon_\\Phi \\approx 0.054$ and an average disk ellipticity of $\\epsilon \\approx 0.076$ for their sample. Taking into account the selection criteria of \\citet{an02}, the implied median disk ellipticity is $\\epsilon \\approx 0.10$ \\citep{ry04}.  Much of what we know about the ellipticity of disks thus comes from looking at galaxies with low inclination; much of what we know about disk thickness comes from looking at galaxies with high inclination. In principle, both ellipticity and thickness can be determined by looking at a sample of disk galaxies with random inclinations \\citep{fa93}. However, at visible wavelengths, thanks to the effects of dust, it's difficult to create a sample in which the disk inclinations are completely random. Because disk galaxies are not fully transparent, a flux-limited sample will be deficient in nearly edge-on galaxies \\citep{hu92}. Because disk galaxies are not fully opaque, an angular diameter-limited sample will have an excess of nearly edge-on galaxies \\citep{hu92,ma03}.  To minimize the effects of dust, I will be examining the apparent shapes of disk galaxies in the Two Micron All-Sky Survey (2MASS). In \\S\\ref{sec-sample}, I extract a randomly inclined sample of spiral galaxies from the 2MASS Large Galaxy Atlas. In \\S\\ref{sec-axisym}, I make both nonparametric and parametric estimates of $f(\\gamma)$, the distribution of disk thickness, assuming the disks are axisymmetric. Separate estimates are made for early-type spirals and for late-type spirals. In \\S\\ref{sec-nonaxi}, I make nonparametric estimates of $f(\\epsilon)$, the distribution of disk ellipticity for nonaxisymmetric disks. Both Gaussian and lognormal functions are found to give satisfactory fits to $f(\\epsilon)$. In \\S\\ref{sec-disc}, I examine the influence of the deduced disk ellipticity on the scatter of the Tully-Fisher relation, and consider the possible origins of disk ellipticity for early-type and late-type spirals. ",
        "conclusions": "\\label{sec-disc}  The 2MASS Large Galaxy Atlas provides a sample of nearby spiral galaxies with large angular size ($R_{3\\sigma} > 60 \\arcsec$ at $\\lambda \\sim 2$ microns). In this paper, I have used this relatively small sample of very well-resolved galaxies to estimate the distribution of apparent and intrinsic shapes of spiral galaxies. An intriguing result of this study is the difference in disk ellipticity $\\epsilon$ between disk galaxies of early and late Hubble type.  If the entire population of 2MASS LGA spirals is considered, summed over all Hubble types, the distribution of ellipticities in the $B$ band is consistent with earlier photometric studies using larger sample sizes \\citep{la92,fa93,ry04}. The $K_s$ band ellipticity is greater, but the shape of the $3\\sigma$ isophote in the $K_s$ band can be influenced by light from triaxial bulges. The distribution $N_\\epsilon (\\epsilon)$ for shapes in the $B$ band can be modeled as a Gaussian peaking at $\\mu_\\epsilon \\approx 0$ and width $\\sigma_\\epsilon \\approx 0.16$ or as a lognormal distribution with $\\mu_\\eta \\approx -2.4$ and $\\sigma_\\eta \\approx 1.0$; both these distributions imply a mean ellipticity $\\epsilon \\approx 0.1$. When the late spirals, consisting mainly of Hubble type Sc, are examined, they are found to be perfectly consistent with axisymmetry in both the $B$ and the $K_s$ band. The early spirals, consisting mainly of types Sb and Sbc, show a high degree of ellipticity, even in the $B$ band, in which the mean ellipticity is $\\epsilon \\approx 0.18$.  It has been noted by \\citet{fr92} that a disk ellipticity of $\\epsilon > 0.1$ cannot be reconciled with the observed small scatter in the Tully-Fisher relation \\citep{tu77} if the disk ellipticity is assumed to trace the potential ellipticity. Suppose that stars and gas are on closed orbits in a logarithmic potential with rotation speed $v_c$. If the potential has a small ellipticity $\\epsilon_\\phi \\ll 1$ in the orbital plane, then the integrated line profile from the disk will have a width $W = 2 v_c (1 - \\epsilon_\\phi \\cos 2 \\phi ) \\sin \\theta$ when viewed from a position angle $(\\theta,\\phi)$ \\citep{fr92}. The difference in line width from that produced in a purely circular disk will produce a scatter in the observed Tully-Fisher relation. If all disks have $\\epsilon_\\phi = 0.1$, the expected scatter is 0.3 mag, even if the inclination has been determined accurately from kinematic information \\citep{fr92}. In most studies of the Tully-Fisher relation, the inclination is determined photometrically, from the apparent axis ratio of the disk, assuming (perhaps erroneously) that the disk is axisymmetric. If inclinations are determined in this way, the scatter in the Tully-Fisher relation will be even greater.  If galaxies had infinitesimally thin circular disks, the relation between apparent axis ratio $q$ and inclination $\\theta$ would be $q = \\cos \\theta$. If the thin disks actually have ellipticity $\\epsilon$ and are viewed at high inclination ($\\sin^2 \\theta \\gg 2 \\epsilon$), the apparent axis ratio will be $q \\approx \\cos \\theta ( 1 - \\epsilon \\cos 2 \\phi )$. Thus, if the inclination is estimated by the relation $\\theta_{\\rm phot} \\equiv \\cos^{-1} q$, a fractional error $\\propto \\epsilon \\cos 2 \\phi$ will be introduced into $\\cos \\theta_{\\rm phot}$ and an error $\\propto \\epsilon \\cot^2 \\theta \\cos 2 \\phi$ will be introduced into $\\sin \\theta_{\\rm phot}$. Since the observed velocity width $W$ must be divided by $\\sin \\theta$ to find the rotation speed $v_c$, the error produced by using $\\sin \\theta_{\\rm phot}$ instead of $\\sin \\theta$ will be negligible only when $\\cot^2 \\theta \\ll 1$. Thus, most studies of the Tully-Fisher relation use galaxies of high inclination ($\\theta_{\\rm phot} \\gtrsim 45^\\circ$). The Tully-Fisher relation for nearly face-on galaxies (see, for instance, \\citet{an03}) can only be determined if the inclinations are determined kinematically.  In Figures~\\ref{fig:gaus_non} and \\ref{fig:lognorm_non}, the dashed lines indicate the expected amount of scatter in the Tully-Fisher relation if the potential ellipticity is given by either a Gaussian distribution (Figure~\\ref{fig:gaus_non}) or by a lognormal distribution (Figure~\\ref{fig:lognorm_non}). Starting at the lower left corner of each panel, the contours are drawn at the levels 0.25 mag, 0.5 mag, 0.75 mag, and 1.0 mag. The inclinations are assumed to be estimated from the apparent axis ratio. Since the disks are not always infinitesimally thin, the inclination $\\theta_{\\rm phot}$ is estimated using the relation \\begin{equation} \\cos^2 \\theta_{\\rm phot} = {q^2 - \\mu_\\gamma^2 \\over 1 - \\mu_\\gamma^2} \\ . \\end{equation} Ultrathin galaxies with $q < \\mu_\\gamma$ are assumed to have $\\cos \\theta = 1$. All galaxies with $\\theta_{\\rm phot} < 45^\\circ$ are discarded.  If all spirals are considered together, the best fitting ellipticity distributions -- either Gaussian or lognormal -- imply approximately 1 mag of scatter if the potential ellipticity equals the disk ellipticity: this can be seen in the upper panels of Figures~\\ref{fig:gaus_non} and \\ref{fig:lognorm_non}. The best fitting ellipticity for the late spirals, as seen in the middle panels of Figures~\\ref{fig:gaus_non} and \\ref{fig:lognorm_non}, implies only 0.3 mag of scatter. Finally, the best fitting ellipticity for the early spirals, seen in the lower panels of Figures~\\ref{fig:gaus_non} and \\ref{fig:lognorm_non}, leads to 1.4 mag of scatter in the Tully-Fisher relation. By comparing the predicted scatter from disk ellipticity to the actual scatter in the Tully-Fisher relation for these galaxies, I can place constraints on how much of the scatter can be attributed to the disk ellipticity.  To create a Tully-Fisher plot for the spirals in the 2MASS LGA sample, I took the relevant astrophysical parameters from the HyperLeda database.\\footnote{ See http://leda.univ-lyon1.fr/} As a measure of galaxy luminosity, I used the $I$ band absolute magnitude; multiband studies of the Tully-Fisher relation indicate that the observed scatter is minimized in or near the near-infrared.  The $I$ band absolute magnitude is computed from the apparent magnitude $m_I$ (corrected for galactic extinction and internal extinction) and the distance modulus $m-M$. The distance modulus was computed from the radial velocity $v_{\\rm Vir}$ corrected for infall of the Local Group toward Virgo, assuming a Hubble constant $H_0 = 70 {\\rm\\,km}{\\rm\\,s}^{-1} {\\rm\\,Mpc}^{-1}$; for galaxies with $v_{\\rm Vir} < 500 {\\rm\\,km} {\\rm\\,s}^{-1}$, the distance modulus was taken from the literature, with preference given to Cepheid distances. As a measure of rotation speed, I used the maximum rotation velocity $v_m$ determined from the 21 cm line of neutral hydrogen. In computing $v_m$, the inclination used is the photometric estimate $\\theta_{\\rm phot}$, using the axis ratio of the $25 {\\rm\\,mag}{\\rm\\,arcsec}^{-2}$ isophote as the apparent axis ratio $q$.  Figure~\\ref{fig:tf} shows $M_I$ as a function of $\\log v_m$ for the 2MASS LGA spirals with $R_{20}^o > 60 \\arcsec$, assuming axisymmetry; galaxies with $\\theta_{\\rm phot} < 45^\\circ$ are omitted, as are galaxies without $I$ band photometry in HyperLeda. A total of $n = 128$ galaxies meet all these criteria. The solid line in Figure~\\ref{fig:tf} is the best fitting Tully-Fisher relation, found using an unweighted inverse fit \\citep{sc80,tu00,ka02}. The best Tully-Fisher relation for the sample of $n = 128$ galaxies is \\begin{equation} M_I ( {\\rm all} ) = (-22.43 \\pm 0.09) - (9.42 \\pm 0.77) \\log_{10} \\left( {v_m \\over 200 {\\rm\\,km}{\\rm\\,s}^{-1} } \\right) \\ , \\end{equation} with a scatter in $M_I$ of $\\sigma_{\\rm rms} = 0.94$ mag. If only the $n = 37$ late spirals are considered, an unweighted inverse fit yields \\begin{equation} M_I ( {\\rm late} ) = (-22.54 \\pm 0.15 ) - ( 9.42 \\pm 0.85 ) \\log_{10} \\left( {v_m \\over 200 {\\rm\\,km} {\\rm\\,s}^{-1} } \\right) \\ , \\end{equation} with a scatter in $M_I$ of $\\sigma_{\\rm rms} = 0.62$ mag. Finally, if only the $n = 91$ early spirals are considered, the fit is \\begin{equation} M_I ( {\\rm early} ) = ( -22.40 \\pm 0.12 ) - ( 9.97 \\pm 1.38 ) \\log_{10} \\left( {v_m \\over 200 {\\rm\\,km} {\\rm\\,s}^{-1} } \\right) \\ , \\end{equation} with a scatter in $M_I$ of $\\sigma_{\\rm rms} = 1.10$ mag. The late and early type spirals have Tully-Fisher relations with statistically indistinguishable slopes; however, the early spirals have a much greater scatter.  The observed scatter, $\\sigma_{\\rm rms}$, is not due entirely to the nonaxisymmetry of disks. In addition to an intrinsic scatter $\\sigma_{\\rm ell}$, due to the ellipticity of the disk and the potential, there is also a contribution $\\sigma_{\\rm err}$ due to errors in observation and interpretation. For instance, the absolute magnitude $M_I$ is subject to errors in flux measurement, errors in determining distance, and errors in extinction correction. The rotation speed $v_m$ is subject to errors in measuring 21 cm line widths and errors in converting line widths to deduced rotation speeds. Let us suppose that the two sources of scatter add in quadrature: $\\sigma_{\\rm rms}^2 = \\sigma_{\\rm ell}^2 + \\sigma_{\\rm err}^2$. If $\\sigma_{\\rm err}$ is the same for both early and late spirals, then \\begin{equation} \\sigma_{\\rm ell} ({\\rm early})^2 - \\sigma_{\\rm ell} ({\\rm late})^2 = \\sigma_{\\rm rms} ({\\rm early})^2 - \\sigma_{\\rm rms} ({\\rm late})^2 = ( 0.91 {\\rm\\,mag})^2 \\ . \\end{equation} If late spirals are perfectly axisymmetric, which is permitted by the data, then $\\sigma_{\\rm ell} ({\\rm late}) = 0 {\\rm\\,mag}$ and $\\sigma_{\\rm ell} ({\\rm early}) = 0.91 {\\rm\\,mag}$. The largest possible values of $\\sigma_{\\rm ell}$ come if I na{\\\"\\i}vely assume that $\\sigma_{\\rm err} = 0$; in this limiting case, $\\sigma_{\\rm ell} ({\\rm late}) = 0.62 {\\rm\\,mag}$ and $\\sigma_{\\rm ell} ({\\rm early}) = 1.10 {\\rm\\,mag}$.  For late spirals, the Tully-Fisher relation thus implies that the intrinsic scatter due to ellipticity lies in the range $\\sigma_{\\rm ell} = 0 \\to 0.62 {\\rm\\,mag}$. This is consistent with the value $\\sigma_{\\rm ell} = 0.3 {\\rm\\,mag}$ derived from the best fitting distributions of disk ellipticity. For early spirals, the Tully-Fisher relation implies an intrinsic scatter due to ellipticity in the range $\\sigma_{\\rm ell} = 0.91 \\to 1.10 {\\rm\\,mag}$. This is smaller than the value $\\sigma_{\\rm ell} = 1.4 {\\rm\\,mag}$ derived from the best fitting distributions of disk ellipticity. However, in the lower panels of Figures~\\ref{fig:gaus_non} and \\ref{fig:lognorm_non}, the band where $\\sigma_{\\rm ell} = 0.91 \\to 1.10$ overlaps the region where $P > 0.1$; thus, the ellipticities derived from the apparent shapes and from the Tully-Fisher relation are not discrepant at a high confidence level.  The difference in disk ellipticity between earlier and later spirals, and the resulting difference in Tully-Fisher scatter, helps to explain the apparent discrepancy between the relatively large disk ellipticity ($\\epsilon \\sim 0.1$) seen in photometric studies \\citep{bi81,gr85,la92,fa93,al02,ry04} and the relatively small intrinsic scatter ($\\sigma \\lesssim 0.3 {\\rm\\,mag}$) seen in some studies of the Tully-Fisher relation. For instance, \\citet{ve01}, in his study of spirals in the Ursa Major cluster, found that his DE (distance estimator) subsample of spirals had a Tully-Fisher relation with 0.26 magnitudes of scatter in the $I$ band; the intrinsic scatter had a most likely value of only 0.06 magnitudes. This small scatter was partly due to the strict selection criteria imposed; galaxies were excluded if they were obviously interacting, had prominent bars, or had irregular outer isophotes. Galaxies of Hubble types earlier than Sb or later than Sd were also excluded. The scatter was further minimized by the fact that the resulting subsample of 16 spirals contained 10 galaxies of type Sc, Scd, and Sd (late spirals, in my classification) and only 6 galaxies of type Sb and Sbc. If the spirals in the 2MASS LGA are characteristic of all spiral galaxies, then scatter in the Tully-Fisher relation can be reduced by excluding spirals of Hubble type earlier than Sc.  In addition to the practical implications for minimizing scatter in the Tully-Fisher relation, the difference in ellipticity between early and late spirals provides an intriguing clue for the formation and evolution of spiral galaxies. Why should the Sb and Sbc galaxies dominating the early spiral sample be more elliptical at $R_{25}$ than the Sc galaxies? One possibility is that the triaxiality of a dark halo's potential affects the Hubble type of the spiral embedded within it. \\citet{ko04} estimate that the majority of spirals with Hubble type of Sb and later contain ``pseudobulges'' rather than classical bulges. Pseudobulges are formed as gas is transported to small radii by nonaxisymmetric structures such as triaxial halos, spiral structure, elliptical disks, and bars. Simulations of disk galaxies in triaxial dark halos \\citep{el98,el01} indicate that pseudobulges grow by secular evolution on a timescale of a few Gyr. The Hubble type of the resulting galaxy depends on both the halo core radius and the potential asymmetry, with greater halo nonaxisymmetry leading to larger bulges.  High resolution n-body simulations of virialized dark halos in the mass range $10^{12} {\\rm\\,M}_\\odot \\lesssim M \\lesssim 10^{14} {\\rm\\,M}_\\odot$ indicate that virialized halos are well described as triaxial shapes \\citep{ji02,al05}. More massive halos are further from spherical, on average. In a lambda CDM simulation, the mean short-to-long axis ratio of a halo with mass $M$ at the present day ($z = 0$) is $\\langle \\gamma \\rangle \\approx 0.54 ( M / M_* )^{-0.05}$, where $M_* = 1.2 \\times 10^{13} {\\rm\\,M}_\\odot$ is the characteristic nonlinear mass scale \\citep{al05}. The mean intermediate-to-long axis ratio also decreases with increasing mass, going from $\\langle \\beta \\rangle \\approx 0.76$ at $M = 0.1 M_*$ to $\\langle \\beta \\rangle \\approx 0.71$ at $M = M_*$. (Dissipation by baryons tends to compress the dark halo into a more nearly spherical shape in its central regions \\citep{ka04}, but will not destroy the trend that lower-mass halos are rounder on average in the plane of the baryonic disk.) At a given halo mass, there exists a fairly wide spread in $\\beta$ around the mean value, with standard deviation $\\sigma_\\beta \\sim 0.1$ for $\\langle \\beta \\rangle \\sim 0.7$ \\citep{al05}. The nearly axisymmetric halos will give rise to nearly circular disks and smaller pseudobulges, such as those seen in Sc galaxies. The more nonaxisymmetric halos will give rise to elliptical disks and larger pseudobulges, such as those seen in Sb and Sbc galaxies. The mean mass of Sb spirals is greater than that of Sc spirals; however, the spread in masses for a given Hubble type is greater than the difference in the mean \\citep{ro94}. This would be expected if the difference between early (Sb and Sbc) spirals and late (Sc) spirals is not due primarily to differences in halo mass, but rather is due to differences in halo nonaxisymmetry."
    },
    "0512/hep-ph0512197_arXiv.txt": {
        "abstract": "We study the stau lifetime in a scenario with the LSP taken to be a neutralino and the NLSP being a stau, based on the minimal supersymmetric Standard Model. The mass difference between the LSP and NLSP, $\\delta m$, must satisfy $\\delta m/m_{\\tilde{\\chi}} \\sim$ a few \\% or less for coannihilation to occur, where $m_{\\tilde{\\chi}}$ is the neutralino mass. We calculate the stau lifetime from the decay modes $\\tilde{\\tau}\\rightarrow \\tilde{\\chi}\\tau$, $\\tilde{\\chi}\\nu_\\tau\\pi$, and $\\tilde{\\chi}\\nu_\\tau\\mu(e)\\nu_{\\mu(e)}$ and discuss its dependence on various parameters. We find that the lifetime is in the range $10^{-22}$--$10^{16}$ sec for $10^{-2} \\le \\delta m \\le 10$ GeV. We also discuss the connection with lepton flavor violation if there is mixing between sleptons. ",
        "introduction": "The existence of non-barionic dark matter is now confirmed and its density has been quantitatively estimated~\\cite{Spergel:2003cb,Bennett:2003bz}. However its identity is still unknown. One of the most prominent candidates is the weakly interacting massive particle (WIMP)~\\cite{Jungman:1995df,Bergstrom:2000pn,Munoz:2003gx,Bertone:2004pz}.  As is well known, the supersymmetric extension of the Standard Model provides a stable exotic particle, the lightest supersymmetric particle (LSP), if R parity is conserved.  Among LSP candidates, the neutralino LSP is the most suitable for non-barionic dark matter since its nature fits that of the WIMP~\\cite{Goldberg:1983nd,Ellis:1983ew}. Neutralinos are a linear combination of the supersymmetric partners of the U(1) and SU(2) gauge bosons (bino and wino) and the Higgs bosons (Higgsino).  They have mass in a range from 100 GeV to several TeV and are electrically neutral. The lightest neutralino is stable if R parity is conserved.  Since the supersymmetric extension of the Standard Model is the most attractive theory, the nature of neutralino dark matter has been studied extensively~\\cite{Ellis:1999mm}. % In many scenarios of the supersymmetric model, the LSP neutralino consists mainly of the bino, the so-called bino-like neutralino. In this case, naive calculations show that the predicted density in the current universe is too high and it is necessary to find a way to reduce it. One mechanism to suppress the density is coannihilation~\\cite{Griest:1990kh}. If the next lightest supersymmetric particle (NLSP) is nearly degenerate in mass with the LSP, the interaction of the LSP with the NLSP is important in calculating the LSP annihilation process. For coannihilation to occur tight degeneracy is necessary, since without coannihilation the LSP decouples from the thermal bath at $T \\sim m/20$~\\cite{Lee:1977ua}, where $m$ is the LSP mass. Therefore the mass difference $\\delta m$ must satisfy $\\delta m/m<$ a few \\%, otherwise the NLSP decouples before coannihilation becomes dominant. Furthermore, if the degeneracy is much tighter, we would observe a line spectrum of photons from pair annihilation of dark matter~\\cite{Hisano:2003ec,Matsumoto:2005ui}, since the annihilation cross section of dark matter would be strongly enhanced due to the threshold correction.  A candidate for the NLSP is the stau or stop in many class of MSSM, and in this paper we study the lifetime of the stau-like slepton having mass degenerate with the LSP neutralino. For the neutralino LSP to be dark matter, very tight degeneracy is required in mass between the NLSP and the LSP neutralino. In particular the heavier the LSP is, the tighter the degeneracy must be~\\cite{Ellis:1999mm}. For such a degeneracy, the NLSP is expected to have a long lifetime due to phase space suppression~\\cite{Profumo:2004qt,Gladyshev:2005mn}.  An alternative scenario for a long-lived scalar particle is the gravitino LSP.  Considerable work has been devoted to the long life NLSP in the context of the gravitino LSP. In this case, due to the small coupling between a superWIMP (including the gravitino) and the NLSP, the lifetime of the NLSP becomes very long~\\cite{Feng:2004mt,Hamaguchi:2004ne,Bi:2004ys}. To determine the most likely candidate for the LSP, we can accumulate and identify~\\cite{Hamaguchi:2004df,Feng:2004yi} the candidate for long-lived NLSPs and compare the nature of the particles including couplings.  An implication on the fundamental feature of quantum mechanics would be brought by a long life stau. As noted in ref.~\\cite{Jittoh:2004bz}, a Small-Q-value S-wave (SQS) decay can exhibit non-exponential decay. A small Q value naively implies a small mass difference. For a gravitino LSP, the decay must occur in a P wave since the gravitino has spin 3/2. In contrast, in our case the decay can be S wave since all daughter particles have spin 0 or spin 1/2. Therefore, to extract fundamental parameters from observations we must take special care in interpreting the results. This study offers the opportunity to examine a fundamental problem of quantum physics in collider physics. Hence it is worthwhile studying the proposition that the stau is the NLSP and the lightest neutralino is the LSP, and that their masses are tightly degenerate.  In Sec.~\\ref{model}, we present the relevant Lagrangian and calculate the decay rate of the stau. In Sec.~\\ref{param}, we study the parameter dependences of the lifetime. We then consider the connection with lepton flavor violation (LFV) in Sec.~\\ref{lfv}. Finally we summarize our results in Sec.~\\ref{summary}. ",
        "conclusions": "\\label{summary} We have studied an MSSM scenario in which the LSP and NLSP are a bino-like neutralino and a stau, respectively. Since the mass difference is, in many cases, assumed to be degenerate, from the requirement of coannihilation, we paid special attention to the very small $\\delta m$ case. We calculated the partial lifetimes for the decay modes shown in Fig.~\\ref{diagrams}.  We have investigated the stau lifetime dependence on $\\delta m$, $\\theta_\\tau$, $\\gamma_\\tau$, and $m_{\\tilde{\\chi}}$, considering cosmological constraints. The lifetime strongly depends on $\\delta m$ and $\\theta_\\tau$, while it is almost independent of $\\gamma_\\tau$ and $m_{\\tilde{\\chi}}$.  The stau lifetime dependence on $\\delta m$ changes as each threshold is crossed. When $\\delta m$ is larger than $m_\\tau$, the lifetime increases in proportion to $(\\delta m)^{-2}$ as $\\delta m$ decreases. In the range $m_\\tau>\\delta m>m_\\pi$ the lifetime obeys the scaling $\\sim (\\delta m)^{-6}$. Below $m_\\pi$, it grows with $(\\delta m)^{-8}$. The $\\delta m$ dependence of the stau lifetime can be largely understood by counting the mass dimension of phase space and the squared amplitude in the massless limit of Standard Model particles. While the massless limit is a good approximation in regions far from the thresholds, the $\\delta m$ dependence near the thresholds are given by Eq.~\\eqref{appro_thresh}.  The lifetime also strongly depends on $\\theta_\\tau$, as shown in Fig.~\\ref{enlargedlifetime}. $\\tilde{\\tau}_R$ contributes to 3- and 4-body decay processes by picking up the $m_\\tau$ term in the intermediate $\\tau$ propagator, while $\\tilde{\\tau}_L$ picks up the $p_\\tau$ term. Since $p_\\tau \\sim \\delta m \\ll m_\\tau$, the contribution for $\\tilde{\\tau}_R$ is much larger and hence there is a strong dependence on $\\theta_\\tau$.  As is seen in Fig.~\\ref{lifetime300}, if $\\delta m$ is smaller than $m_\\tau$,  a stau can be very long-lived. This fact may need to be taken into account in studies at the Large Hadron Collider (LHC) and the International Linear Collider (ILC). The superWIMP case is similar. However, in the gravitino scenario an energetic $\\tau$ is apparently produced, while in our scenario we should observe very low energy $\\pi$, $\\mu$, or $e$. Therefore a clean experiment is required. The ILC would be most suitable for investigating the nature of the NLSP slepton.  Another candidate for the NLSP is the stop. However, it is more complicated to investigate this possibility. A stop must always be accompanied by quark(s); it forms a meson-like fermion with one quark or a baryon-like boson with two quarks. It is almost impossible to calculate the exact mass eigenvalue using QCD. It is beyond the scope of this paper and is left for future work.  We have also discussed lepton flavor violation due to slepton mixing. If there is even a tiny component of a scalar electron or a scalar muon in the NLSP ``stau'', the decay signal of the NLSP will be completely different from the pure stau case. The NLSP slepton undergoes 2-body decay into the accompanying electron or muon. Since it is a 2-body process, it occurs very quickly, $\\sim ~(10^{-20})N_1^{-2}$ sec where $N_1$ represents the portion of the scalar electron or scalar muon, as shown in Eq.~(\\ref{eq:LFVNLSP}).  As this mixing causes charged/neutral lepton flavor violation, it is very important to compare the NLSP slepton decay with other processes such as $\\tau\\rightarrow e(\\mu)\\gamma$. To fully clarify the nature of the NLSP, we need to interpret all the LFV processes together.  SQS decay takes place in a case of small mass difference. It will appear strongly, if decay rate is small. Since long lifetime means small decay rate, long-lived stau will give a good chance to observe a non-exponential decay. In contrast, if there is no LFV in slepton mixing, the pure stau has a very long lifetime and it is then possible to experimentally observe SQS decay~\\cite{Jittoh:2004bz}. This is important to check the fundamental feature of quantum mechanics.  Hence, the small $\\delta m$ case is very interesting and its study is very important."
    },
    "0512/astro-ph0512363_arXiv.txt": {
        "abstract": "We build an accurate database of 5200 HCN and HNC rotation-vibration energy levels, determined from existing laboratory data. 20~000 energy levels in the \\citet{linepaper} linelist are assigned approximate quantum numbers. These assignments, lab determined energy levels and \\citet{linepaper} energy levels are incorporated in to a new energy level list. A new linelist is presented, in which frequencies are computed using the lab determined energy levels where available, and the \\ai\\  energy levels otherwise.  The new linelist is then used to compute new model atmospheres and synthetic spectra for the carbon star WZ~Cas. This results in better fit to the spectrum of WZ~Cas in which the absorption feature at 3.56 \\mum\\  is reproduced to a higher degree of accuracy than has previously been possible. We improve the reproduction of HCN absorption features by reducing the abundance of Si to [Si/H] = --0.5 dex, however, the strengths of the $\\Delta v=2$ CS band heads are over-predicted. ",
        "introduction": "\\label{sec:Intro}  Cool AGB carbon stars often show a strong absorption feature at 3 microns. This feature was discovered \\citep{Johnson} and later identified as the C-H stretch mode of HCN and/or C$_2$H$_2$ \\citep*{Fay,Ridgeway}. Subsequent to the discovery of HCN in observed carbon star spectra, it was found that the line opacity of HCN has a strong effect upon the structure of model atmospheres \\citep{Eriksson,Jorgensen1}. \\citet*{linepaper}, referred to as HPT, have made publicly available an extensive and accurate {\\em ab initio} HCN and HNC linelist. This linelist has been used to compute model atmospheres and synthetic spectra that reproduce the HCN features in the observed spectra of the carbon stars TX~Psc and WZ~Cas and has also allowed the identification off HNC absorption at 2.9 \\mum\\  \\citep{paper1}. However, the HPT linelist has inherent inaccuracies which result in the line frequencies deviating from laboratory measurements by 3 \\cm\\  or more. This error is sufficiently large to be noticeable at the resolving power of the ISO SWS spectrometer. In particular \\citet{paper1} fail to accurately reproduce the observed HCN Q-branch absorption feature at 3.56-3.62 \\mum.  In this work we assign approximate quantum numbers to 20~000 of the HPT \\ai\\  energy levels. The quality of the HPT linelist is improved by substituting energy levels derived from laboratory measurements of line frequencies, for \\ai\\  energy levels. This improves the frequencies of the lines for which experimental data is available, to spectroscopic precision (0.1 \\cm) and allows the calculation of more accurate C-star synthetic spectra. These new synthetic spectra significantly improve the reproduction of the HCN Q-branch absorption feature at 3.56-3.62 \\mum. ",
        "conclusions": "\\label{sec:conc}  We have assigned approximate quantum numbers to 20~000 of the {\\em ab initio} energy levels found in the HCN/HNC linelist of \\citet{linepaper}. Existing laboratory measurements of HCN and HNC line frequencies are used to accurately determine 5200 energy levels. Both the energy level assignments and the \\lee\\  energy levels have been included along with the \\ai\\  energy levels of \\citet{linepaper} in a new energy level list. A combination of the \\lee\\  and \\ai\\  energy levels are used in conjunction with the \\citet{linepaper} Einstein A coefficients to produce an improved linelist. We make the new energy level file which includes assignments, laboratory determined energy levels and the resulting improved HCN and HNC linelist publicly available.  The new linelist has been incorporated into our computations of C-rich model stellar atmospheres and synthetic spectra. The new synthetic spectra are a clear improvement over our earlier synthetic spectra \\citep{paper1}, and show better agreement with the observed spectrum of WZ~Cas. The improvement is clearest in the range 3.56-3.62 \\mum. We further improve the fit to HCN features by reducing the Si abundance to [Si/H]=--0.5 dex, however the CS absorption feature at 3.8-4.0 microns is still poorly reproduced."
    },
    "0512/astro-ph0512649_arXiv.txt": {
        "abstract": "The use of Gamma Ray Bursts as ``standard candles'' has been made possible by the recent discovery of a very tight correlation between their rest frame intrinsic properties. This correlation relates the GRB prompt emission peak spectral energy $E_{\\rm peak}$ to the energy $E_{\\gamma}$ corrected for the collimation angle $\\theta_{\\rm jet}$ of these sources. The possibility to use GRBs to constrain the cosmological parameters and to study the nature of Dark Energy are very promising. ",
        "introduction": "The extremely large luminosity of Gamma Ray Bursts (GRBs) makes them detectable, in principle, out to very large redshifts $z<20$ (e.g. \\cite{Lamb2000}).  The present redshift distribution for $\\sim$60 GRBs, which extends out to $z=$ 6.29 for GRB 050904 (\\cite{Antonelli2005}) would make GRBs exquisite potential tools for observational cosmology.  They can have a profound impact on: (i) the study the epoch of re-ionization; (ii) the characterization of the properties of the cosmic intervening medium; (iii) the study of the cosmic star formation history back to unprecedented epochs; (iv) the description of the geometry of the Universe and (v) the investigation of the nature of Dark Energy (DE).  However, the last two points require a class of ``standard candles'' whose spread in the Hubble diagram is comparable (and even smaller) than the precision on the measure of their luminosity distance. At first glance GRBs are everything but standard candles: their intrinsic isotropic emitted energies span more than 4 orders of magnitudes, and even the collimation corrected energy span about two orders of magnitudes.  This has prevented, until recently, their application as cosmological tools (\\cite{Bloom2001}; \\cite{Frail2001}; \\cite{Schaefer2003})  However, the discovery of a very tight (so called ``Ghirlanda'') correlation, with a scatter less than 0.1 dex, between the rest frame spectral peak energy $E_{\\rm peak}$ and the collimation corrected energy $E_{\\gamma}$ (\\cite{Ghirlanda2004}) allowed a very accurate measurement of the true GRB energetics and made them usable as ``standard candles'' to constrain the cosmological parameters $\\Omega_{M},\\Omega_{\\Lambda}$ and to study the DE equation of state (\\cite{Ghirlanda2004a,Firmani2005}). \\begin{center} \\begin{figure} \\includegraphics[height=0.6\\textheight]{GGL04_updated.ps} \\caption{ Rest frame peak energy $E_{\\rm peak}$ versus isotropic (open symbols) and collimation corrected (filled symbols) energy.  The black open circles are the 18 GRBs (\\cite{Nava2005}) with measured $z$, $E_{\\rm peak}$ and $t_{\\rm jet}$ for which the collimation corrected energy could be computed (red filled circles).  Upper/lower limits on either one of the variables are shown by the blue filled triangles. The best fit powerlaw to the red filled circles (Eq. 1) is represented by the solid (blue) line and its uncertainty by the thin yellow shaded region.  The large (light orange) shaded region represents the $3\\sigma$ gaussian scatter of the data points around the correlation. Also shown is the Amati correlation obtained by fitting the open black data points and the open grey triangles (which are not upper/lower limits) either accounting for the errors on both the coordinates (long dashed line - slope = 0.54 and $\\chi^2_{\\rm r}=8.4$ for 26 dof) and by a linear regression (dotted line - slope = 0.4).  Note the two XRF (030723 and 020903) and the two GRBs 980425/SN1998bw and 031203/SN2003lw which are outliers with respect to both the Ghirlanda and the Amati correlations.  } \\label{fig_gg1} \\end{figure} \\end{center} There are some issues to be discussed about the use of the Ghirlanda correlation for cosmology and the existence of the $E_{\\rm peak}$-$E_{\\gamma,\\rm iso}$ correlation (so called ``Amati'' ref \\cite{Amati2002}).  The cosmological use of the $E_{\\rm peak}$-$E_{\\gamma}$ correlation suffers from the so called ``circularity'' problem (\\cite{Ghirlanda2004}, see also \\cite{Ghisellini2005}) due to the fact that the correlation is not calibrated with (for instance) low redshift GRBs and, therefore, its slope and normalization are cosmology dependent.  The latter problem could be solved with either a number of low redshift GRBs or through a convincing theoretical interpretation (see \\cite{Eichler2004,Levinson2005}, \\cite{Yamazaki2004} and \\cite{Rees2005}) for possible interpretations) fixing at least its slope.  In the meanwhile different approaches have been adopted (\\cite{Ghirlanda2004a,Firmani2005}) to circumvent this problem.    The Amati correlation was derived with a very limited sample of GRBs (originally only 9 \\cite{Amati2002}, then 24 \\cite{Ghirlanda2004}) and it is possible that it is affected by some selection effect connected to the need to have a measured spectroscopic redshift (which may select the brightest bursts).  While these possible selection effects are still a matter of debate, this correlation (as well as the Ghirlanda one) appears to be satisfied by the (few) newly discovered Hete--II and Swift GRBs.  Furthermore, more sophisticated statistical tests, based on the large BATSE sample of GRBs (\\cite{Nakar2005,Band2005} and \\cite{Ghirlanda2005,Bosnjak2005}), have been performed. ",
        "conclusions": ""
    },
    "0512/astro-ph0512155_arXiv.txt": {
        "abstract": " ",
        "introduction": "In a recent paper, Zavattini  et al. \\cite{Zavattini} have reported a rotation of polarization of light in vacuum in the presence of a transverse magnetic field. If we interpret this rotation in terms of the coupling of a light pseudoscalar particle to photons, \\begin{equation} {\\cal L}_{\\phi\\gamma\\gamma} = {1\\over 4 M_\\phi} \\phi F_{\\mu\\nu}\\tilde F^{\\mu\\nu}\\ , \\label{Lphigg} \\end{equation} we find that the allowed range of parameters is $1\\times 10^5\\ {\\rm GeV} \\le M_\\phi \\le 6\\times 10^5\\ {\\rm GeV}$ and $0.7\\ {\\rm meV} \\le m_\\phi \\le 2\\ {\\rm meV}$ \\cite{Zavattini}, where $m_\\phi$ is the mass of the pseudoscalar. This range of parameters does not conflict with any laboratory bounds \\cite{Cameron}. However these values are ruled out by astrophysical considerations \\cite{astrobounds} if we assume that the pseudoscalar is an axion. For a review of axion physics see, for example, Ref. \\cite{review}. The most stringent astrophysical limit comes from SN1987A which suggests that the mass of the axion cannot be greater than 0.01 eV \\cite {PDG}. Assuming that this particle is an invisible axion, this in turn implies an upper bound on the coupling, $g_{\\phi\\gamma\\gamma} < 2 \\times 10^{-12} \\rm{GeV}^{-1}$, upto a model dependent factor of order unity. The limit on the axion mass arises by demanding that axion emission should not lead to too much energy loss from the core. Similar but less stringent limits can be obtained by considering energy loss from the core of the sun. The coupling, Eq. \\ref{Lphigg}, also leads to several interesting astrophysical polarization effects \\cite{polarization}. For standard axion, its mass and coupling to photons are both related to the Peccei-Quinn scale \\cite{PQ}. It is clearly of interest to see if the astrophysical bounds can somehow be evaded \\cite{Ringwald}. All the astrophysical bounds assume that the pseudoscalar couplings are so small that once produced it will freely escape from the source, which may be the sun or a red giant or a supernova. In the present paper we examine whether the pseudoscalar particles might be trapped inside the sun. ",
        "conclusions": "In conclusion we find that it is possible to evade all astrophysical limits on pseudoscalar photon coupling if the self coupling of pseudoscalars is sufficiently strong. The pseudoscalars start accumulating inside the sun due to scattering with other pseudoscalars. The energy per pseudoscalar degrades due to higher order processes as well as by energy loss due to scattering on electrons, nucleons and pseudoscalars. The pseudoscalar number density eventually reaches steady state due to conversion of pseudoscalars back into photons. We find that the mean free path of these particles inside the core is much smaller in comparison to the mean free path of photons. Hence they contribute almost negligibly to radiative transport. In our analysis we have used the pseudoscalar coupling and mass extracted by the PVLAS experiment. Our results show that the results of the PVLAS experiment are consistent with astrophysical limits if we allow strong self couplings.  \\bigskip \\noindent {\\bf Acknowledgements:} We thank Pasquale D. Serpico and Javier Redondo for a useful communication."
    },
    "0512/astro-ph0512225_arXiv.txt": {
        "abstract": "We report high spatial resolution VLA observations of the low-mass star-forming region IRAS 16293-2422 using four molecular probes: ethyl cyanide (CH$_3$CH$_2$CN), methyl formate (CH$_3$OCHO), formic acid (HCOOH), and the ground vibrational state of silicon monoxide (SiO).  Ethyl cyanide emission has a spatial scale of $\\sim$20$''$ and encompasses binary cores A and B as determined by continuum emission peaks.  Surrounded by formic acid emission, methyl formate emission has a spatial scale of  $\\sim$6$''$and is confined to core B.  SiO emission shows two velocity components with spatial scales less than  2$''$ that map $\\sim$2$''$ northeast of the A and B symmetry axis.  The redshifted SiO is $\\sim$2$''$ northwest of blueshifted SiO along a position angle of $\\sim$135$^o$ which is approximately parallel to the A and B symmetry axis.  We interpret the spatial position offset in red and blueshifted SiO emission as due to rotation of a protostellar accretion disk and we derive $\\sim$1.4 M$_{\\odot}$ interior to the SiO emission.  In the same vicinity, Mundy et al.\\ (1986) also concluded rotation of a nearly edge-on disk from OVRO observations of much stronger and ubiquitous $^{13}$CO emission but the direction of rotation is opposite to the SiO emission findings. Taken together, SiO and $^{13}$CO data suggest evidence for a counter-rotating disk.  Moreover, archival BIMA array $^{12}$CO data show an inverse P Cygni profile with the strongest absorption in close proximity to the SiO emission, indicating unambiguous material infall toward the counter-rotating protostellar disk at a new source location within the IRAS 16293-2422 complex.  The details of these observations and our interpretations are discussed. ",
        "introduction": "IRAS 16293-2422 is a low-mass star forming region located in the $\\rho$ Ophiuchus cloud complex at a heliocentric distance of 160 pc.  It contains an undetermined number of protostellar objects, high velocity outflows in the E-W and NE-SW directions (Stark et al.\\ 2004), and the region is dominated by two radio continuum peaks designated cores A and B. These two principal cores are separated by $\\sim$5$''$ and are often referred to as the ``binary'' system with a mass of 0.49 M$_{\\odot}$ for core A and 0.61 M$_{\\odot}$ for core B (Looney, Mundy \\& Welch 2000). Moreover, the gas and dust toward core A appears to be at a higher temperature than core B and most of the molecular emission lines of high energy transitions have been detected exclusively toward core A.  The measured dust temperature toward core B is $\\sim$40 K (Mundy et al.\\ 1986) and the measured temperature toward core A from many different molecular species is between 80 and 200 K (Chandler et al.\\ 2005).  Spectral observations with the IRAM 30-m radio telescope toward the low-mass protostellar system IRAS 16293-2422 (Cazaux et al.\\ 2003) demonstrated an extremely rich organic inventory with abundant amounts of complex oxygen and nitrogen bearing molecules including formic acid (HCOOH), methyl formate (CH$_3$OCHO), acetic acid (CH$_3$COOH) and methyl cyanide (CH$_3$CN) which are archetypal species often found in massive hot molecular cores (HMCs). Subsequent spatial observations by Kuan et al.\\ (2004) with the Submillimeter Array (SMA) mapped emission from high energy transitions of CH$_3$OCHO, vibrationally excited vinyl cyanide (CH$_2$CHCN) and other large molecular species. Additionally, Bottinelli et al.\\ (2004) mapped slightly lower energy transitions of CH$_3$CN and CH$_3$OCHO with the Plateau de Bure interferometer, noting that all the molecular emission is contained toward cores A and B and that no large molecule emission is seen associated with the molecular outflows.  Recently, Chandler et al.\\ (2005) detected large abundances of a number of different organic molecular species primarily located toward core A from SMA observations.  Moreover, from observations of sulfur monoxide (SO) absorbing against the background dust continuum near binary core B, Chandler et al.\\ (2005) suggest that there must be material from an extended molecular envelope falling onto an embedded dust disk and that there may be a low-velocity outflow coming from an embedded protostar.  The discovery of large oxygen and nitrogen bearing molecules in a low-mass star forming region is significant because the molecular complexity appears to be similar to high-mass star forming regions. Such complex molecules undoubtedly are incorporated into subsequent formation of protoplanetary bodies and would provide a rich organic inventory that could be supplied to early planetary systems similar to that expected to have occurred in our own pre-solar nebula. With this in mind, we conducted high resolution observations of IRAS 16293-2422 using the NRAO\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation, operated under cooperative agreement by  Associated Universities, Inc.} Very Large Array (VLA) using four molecular probes: ethyl cyanide (CH$_3$CH$_2$CN), methyl formate (CH$_3$OCHO), formic acid (HCOOH) and the ground vibrational state of silicon monoxide (SiO) to determine the implication of their relative spatial locations.  Surprisingly, the location of SiO is not co-spatial with either core A or B.  This prompted us to process archival BIMA array $^{12}$CO data for an explanation. ",
        "conclusions": "To summarize, we find a low temperature envelope of molecular CH$_3$CH$_2$CN emission that is encompassing the IRAS 16293-2422 cores A and B.  Furthermore, low frequency, low energy transitions of large oxygen bearing molecules (i.e. HCOOH and CH$_3$OCHO) are seen toward core B where an outflow may be responsible for a low velocity shock that is leading to enhanced abundances of large oxygen bearing species.  Also,  the SiO emission shows two velocity components with spatial scales less than  2$''$ and a velocity separation of $\\sim$5.6 km s$^{-1}$. We interpret the spatial position offset in red and blueshifted SiO emission as due to the rotation of a protostellar accretion disk and we derive $\\sim$1.4 M$_{\\odot}$ interior to the SiO emission.  While it could be that the SiO complex we observe is due to outflow, evidence is accumulating to the contrary that rotation is the explanation.  Such evidence is best seen on a small scale.  The SiO observations in the present work show two velocity components each with linewidths of $\\sim$5 km s$^{-1}$, suggesting rotation of a disk seen edge-on since the two components are not cospatial, are symmetric about an infall location, and have a velocity gradient perpendicular to large-scale CO outflow.  Moreover, images of other molecules on a small-scale also show that their velocity gradients are orthogonal to the direction of outflow seen on a large scale (e.g., Sch\\\"{o}ier et al.\\ 2005).  In this work, we have been very careful to match the synthesized beam to the source emission or absorption to pin-point the location of a new source of infall $\\sim$3.$''$5 north of core source A and $\\sim$2.$''$5 west of core source B.  We have done this by two independent sets of observations (i.e., the VLA in SiO emission and the BIMA array in $^{12}$CO absorption). While there are no previous reports of continuum emission at this location, an inverse P Cygni profile is absolute proof of absorption against background continuum, and our BIMA data show that the location is smaller than the synthesized beam (i.e., 3.2$''$$\\times$0.8$''$).  Thus, the continuum source is undoubtedly small (e.g. a disk seen edge-on) and may eventually be detected by an interferometer at high resolution with enough integration time.  However, the compact disk containing the SiO emission is {\\it counter-rotating} with respect to the $^{13}$CO and CS emission seen by other investigators. This is the first report of evidence for a counter-rotating accretion disk toward a low-mass protostellar complex. Moreover, archival BIMA array $^{12}$CO data show an inverse P Cygni profile with the strongest absorption in close proximity to the SiO emission, indicating unambiguously, material infalling toward the counter-rotating protostellar disk.  The infall location is compelling evidence for a new protostellar source within the IRAS 16293-2422 complex.  We thank F. Sch\\\"{o}ier and C. Ceccarelli for valuable comments on this work, and R. Crutcher for permission to utilize the BIMA array data archive to explore $^{12}$CO data toward IRAS 16293-2422. J.M.H.\\ gratefully acknowledges research support from H.A.\\ Thronson while assigned to the NASA Science Mission Directorate. \\clearpage"
    },
    "0512/astro-ph0512539_arXiv.txt": {
        "abstract": "We performed a spectroscopic search for binaries among hot Horizontal Branch stars in globular clusters. We present final results for a sample of 51 stars in NGC\\,6752, and preliminary results for the first 15 stars analyzed in M\\,80. The observed stars are distributed along all the HBs in the range 8000~$\\leq$~T$_{\\mathrm{eff}}$~$\\leq$~32000~K, and have been observed during four nights. Radial velocity variations have been measured with the cross-correlation technique. We carefully analyzed the statistical and systematic errors associated with the measurements in order to evaluate the statistical significance of the observed variations. No close binary system has been detected, neither among cooler stars nor among the sample of hot EHB stars (18 stars with T$_{\\mathrm{eff}}\\geq$~22000~K in NGC\\,6752). The data corrected for instrumental effects indicate that the radial velocity variations are always below the 3$\\sigma$ level of $\\approx$~15~km~s$^{-1}$. These results are in sharp contrast with those found for field hot subdwarfs, and open new questions about the formation of EHB stars in globular clusters, and possibly of the field subdwarfs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512296_arXiv.txt": {
        "abstract": "{We analysed the fine structure of the phase space distribution function of nearby subdwarfs using data extracted from various catalogues. Applying a new search strategy based on Dekker's theory of galactic orbits, we found four overdensely populated regions in phase space. Three of them were correlated with previously known star streams: the Hyades--Pleiades and Hercules streams in the thin disk of the Milky Way and the Arcturus stream in the thick disk. In addition we find evidence for another stream in the thick disk, which resembles closely the Arcturus stream and probably has the same extragalactic origin. ",
        "introduction": "Fine structure in the velocity distribution of stars in the Milky Way was discovered and studied by O.J.~Eggen during almost all of his career (Eggen 1996 and references therein). Some of Eggens's star streams are associated with young open clusters and can be naturally interpreted as clouds of former members, now unbound and drifting away from the clusters. Other streams contain only stars older than 10 Gyrs. Since for many members distances were not known but had to be assumed in order to construct space velocities, the real existence of such old streams has often been doubted.  However, modern data seem to confirm the concept of old star streams. Helmi et al.~(1999) found  the signature of a cold stream in the velocity distribution of the halo stars of the Milky Way when analyzing Hipparcos data. This was confirmed later by Chiba \\& Beers (2000) using their own data (Beers et al.~2000). Helmi et al.~(1999) interpreted this stream as part of the tidal debris of a disrupted satellite galaxy accreted by the Milky Way, which ended up in the halo. Indeed, numerical simulations have shown that relic stars from disrupted satellites can stay on orbits that are close together for many Gyrs (Helmi et al.~2003, Helmi 2004). These then show up as overdensities in phase space. In the same vein Navarro et al.~(2004) argue that Eggens's (1996) Arcturus group is another such debris stream, but in the thick disk of the Milky Way, dating back to an accretion event 5 to 8 Gyrs ago. These observations complement observations of ongoing satellite accretion such as of the Sagittarius dwarf galaxy (Ibata et al.~1994) or very recent accretion in the form of the Monoceros stream discovered in the outer disk of the Milky Way with SDSS data (Newberg et al.~2002, Yanny et al.~2003, Rocha--Pinto et al.~2003, Pe\\~{n}arrubia et al.~2005). Extended periods of the accretion of satellites onto massive galaxies are also theoretically expected. For instance, recent sophisticated simulations of the formation of a disk galaxy in the framework of cold dark matter cosmology and the cosmogony of galaxies by Abadi et al.~(2003a, b) suggest that disrupted satellites significantly contribute not only to the stellar halo but also to the disk of a galaxy.  Old moving groups are also observed in the velocity distribution of thin disk stars in the solar neighbourhood. Using {\\sf Hipparcos} parallaxes and proper motions, Dehnen (1998) found  new evidence of the Sirius--UMa, Pleiades--Hyades, and Hercules star streams by statistical methods. Even more convincingly these streams show up in the extensive data sample of the three--dimensional kinematical data of F and G stars in the solar neigbourhood by Nordstr\\\"om et al.~(2004, hereafter NMA+). The crowding of these stars on orbits in certain parts of velocity space is attributed to dynamical effects. Dehnen (2000) and Fux (2001) demonstrate that the Hercules stream may well be due to an outer Lindblad resonance of the stars with the central bar of the Milky Way. The Sirius--UMa and Pleiades--Hyades streams, on the other hand, are probably due to orbital resonances of stars in the solar neighbourhood with spiral density waves in the Milky Way disk (De Simone et al.~2004, Quillen \\& Minchev 2005). However, there are also hints that further overdensities in velocity space might be relics of accreted satellites (Helmi et al.~2005).  Here we use our own data (Arifyanto et al.~2005; hereafter AFJW) on the kinematics of nearby subdwarfs and develop a new strategy to search for signatures of old star streams in the phase space distribution of the stars. We then cross check our findings with the NMA+ data. ",
        "conclusions": "We split our sample up into two subsets with metallicities of [Fe/H] $>$ -0.6 and [Fe/H] $\\leq$ -0.6, respectively.  \\subsection{Thin disk}  The stars in our sample with metallicities [Fe/H] $>$ -0.6 dex have kinematics of the old thin disk of the Milky Way. Of course the metallicity cut is somewhat arbitrary, because the thin and thick disk populations do not have a bimodal metallicity distribution, but the transition is quite gradual. In Fig.~1 we show the distribution of 309 stars, which have $|W|$ velocities $<$ 50 km/s, over $\\sqrt{U^2+2\\,V^2}$ versus $V$ and $|W|$ versus $V$, respectively. The space velocities have been reduced to the local standard of rest by adding the solar motion $(U, V, W)_{\\odot} = (10.0, 5.2, 7.2)$ km/s (Dehnen \\& Binney 1998) to the observed space velocities. Instead of scatter plots we show colour coded wavelet transforms of our data in Fig.~1. For this purpose we used the two--dimensional Mexican--hat wavelet transform described by Skuljan et al.~(1999). After some experimentation we found that a wavelet scale of 10 km/s showed the overdensities in the data samples in the clearest way. The Hercules stream at $V \\approx$ -40 km/s is clearly visible as is, to a lesser degree, the Hyades--Pleiades stream at $V \\approx$ -15 km/s, and in both cases exactly where expected (Dehnen 2000, NMA+). Since these streams have been discussed widely in the literature, we do not go into any further details. We present them mainly to demonstrate that by recovering previously known streams our method is well--suited to searching for cold star streams. \\begin{figure} \\centering \\vbox{ {\\includegraphics[scale=0.5]{4355f1.eps}} \\\\ {\\includegraphics[scale=0.5]{4355f2.eps}}} \\caption{Wavelet analysis of the distribution of thin disk stars over $\\sqrt{U^2+2\\,V^2}$ versus $V$ (top panel) and over $|W|$ versus $V$ (bottom panel). The wavelet scale of the Mexican hat kernel is 10 km/s and a linear colour table from black over lilac, green, yellow to red is adopted.} \\label{fig:1}       % \\end{figure} On the other hand, it is instructive in order to assess the reality of such overdensities in phase space to compare the observed distribution with Monte Carlo simulations of realisations of a smooth distribution. In Fig.~2 we show such a Monte Carlo simulation analysed in the same way as the observations. Three hundred nine stars were distributed in the range -200 km/s $<$ $V$ $<$ 50 km/s and $\\sqrt{U^2+2\\,V^2}$ $<$ 300 km/s, respectively, according to a Schwarzschild distribution \\begin{equation} f \\propto \\exp{-\\frac{1}{2}\\left[ \\left( \\frac{U}{\\sigma_{\\rm U}}\\right)^2 + \\left( \\frac{V - \\bar{V}}{\\sigma_{\\rm V}}\\right)^2 \\right] }\\,, \\end{equation} with parameters $\\sigma_{\\rm U}$ = 45 km/s, $\\sigma_{\\rm V}$ = 32 km/s, and $\\bar{V}$ = 26 km/s, which have been derived from the {\\sf CNS4} catalogue as representative of disk stars in the solar neighbourhood (Jahrei{\\ss} \\& Wielen 1997). As can be seen from Fig.~2 the simulated distributions looks generally smoother than the observed distribution but also show considerable Poisson fluctuations that can be confused with real cold star streams. The only remedy for detecting real star streams is obviously to search for such streams in separate data sets. \\begin{figure} \\centering \\vbox{ {\\includegraphics[scale=0.5]{4355f8.eps}} } \\caption{Same as Fig.~1, but for 309 stars distributed randomly in phase space according to a smooth Schwarzschild distribution} \\label{fig:2}       % \\end{figure}  \\subsection{Thick disk}  The remaining stars of our sample with metallicities [Fe/H] $\\leq$ -0.6 dex belong to the thick disk and halo of the Milky Way. The distribution of 382 stars is shown in Fig.~3 in the same way as above, but now restricted to $|W|$ $<$ 100 km/s. There are two distinct features in the phase--space distribution function. The lesser feature at $V\\approx$ -125 km/s corresponds to the familiar Arcturus stream (Eggen 1996, Navarro et al. 2004). The stars in this phase space region are listed in Table 1 (available only in electronic form) giving all relevant data. The kinematics and metallicities of the stars listed in Table 1 can be compared with those of the stars considered by Navarro et al.~(2004; their Figs.~2 and 3) as members of the Arcturus group. Actually there is one common star, G2-34. The kinematics and metallicities agree so well with each other that, even though the reality of low number overdensities is difficult to assess, we are confident that both investigations have identified the same stream. Arcturus itself, although not a CLLA star, lies in Fig.~3 at $V$ = -114 km/s, $\\sqrt{U^2+2\\,V^2}$ = 165 km/s, and  $|W|$ = 4 km/s, respectively. With a metallicity of [Fe/H] = -0.55 (Luck \\& Heiter 2005), it fits well to the rest of the presumed stream members. We place the centre of the stream at $V$ = -125 km/s and $\\sqrt {U^2+2V^2}$ = 185 km/s implying $|U|$ = 55 km/s. According to Eq.~(7) the guiding centre radius of the orbits of the stars now passing close to the Sun is $R_0 = 0.43\\,R_{\\odot}$ = 3.5 kpc. The eccentricity is $e_{R_{0}}$ = 0.59 implying an outer turning radius of $R_t = 2.5\\,R_0$ = 8.5 kpc. The stars are apparently close to apogalacticon, when they are at their slowest on their orbits and the detection probability is highest. In Fig.~4 we show a colour--magnitude diagram of the presumed members of the Arcturus stream listed in Table 1. Overlaid are theoretical isochrones of subdwarfs with an age of 12 Gyrs calculated for metallicities [Fe/H] = -0.5, -1, and -1.5, respectively (Yi et al.~2001). The good fit of the isochrones indicates that the selected stars must be very old. In particular the tip of the main sequence fits well to the turn--off points of the isochrones. This is not a selection effect in the AFJW sample, because the brightest stars in the sample have absolute magnitudes of $M_{\\rm V}$ $<$ 4 mag. Judging from the ages and metallicities of the stars and the similarity of their kinematics with that of debris from a disrupted satellite in the vicinity of the Sun, we follow Navarro et al.~(2004) in concluding that the members of the Arcturus stream are of extragalactic origin. \\begin{figure} \\centering \\vbox{ {\\includegraphics[scale=0.5]{4355f3.eps}} \\\\ {\\includegraphics[scale=0.5]{4355f4.eps}}} \\caption{Same as Fig.~1, but for thick disk stars.} \\label{fig:3}       % \\end{figure}  \\begin{figure} {\\includegraphics[scale=0.5]{4355f5.ps}} \\caption{Colour--magnitude diagrams of the presumed members of the Arcturus stream (left panel) and the proposed new stream (right panel). Overlaid are theoretical isochrones for subdwarfs with an age of 12 Gyrs and and metallicities of [Fe/H] = -0.5, -1, and -1.5 (from right to left).}  \\label{fig:4}       % \\end{figure}  As can be seen in Fig.~3 there is a second strong feature in the phase--space distribution of the thick--disk stars. This seems to be even more significant than the overdensity in the Arcturus region. The stars in this overdensely populated region are listed in Table 2. To our knowledge the existence of a cold star stream in this part of phase space has not been suggested before. Comparing Figs.~3 and 1 we find a clear indication of a corresponding density enhancement at $V$ $\\approx$ -70 km/s in Fig.~1. This phase--space feature can thus also be traced among the more--metal rich stars, but is more prominently seen in the metal--poor population.  In order to test the robustness of our findings we analysed the {\\sf Copenhagen--Geneva Survey} of nearby F and G stars (NMA+). This is based on {\\sf Hipparcos} parallaxes, {\\sf Tycho--2} proper motions, radial velocities, and Str\\\"omgren photometry measured by the authors themselves. We have drawn all those stars from the catalogue with metallicities [Fe/H] $<$ -0.6 and show a wavelet analysis of the phase--space distribution function of 591 stars in Fig.~5 in the same way as in Figs.~1 and 3. \\begin{figure} \\centering \\vbox{ {\\includegraphics[scale=0.5]{4355f6.eps}} \\\\ {\\includegraphics[scale=0.5]{4355f7.eps}}} \\caption{Same as Fig.~1, but for stars  drawn from the Nordstr\\\"om et al.~(2004) sample with metallicities [Fe/H] $<$ -0.6.} \\label{fig:5}       % \\end{figure} Comparing Figs.~5 and 3 it becomes immediately clear that the NMA+ sample is much more fully populated at low space velocities. This reflects that the latter is kinematically unbiased, whereas the CLLA sample is biased towards high proper motion stars. Eggen's classical moving groups show up very clearly: the Sirius--UMa group at positive $V$ velocities, the Pleiades--Hyades stream at $V$ $\\approx$ -20 km/s, and the Hercules stream at $V$ $\\approx$ -40 km/s.  Even though these moving groups were originally devised by Eggen among thin disk stars with metallicities of [Fe/H] $\\approx$ 0, they are also very prominent among the metal--poor stars. This indicates that the origins of the moving groups cannot be dissolving open clusters, but must be due to non--axisymmetric perturbations of the gravitational potential of the Milky Way (Dehnen 2000, Quillen \\& Minchev 2005). In Fig.~5 there is a weak but significant sign of the Arcturus stream, and some of the stars found as members of the Arcturus stream in the AFJW sample appear in the NMA+ catalogue. In our view this is due to the kinematical bias in the AFJW sample, so that the phase space is more richly populated at these negative $V$--velocities than in the NMA+ sample. For instance, in the NMA+ sample there are 144 stars with -200 km/s $<$ $V$ $<$ -50 km/s and $\\sqrt{U^2+2\\,V^2}$ $<$ 400 km/s, respectively, whereas in the AFJW sample there are 191 stars in the same range. However, the overdensity between -100 km/s $<$ $V$ $<$ -60 km/s is clearly discernible in Fig.~5, which confirms the detection of the new cold star stream claimed above.  Our data as given in Tables 1 and 2 show that the velocity and metallicity distributions of the members of the proposed new stream and the Arcturus stream are practically identical. Also the colour--magnitude diagrams shown in Fig.~4 seem to indicate that the stars stem from the same population. We place the centre of the proposed new stream at $V$ = -80 km/s and $\\sqrt{U^2+2V^2}$ = 130 km/s implying $|U|$ = 64 km/s. The mean guiding centre radius of the orbits of these stars now passing close to the sun is $R_0=0.64\\,R_{\\odot}$ = 5.1 kpc. The eccentricity is $e_{R_{0}}$ = 0.42 and the outer turning radius is at $R_{\\rm t} = 1.7\\,R_0$ = 8.7 kpc. Thus the stars of the proposed new stream are also on their orbits close to apogalacticon.  Their orbits are actually very similar to the orbits of the presumed members of the Arcturus stream. We can at present only speculate about the possible origin of the stream. However, the similarity of the characteristics of the new stream with the Arcturus stream seems to point to an extragalactic origin. Moreover, both streams are probably related to each other. Indeed, Helmi et al.~(2005) show that in numerical simulations of the disruption of a satellite galaxy falling into its parent galaxy, the satellite debris can end up in several cold star streams with roughly the same characteristic eccentricities of their orbits. Precisely this seems to be the case here, so that both streams can have very well originated from the same accretion event of a dwarf galaxy into the Milky Way. How the star streams discussed here are related to the star streams reported by Helmi et al.~(2005), especially their groups 2 and 3, has yet to be explored."
    },
    "0512/astro-ph0512069_arXiv.txt": {
        "abstract": "Using mid-IR and optical data, we deduce the total infrared (IR) luminosities of galaxies in the Coma cluster and present their infrared luminosity function (LF). The shape of the overall Coma IR LF does not show significant differences from the IR LFs of the general field, which indicates the general independence of global galaxy star formation on environment up to densities $\\sim$ 40 times greater than in the field (we cannot test such independence above $L_{ir} \\approx 10^{44}~{\\rm ergs~s}^{-1}$). However, a shallower faint end slope and a smaller $L_{ir}^{*}$ are found in the core region (where the densities are still higher) compared to the outskirt region of the cluster, and most of the brightest IR galaxies are found outside of the core region. The IR LF in the NGC 4839 group region does not show any unique characteristics. By integrating the IR LF, we find a total star formation rate in the cluster of about 97.0 $M_{\\sun}{\\rm yr}^{-1}$. We also studied the contributions of early- and late-type galaxies to the IR LF. The late-type galaxies dominate the bright end of the LF, and the early-type galaxies, although only making up a small portion ($\\approx$ 15\\%) of the total IR emission of the cluster, contribute greatly to the number counts of the LF at $L_{ir} < 10^{43}~{\\rm ergs~s}^{-1}$. ",
        "introduction": "Galaxy evolution is largely the story of how the masses, morphologies, and patterns of star formation in these objects vary with environment and cosmological epoch. Luminosity functions (LFs) are very important statistical tools for studying evolutionary changes in galaxy populations and provide key observational constraints on galaxy evolution. The infrared (IR) LF, in particular, by tracing the IR emission of dust heated by star forming activities, remains one of the best probes to study the global evolution of the star formation rate (SFR) with environment and redshift. All the indicators show that the SFR is a strong function of epoch. The average SFR has declined by an order of magnitude to the present epoch from a peak near $z\\approx 1$ \\citep{Lilly96, Madau96, Pablo05, LeFloch05}.  We also expect the SFR to depend on environment. In crowded regions, e.g., the cores of dense clusters, galaxies experience many interactions and mergers, which drive the accumulation of mass and redistribute the gas. The hierarchical galaxy formation models \\citep{Somerville99, Cole00} suggest that galaxies in a cluster gradually lose their gas reservoir as they are accreted into the cluster center. Those environmental effects are expected to result in a systematic modification of the patterns of star formation in clusters. In fact, a number of studies \\citep{Gomez03, Balogh98} have shown such effects, with a lower SFR in clusters compared with the field. The infrared LFs of clusters, as a global indicator of SFR, may also carry such an imprint and differ from the IR LFs of field galaxies.  Since the launch of the {\\it Infrared Astronomical Satellite} (IRAS) and the {\\it Infrared Space Observatory} (ISO), there have been many studies of the infrared LF of galaxies \\citep[e.g.,][]{Rush93, Serjeant01,Takeuchi03}. However, these works concentrate on general field surveys and there is no systematic study of IR LFs with epochs devoted to the cluster population. Such a study will enable us to disentangle the evolution of the SFR in different environments and help us understand where and why the changes in SFR occurred.  The recently launched $Spitzer$ Space Telescope with the Multiband Imaging Photometer \\citep[MIPS,][]{Rieke04}, capable of high resolution, large sky surveys with high sensitivity, provides us the capabilities necessary for thorough studies of the star forming properties of dense galaxy clusters at different redshifts. In this paper, the first one of a series studying the IR LF of clusters up to $z \\approx 0.8$ using MIPS 24 and 70 \\micron\\ observations, we present the IR LF of the Coma cluster. The Coma cluster, as the nearest rich cluster, provides us an excellent chance to study the properties of the IR LF down to a very faint limit. It also enables us to study the change of IR LF in different regions of the cluster, as well as the contributions to the total IR LF from different types of galaxies. This paper will show how the cluster environment shapes the current star formation in this prototypical dense cluster and it will provide a foundation for the future studies of the SFR patterns in other clusters, both nearby and at high redshift. In the paper, we use the cosmological parameter set $(h,\\Omega_{0},\\lambda_{0}) = (0.7,0.3,0.7)$. We assume a distance modulus of m - M = 35.0 mag for the Coma cluster at z=0.023 \\citep{Struble99}. ",
        "conclusions": "Using MIPS 24 \\micron\\ observations and two spectroscopic surveys of the Coma cluster, we present the IR LF of the cluster. The shape of the Coma cluster LF does not differ from that of the general field significantly. The $L^{*}_{ir}$ value of our LF is very similar to those given by \\citet{Rush93} and \\citet{Pablo05}, which are both based on surveys of general fields. The faint end slope of the Coma cluster is shallower than the slope of \\citet{Rush93} but steeper than that of \\citet{Pablo05} and \\citet{Takeuchi03}, again indicating little variation between field and cluster. In addition, the overall proportion of IR-active galaxies in the cluster is only slightly less than in the field. Thus, the overall pattern of star formation in cluster members is surprisingly similar to that in the field, despite an increased galaxy space density by an average factor of $\\sim$ 40.   However, in the cluster core where the galaxy density is six times higher still, we found a shallower faint end slope and a smaller $L^{*}_{ir}$ compared to the outer region of the cluster, which indicates a decrease in the number of faint IR galaxies as well as in the very bright ones. The IR-bright galaxies are distributed around the outer region of the cluster. All the galaxies with $L_{ir} > 10^{44}~{\\rm ergs~s}^{-1}$ lie outside of the core region, e.g, $r > 340$ kpc. No special feature of the IR LF was found in the NGC 4839 region.  In determining the LF of different morphological types, we found that early type galaxies only make about a 15\\% contribution to the total IR luminosity density, but they dominate the number density at the low luminosity end. The global SFR density in the cluster is about 0.88 $M_{\\sun}{\\rm yr}^{-1}$Mpc$^{-3}$ and the total SFR in the 8.5 Mpc$^{2}$ area of the central cluster is about 97.0 $M_{\\sun}{\\rm yr}^{-1}$."
    },
    "0512/astro-ph0512543_arXiv.txt": {
        "abstract": "From a newly obtained VLT/UVES spectrum we have determined the oxygen abundance of HE~1327$-$2326, the most iron-poor star known to date. UV-OH lines yield a 1D LTE abundance of $\\mbox{[O/Fe]}_{\\rm{OH}}=3.7$ (subgiant case) and $\\mbox{[O/Fe]}_{\\rm{OH}}=3.4$ (dwarf case). Using a correction of $-$1.0\\,dex to account for 3D effects on OH line formation, the abundances are lowered to $\\mbox{[O/Fe]}=2.8$ and $\\mbox{[O/Fe]}=2.5$, respectively, which we adopt. Without 3D corrections, the UV-OH based abundance would be in disagreement with the upper limits derived from the O\\,I triplet lines: $\\mbox{[O/Fe]}_{\\rm{trip}}<2.8$ (subgiant) and $\\mbox{[O/Fe]}_{\\rm{trip}}<3.0$ (dwarf). We also correct the previously determined carbon and nitrogen abundances for 3D effects. Knowledge of the O abundance of HE~1327$-$2326 has implications for the interpretation of its abundance pattern. A large O abundance is in accordance with HE~1327$-$2326 being an early Population II star which formed from material chemically enriched by a first generation supernova. Our derived abundances, however, do not exclude other possibilities such as a Population III scenario. ",
        "introduction": "\\citet{HE1327_Nature} recently reported the discovery of the dwarf or subgiant HE~1327$-$2326, the most iron-poor star known to date (with $\\mbox{[Fe/H]}_{\\rm{NLTE}}=-5.4$). Abundances were derived for nine elements and upper limits for a further eight \\citep{Aokihe1327}, including oxygen ($\\mbox{[O/Fe]}<4.0$). No detection of molecular OH lines in the UV was possible from their Subaru/HDS spectrum. A new attempt is presented here to measure the oxygen abundance of HE~1327$-$2326 from different O indicators: UV-OH lines, the [O\\,I] at 6300\\,{\\AA} and the O\\,I triplet at 7774\\,{\\AA}, using a higher quality VLT/UVES spectrum. A measurement of the O abundance of HE~1327$-$2326 is desired for the investigation into the origin of the star.  As the third most common element in the Universe, O generally is an ideal tracer of its chemical history. Hence, it has been studied in extensive detail in metal-poor stars to unravel the earliest evolutionary phases of the Galaxy which is crucial for an understanding of the formation mechanism of the first generations of stars. It is not clear how the first low-mass stars could form in the early Universe. A possibility might involve the C and O yields from Population III supernova which act as sufficient cooling sources in star-forming gas clouds producing the first low-mass stars (e.g. \\citealt{UmedaNomotoNature, brommnature}). However, the picture which emerged from the observational studies is not free from inconsistencies, making the scientific interpretation difficult. A discrepancy of the O abundances derived from different O indicators poses a serious, not yet resolved, problem (for a recent discussion see \\citealt{asplund_araa}).  Our new observations of HE1327-2326 are presented in \\S 2 and the O abundance measurements are described in \\S 3. We discuss the implications in \\S 4. ",
        "conclusions": "\\subsection{Implications of a High Oxygen Abundance} In order to learn about the earliest stages of star formation in the Universe it is very important to identify the origin of the elements observed in HE~1327$-$2326. Oxygen is a key element in this quest because it provides strong constraints on the different origin scenarios previously invoked for the star. Of particular importance is whether HE~1327$-$2326 is an early Population II or a Population III star. Recently, \\citet{iwamoto_science} made an attempt to explain the abundance pattern of HE~1327$-$2326. They invoke a pre-enrichment scenario in which a faint $25\\,M_{\\odot}$ Population~III supernova undergoes a mixing and fallback process producing ejecta containing little iron and large amounts of CNO. Based on the 1D LTE abundances of HE~1327$-$2326 and constrained by the 1D LTE upper limit of oxygen \\citep{HE1327_Nature} they compute an O abundance of $\\mbox{[O/Fe]}\\sim4.0$ which is close to our 1D~LTE abundance derived from OH lines. However, our adopted 3D abundance is significantly lower. It remains to be seen if their model could also reproduce our new CNO values since it might be difficult to simultaneously fit a lower O together with e.g. the high Mg abundance (potential 3D corrections for Mg are expected to be less severe than for OH).  \\citet{meynet05} predict a similarly high oxygen abundance ($\\mbox{[O/Fe]}=3.5$) based on their combined stellar wind and supernova ejecta of their rotating $\\mbox{[Fe/H]}=-6.6$ stellar models. This is in qualitative agreement with the observed excesses of O in HE~1327$-$2326 and other metal-poor stars.  Following \\citet{suda}, a Population~III scenario might explain the origin of HE~1327$-$2326 in terms of a binary system. It would then have accreted its heavier elements from the interstellar medium and the lighter elements from an erstwhile AGB companion in a binary system. However, the absence of significant radial velocity variations (see Figure \\ref{radvel}) over a period of just over one year does not support this idea. Within the overall error there is no change to report in the radial velocity so far. The slight offset between the Subaru and UVES data points in Figure \\ref{radvel} can be accounted to uncertainties in the wavelength calibrations.  Further work is required to ascertain whether the O abundance of \\he can be explained in this manner.  Despite the uncertainties of the corrections to the 1D LTE analysis, it is clear that HE~1327$-$2326 belongs to the group of stars displaying very large CNO abundances. It appears that the majority of these objects have very low metallicities (i.e. $\\mbox{[Fe/H]}<-3.0$) and that HE~1327$-$2326 is the most extreme example of the group. However, HE~1327$-$2326 has a similar overall CNO abundance pattern compared to the only other known star having $\\mbox{[Fe/H]}<-5.0$, HE~0107$-$5240 \\citep{HE0107_Nature, O_he0107}. The unusually high excesses of O of these objects underline that there is no defined trend amongst the stellar O abundances at the lowest metallicities. This suggests that there might not be a simple explanation for the origin of O in the very earliest phases of the Galaxy.  \\subsection{Concluding Remarks} In summary, we adopt the final O abundance to be $\\mbox{[O/Fe]}=2.8\\pm0.2$ (subgiant) or $\\mbox{[O/Fe]}=2.5\\pm0.2$ (dwarf). These values are consistent with the upper limits derived from the O\\,I triplet at $\\sim 7775$\\,{\\AA} and the [O\\,I] line at 6300\\,{\\AA}. This would not be the case if the 1D LTE abundances derived from OH had been adopted. We note here that atomic diffusion might have modified the abundances of HE~1327$-$2326. According to theoretical calculations of \\citet{richard2002}, the O/Fe ratio might originally have been higher. However, observational confirmation of their calculations is still pending. In any case HE~1327$-$2326 provides strong observational evidence that 3D LTE effects for the O abundances derived from OH lines using 1D LTE model atmospheres have to be taken into account, especially for hotter metal-poor stars.  Where already available, such corrections should generally be applied when deriving O abundances for metal-poor stars. A systematic investigation of newly corrected O abundances with respect to metallicity is clearly desirable.  The newly derived O abundance provides additional constraints on the Population II models proposed for HE~1327$-$2326 and other metal-poor stars. Whether or not the new abundance can be reproduced by those models remains to be seen. Finally we wish to mention that the \\citet{iwamoto_science} model does not include neutron-capture elements. Thus it is not clear whether the high Sr abundance in HE~1327$-$2326 could be accounted for with their model. However, recent computations by \\citet{froehlich} indicate that a Sr excess could be in agreement with the faint SN scenario of Iwamoto et al. We note too that the absence of Li lacks explanation. The Population III binary scenario might account for the low Li abundance \\citep{HE1327_Nature, Aokihe1327}, but radial velocity variations have not yet been detected. Hence, a longer time span is needed to monitor the star for such variations in order to draw a final conclusion. In the absence of such data we favor the Population II interpretation of HE~1327$-$2326."
    },
    "0512/astro-ph0512319_arXiv.txt": {
        "abstract": "{We report a discovery of the Zeeman resolved spectral lines, corresponding to the extremely large magnetic field modulus \\bs\\,=\\,17.5~kG, in the cool Ap star HD~178892. The mean longitudinal field of this star reaches 7.5~kG, and its rotational modulation implies the strength of the dipolar magnetic component $B_{\\rm p}\\ge$\\,23~kG. We have revised rotation period of the star using the All Sky Automated Survey photometry and determined $P$\\,=\\,8.2478~d. Rotation phases of the magnetic and photometric maxima of the star coincide with each other. We obtained Geneva photometric observation of HD~178892 and estimated \\teff\\,=\\,7700$\\pm$250~K using photometry and the hydrogen Balmer lines. Preliminary abundance analysis reveals abundance pattern typical of rapidly oscillating Ap stars.} ",
        "introduction": "\\label{intro}  About fifty years have passed since the discovery of the Zeeman resolved line profiles in intensity spectrum of a magnetic chemically peculiar star (Babcock \\cite{B60}). Currently, 47 such objects, with magnetic field strength in the range of 2.8--33.5~kG, are known (Hubrig et al. \\cite{HNS05}). The majority of the less massive (late A) strongly magnetic stars are slow rotators, with rotation periods of months or years. However, an exceedingly strong (up to $\\sim$8~kG) longitudinal magnetic field, variable with a high amplitude and a period of about 8.3~days, was recently discovered in the poorly studied SrCrEu Ap star HD~178892 (BD+$14\\degr\\,3811$, HIP 94155) by El'kin et al. (\\cite{EKR03}) and Kudryavtsev et al. (\\cite{KRE04}). Such a strong field in combination with a relatively fast rotation is not known for any other low-mass Ap star, thereby making HD~178892 a very interesting target for in-depth analysis of the atmospheric structure and magnetic field geometry.  The lack of information about HD~178892 motivated us to obtain new magnetic and photometric measurements, and to acquire the first high-resolution spectroscopic observations for this star. With these new data, Zeeman resolved lines indicating a field modulus \\bs=17.5~kG are detected in the spectrum of HD~178892, which confirms the presence of an extremely strong field in this star. We describe new observations of HD~178892 in Sect.~\\ref{observ} and determine fundamental stellar parameters in Sect.~\\ref{parameters}. The stellar rotation period and magnetic field are discussed in Sect.~\\ref{rotation}. Sect.~\\ref{abundances} reports preliminary abundance analysis results, and conclusions are given in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl} HD~178892 is a new cool photometric, spectroscopic, and magnetic variable. Unusually short rotation period of 8.25~d and a very large magnetic field make this star a perfect target for detailed study of the field geometry and surface chemical inhomogeneities using magnetic Doppler imaging (Piskunov \\& Kochukhov \\cite{PK02}). The position of HD~178892 in the H-R diagram amongst roAp stars, as well as the observed abundance anomalies, make this star a suitable target for the search of rapid oscillations. New photometric and high-resolution spectropolarimetric observations are needed to improve determination of the rotation period and to study in details the surface magnetic field and abundance distributions."
    },
    "0512/astro-ph0512405_arXiv.txt": {
        "abstract": "We test the ability of the numerical action method (NAM) to recover the individual orbit histories of mass tracers in an expanding universe in a region of radius $26\\hmpc$, given the masses and redshift-space coordinates at the present epoch.  The mass tracers are represented by dark matter haloes identified in a high resolution $N$-body simulation of the standard $\\Lambda$CDM cosmology.  Since previous tests of NAM at this scale have traced the underlying distribution of dark matter particles rather than extended haloes, our study offers an assessment of the accuracy of NAM in a scenario which more closely approximates the complex dynamics of actual galaxy haloes.  We show that NAM can recover present-day halo distances with typical errors of less than 3 per cent, compared to 5 per cent errors assuming Hubble flow distances.  The total halo mass and the linear bias were both found to be constained at the 50 per cent level.  The accuracy of individual orbit reconstructions was limited by the inability of NAM, in some instances, to correctly model the positions of haloes at early times solely on the basis of the redshifts, angular positions, and masses of the haloes at the present epoch. Improvements in the quality of NAM reconstructions may be possible using the present-day three-dimensional halo velocities and distances to further constrain the dynamics.  This velocity data is expected to become available for nearby galaxies in the coming generations of observations by SIM and GAIA. ",
        "introduction": "In the standard cosmological paradigm, the formation of large-scale structure is driven by the gravitational amplification of small initial density fluctuations (e.g. Peebles 1980).  In addition to gravity, hydrodynamical processes can influence the formation and evolution of galaxies, groups and clusters of galaxies.  But since hydrodynamical effects play a minor role on scales larger than the size of galaxy clusters, gravitational instability theory alone can directly relate the present day large-scale structure to the initial density field and provide the framework within which the observations can be analyzed and interpreted.  Gravitational instability is a nonlinear process, making numerical methods an essential tool for understanding the observed large-scale structure.~\\footnote{ Present address~: Racah Institute of Physics, The Hebrew University, Jerusalem Israel}  There are two complementary numerical approaches to studying cosmological structure.  The first relies on $N$-body techniques designed to solve an initial value problem in which the evolution of a self-gravitating system of massive particles is determined by forward numerical integration of the Newtonian differential equations. Because the exact initial conditions are unknown, comparisons between these simulations and observations are mainly concerned with general statistical properties.  The second approach works in the opposite direction, deriving from the observed present-day distribution and peculiar motions of galaxies, and independently of the nature of the dark matter, certain features of the dynamics at earlier times.  The numerical action method (NAM) belongs to this second category of approaches.  It arises from the observation that the present-day distribution of galaxies, combined with the reasonable assumption that their peculiar velocities vanish at early times, presents a boundary value problem that naturally lends itself to an application of Hamilton's principle in which stationary variations of the action are found subject to the boundary conditions. The result is a prediction of the full orbit histories of individual galaxies, either with real space boundary conditions (Peebles 1989, 1990, 1994, 1995) or, after a coordinate transformation, in redshift space (Peebles \\etal 2001, Phelps 2002).  The potential of NAM as a probe of galaxy dynamics and of cosmological parameters has been explored in a number of studies following the introduction of the method in Peebles 1989.  Possible applications include the full nonlinear analysis of orbit histories of nearby galaxies (Peebles 1990, 1994, 1995; Sharpe \\etal 2001), recovering the initial power spectrum of density fluctuations (Peebles 1996), predicting the values of cosmological parameters (Shaya \\etal 1995), and estimating the proper motions of nearby galaxies (Peebles \\etal 2000).  Concerning the latter application, ground and space-based observations will soon make possible the measurement of the full three-dimensional velocities of many nearby galaxies and promise both a rigourous test of NAM predictions and, given the additional dynamical constraints on galaxy motions, the possiblity of using NAM as a probe of individual masses of nearby galaxies.  Since a central result of NAM, the past orbit histories of galaxies, cannot be confirmed by direct observations, $N$-body simulations provide an important test of NAM and its key assumption that galaxies can be approximated as discrete, non-merging objects throughout their history.  It is desirable then to test NAM in a scenario which approximates the complexity of the observational situation but where all of the relevant physical quantities are known. Previous tests of NAM using $N$-body simulations have either been confined to a few dark matter haloes at the scale of the Local Group (Branchini \\& Carlberg 1994, Dunn \\& Laflamme 1995), traced the paths of individual dark matter particles rather than extended haloes (Nusser \\& Branchini 2000), or used simulations which demonstrate in principal the ability of NAM to recover particle orbits to a high degree of accuracy but which do not reproduce the full complexity of extended mass distributions (Phelps 2002).  In this paper we extend the tests of NAM to simulations at a scale approaching that of the local supercluster with a catalogue containing several hundred extended objects modelled as particles.  We begin with an overview of the relevant properties of the $N$-body simulation and the halo catalogue we derived from it, and follow with details of the version of NAM used here, which includes a novel approach to the assignment of halo masses.  We will then test the sensitivity of NAM both as a probe of the total mass as well as of the linear bias, and examine in some detail a representative solution, focusing on the comparison between the NAM predictions and the actual halo orbits. ",
        "conclusions": "We have shown, using a catalogue of over 500 dark matter haloes derived from a large $N$-body simulation at the scale of the local supercluster, that it is possible with the numerical action method to reconstruct the full dynamical histories of dark matter haloes given the masses and the redshift space coordinates at the present epoch. The reconstruction is most successful in recovering the halo distances at the present epoch, with typical errors of less than 3 per cent. Individual orbits paths, including the initial positions as well as the direction of motion of the haloes at the present epoch, are predicted with less accuracy.  By varying the relative contributions to the total mass from the haloes and the background, we have also found a way to use NAM to directly measure the linear bias when the total mass density is known, although with an uncertainty of about 50 per cent.  Given the dynamical complexity of millions of interacting particles, and the sweeping nature of NAM's central simplifying assumption that galaxy haloes can be approximated as discrete, non-merging point masses throughout their evolution, it is remarkable how successfully the dynamics of a many-body system can be reconstructed on the basis of an incomplete catalogue of facts.  The successes of NAM as it has been implemented here are of course partially offset by their weaknesses.  Among these is the relatively poor quality of the reconstruction in the vicinity of massive haloes, clearly seen in Fig.~5, which shows a breakdown in the non-linear regime where NAM, which is itself a fully non-linear method, might have potentially offered the most insight.  In these regions $\\chi^2$ is a good indicator of poorly reconstructed orbits, but preliminary attempts to use this information to nudge haloes into the correct orbits while not imposing any additional formal constraints have so far been unsuccessful.  A second concern is the inability of NAM in many cases to isolate, on the basis of $\\chi^2$ alone, predicted halo orbits which are moving in the wrong direction at the present epoch.  Fig.~9 shows, for example, 34 haloes moving more than $90^o$ in the wrong direction but with good distances at the present epoch and thus low $\\chi^2$.  The above innacuracies may in part arise from the details of our implementation, such as our ad hoc procedure of scaling of halo masses according to linear theory, and there is certainly room here for improvement.  The scale of the catalog is also a factor to consider, and in particlar the density of mass tracers within it.  This analysis should be repeated at the scale of the Local Group, where a larger number of mass tracers acting within a smaller volume may better constrain the dynamics and permit more accurate orbit reconstructions.  It is also possible that, in dynamical systems of this complexity, the angular positions, redshifts and masses are by themselves insufficient to lift the degeneracies in the halo orbits, and that the full three-dimensional velocities at the present epoch will be needed to accomplish this.  This again is work to be undertaken at Local Group scales, where next-generation observations from SIM and GAIA hold out the promise of multiple galaxy proper motion measurements with which to test the NAM predictions.  One related concern is that part of the proper motion data may be needed to recover accurate orbits, leaving fewer remaining free parameters to assist with the more weighty problem of constraining individual galaxy halo masses, although it is possible that only one of the two components of the tangential velocity will be sufficient to break the orbtial degeneracy.  Finally, inaccuracies in the NAM predictions are doubtless due at least in part to intrinsic limitations of the method and its assumptions, although we do not wish to suggest at this stage, given the work that remains to be done, that that an upper limit on NAM accuracy in orbit reconstruction has yet been reached.  We anticipate that work on NAM in the near term will lie principally in two directions. The first is an extension of the above analysis, with further improvements in the implementation, to a high-resolution simulation at the scale of the Local Group, where present-day three-dimensional velocities can provide significant additional dynamical constraints.  The second is a direct comparison of NAM with other reconstruction methods, both in real space (e.g, Nusser \\& Dekel 1992, Gramman 1993, Croft \\& Gazta\\~{n}aga 1998, Frisch \\etal 2002, Mohayaee \\etal 2005) and redshift space (e.g., Narayanan \\& Weinberg 1998, Monaco \\& Efstathiou 1999, Mohayaee \\& Tully 2005), that help to bridge the present-day observations of large-scale structure with the initial conditions prevailing in the early universe.  We acknowledge the support of the Asher Space Research Institute.  We would like to thank Felix Stoehr for providing us with the snapshots from his simulation."
    },
    "0512/astro-ph0512174_arXiv.txt": {
        "abstract": "We propose the development of an instrument by the Martin \\& Puplett-type Fourier Transform Spectrometer to applying the aperture synthesis technique in millimeter and submillimeter waves. We call this equipment the Multi-Fourier Transform interferometer(MuFT). MuFT performs a wide band imaging, spectroscopy and polarimetry in millimeter and submillimeter wavelengths. We describe the fundamentals of MuFT, and give an example of one potential implementation. Full description of the observables by MuFT are provided. A physical explanation of the observability of the complex visibility by MuFT is given. Fundamental restrictions on observations with MuFT, eg. limits on spectral and spatial resolutions and field-of-view, are discussed. The advantages of MuFT are also summarized. ",
        "introduction": "Astronomical observation of submillimeter waves represents one of the last unexplored regions of the electromagnetic spectrum. One of the difficulties with performing observations in this band is owing to a large atmospheric absorption. Since submillimeter waves are a boundary of the radio and infrared, there remain a lot of frontiers in the development of fundamental observational technology.  Radio interferometry and synthesis arrays, which are basically ensembles of two elements interferometers, are used to make measurements of fine angular resolution\\cite{Thompson}. Interferometers have been used at millimeter waves\\cite{ishiguro} and have been considered one of the best instruments at submillimeter wave\\cite{ALMA}. By setting two telescopes with separation larger than the maximum construct-able diameter of the single mirror, higher angular resolution than single dish system is easily achieved with aperture synthesis type interferometers. By combining many telescopes, a large effective area can be attained to increase the sensitivity of the system. Interferometers are robust against atmospheric fluctuations since large scale atmospheric fluctuations do not cause interference signals and are automatically omitted by the interferometers.  However, currently operating millimeter and submillimeter interferometers have the following two fundamental problems because they use heterodyne receivers as focal plane detectors: The receiving systems perform modulation and amplification of the source signal by mixing with a local oscillator signal. They can measure the phase of the incident electromagnetic wave. In the following discussion, the space-borne observation assumes as for an ideal case. The noise temperature of heterodyne receivers at frequency $\\nu$ is written as $T_{sys-het}\\sim 2 ({h\\nu }/{k_B})\\left[ (\\epsilon^2 n^{2}(\\nu,T_0) +\\epsilon n(\\nu,T_0))^{\\frac{1}{2}} +1 \\right]$, where $n$ is the thermal Planck function with $T_0=2.73~K$ due to the cosmic microwavebackground (CMB), $h$ and $k_B$ are Planck constant and Boltzmann constant, respectively, and $\\epsilon$ represents the detection efficiencies. The first term represents the photon noise due to statistical fluctuation of the CMB intensity. The second term represents the quantum limit of the heterodyne receivers which comes from the uncertainty principle between the phase and the number of photons\\cite{Rieke}. This equation shows that the noise temperature is limited by the quantum limit and is increased linearly with the frequency in the frequency range of $\\nu >100$~GHz. Therefore, the sensitivity of heterodyne-based interferometers is severely limited by the quantum limit of the receivers. Another problem is the severely limited field-of-view (FOV) of the interferometers, due to the difficulty of constructing a large format detector array of heterodyne receivers \\cite{Rieke}. In addition, the possible bandwidth of the interferometer is also limited by a bandwidth of Intermediate Frequency(IF) amplifier, since the phase change of the modulated signal after mixing with a local oscillator signal must be followed. %  On the other hand, direct detectors such as bolometer are also used in millimeter and submillimeter systems except for interferometers. Direct detectors measure the intensity of the source signal. They cannot measure the phase of the incident electromagnetic wave. Therefore, they have not been used as focal plane detectors for interferometers. The quantum limit of a direct detector is the minimum detectable power when any internal noise can be neglected, the quantum limit of a direct detector is expressed in noise temperature as $T_{qu-dir}\\sim {h\\nu}/(k_{B} \\Delta t \\Delta \\nu)$, where $\\Delta t$ and $\\Delta \\nu$ are the response time of the detector, which is about $1-100$~msec in the case of bolometer, and the bandwidth of the detector, which is about $10-100$~GHz\\cite{Rieke}. The noise temperature of direct detectors at frequency $\\nu$ is written as $T_{sys-dir}\\sim 2 ({h\\nu }/{k_B})\\left[ (\\epsilon^2 n^{2}(\\nu,T_0) +\\epsilon n(\\nu,T_0))^{\\frac{1}{2}} +1/ \\Delta t \\Delta \\nu \\right]$. The scecond term is negligible compared with the CMB photon noise at millimeter and submillimeter waves. Therefore, noise temperature of direct detectors are much lower than heterodyne receivers in the frequency range of $\\nu >100$~GHz. Actually, the noise power of the most sensitive currently operating bolometers is already lower than the quantum limit of the heterodyne in the frequency range of $\\nu>100$~GHz. In Contrast to the heterodyne, a large format millimeter and submillimeter direct detector array is possible to construct\\cite{benford, Matsuo}. There is essentially no limit on the bandwidth for direct detectors\\cite{Zmuidzinas}.  The above discussion indicates that the bolometric interferometer, which is able to use a direct detector as a focal plane detector for the interferometer, could be one of the best ideal instruments for millimeter and submillimeter astronomy since it shares the advantages of both direct detectors and interferometer. Since direct detectors cannot obtain the information of the phase of electromagnetic waves, two beams captured by two apertures must be mixed before guided into the detectors. There are two ways to mix beams. One is called Fizeau type beam mixing. The simplest example of this method is that two beams obtained by masking a single telescope are mixed at the focus of the telescope. Images of different portions of the sky are focused on different positions of the focal plane. Therefore, the extension of the FOV by using the focal plane detector array is straightforward in this case. However, there is a fundamental problem in applying this method to millimeter and submillimeter observations. We have to use many detector pixels within a single FOV to resolve the fringe pattern inside of the FOV. Since the direct detector arrays in millimeter and submillimeter wave bands are still expensive, the cost performance of applying this method to this waveband is poor. The other method is called the Michelson-type beam mixing. An application of the Fourier Transform Spectrometer (FTS) to aperture synthesis interferometry has been studied as a possible solution of the Michelson-type beam mixing bolometric interferometer. The application was first independently proposed by Itoh \\& Ohtsuka(1986)\\cite{Itoh} and Mariotti \\& Ridway (1988)\\cite{Mariotti} in near infrared (NIR). The system was referred by these authors as ''double Fourier interferometry'',  since the signal obtained by the system is a Fourier transformation of both spectra and intensity distribution on the sky. Itoh \\& Ohtsuka studied single pupil interferometry. Mariotti \\& Ridway studied multi-pupil interferometry for high spatial resolution. As explained in this paper, the extension of the FOV by putting the focal plane detector array is not straightforward, but it is possible with some restrictions. Rinehart et al. (2004)\\cite{Rinehart} have been studying the extension of the FOV by putting a detector array on the focal plane of the double Fourier in laboratory experiments in an optical wave bands by using optical CCD camera. They refer to this experiment as the Wide-field Imaging Interferometry Testbed.  In this paper, we propose the application of a Martin \\& Puplett-type Fourier Transform spectrometer (Martin \\& Puplett(1969)\\cite{Ma-Pu}, hereafter MP-FT) to the aperture synthesis system in millimeter and submillimeter waves(Ohta(2004)\\cite{sen}). A wire grid polarizer (WG) is used as a beam splitter in this system. So, the wavelength dependence of reflectivity of WG is small and is suitable for a wide band measurement system. By setting the two input WGs appropriately, 2D intensity distribution of four Stokes parameters can be measured. The signal obtained by our system is a Fourier transformation of spectra and intensity distributions of four Stokes parameters (multiple components) on the sky. Therefore, we refer to this system as a Multi-Fourier Transform interferometer, abbreviated to MuFT. The abbreviation  also contains the meaning that this instrument measures Mutual correlation of the source signal instead of auto-correlation as in the case of usual FTS.  The plan of this paper is as follows: In Section 2, fundamentals of imaging and spectroscopy by MuFT are explained using scalar waves for simplicity. How complex visibility is measured by MuFT is also explained in Sec. 2. In Sec. 3, details of the components of MuFT are introduced. A full description of the observables by the MuFT and restrictions intrinsic to MuFT are also provided in Sec. 3. In Sec. 4, the advantages of the MuFT are summarized. A summary of the paper is given in Sec. 5. This paper presents the fundamental theory of MuFT. Details of broadband imaging and spectroscopy experiments using MuFT are presented in a forthcoming paper. ",
        "conclusions": "We proposed a new type of bolometric interferometer named MuFT by applying MP-FT to aperture synthesis. Fundamentals of imaging, spectroscopy and polarimetry with this instrument were developed. The MuFT is a system which permits imaging and spectroscopy in a wide band by combining the Wiener-Khinchine Formula\\cite{Born} which makes it possible to extract a spectrum from the auto-correlation and van Cittert Zernike Formula\\cite{Born} which allows imaging from the mutual correlation. The concrete composition of this equipment was proposed. By combining wire grid beam-splitters adequately, source intensity distributions of four Stokes parameters can be acquired in a wide band. Fundamental restrictions in practical use were discussed.  We are planning to perform the observations from good mm and submillimeter observation sites, such as the Nobeyama Radio Observatory in Japan, the Atacama desert in Chile, South Pole etc.. The possibility of having a super-wide FOV by mounting the focal plane array of bolometers, SIS photon detectors and transition edge sensors is attractive for future applications. Optical designs of MuFT for using detector arrays is going to be advanced.  Applying this technique to a space-borne mission is one of the best possibilities for extracting the maximum ability of MuFT, since there is no restriction on the bandwidth from atmospheric absorptions. Mather et al. \\cite{Mather} has been proposing a space-borne FIR observatory based on this kind of technique. The future application of this technique to observations from space could open new and interesting possibilities in FIR astronomy."
    },
    "0512/astro-ph0512497_arXiv.txt": {
        "abstract": "Recent observations demonstrate that dwarf elliptical (dE) galaxies in clusters, despite their faintness, are likely a critical galaxy type for understanding the processes behind galaxy formation. Dwarf ellipticals are the most common galaxy type, and are particularly abundant in rich galaxy clusters.  The dwarf to giant ratio is in fact highest in rich clusters of galaxies, suggesting that cluster dEs do not form in groups that later merge to form clusters. Dwarf ellipticals are potentially the only galaxy type whose formation is sensitive to global, rather than local, environment.  The dominant idea for explaining the formation of these systems, through Cold Dark Matter models, is that dEs form early and within their present environments.   Recent results suggest that some dwarfs appear in clusters after the bulk of massive galaxies form, a scenario not predicted in standard hierarchical structure formation models. Many dEs have younger and more metal rich stellar populations than dwarfs in lower density environments, suggesting processes induced by rich clusters play an important role in dE formation.  Several general galaxy cluster observations, including steep luminosity functions, and the origin of intracluster light, are natural outcomes of this delayed formation. ",
        "introduction": "Although dwarf galaxies are the faintest and lowest mass galaxies in the universe, they likely hold important clues for understanding galaxy formation and the nature of dark matter. The reason for this is quite simple: low-mass galaxies, and particularly low-mass galaxies in clusters,   are the most common galaxies in the nearby universe (Ferguson \\& Binggeli 1994). Any ultimate galaxy evolution/formation theory must be able to predict and accurately describe the properties of these objects. In popular galaxy formation models, such as hierarchical assembly (e.g., Cole et al. 2000), massive dark halos form by the mergers of lower mass ones early in the universe, and the first galaxies likely have low stellar mass.  By understanding dwarf galaxies, we are also thus potentially studying the very first galaxies to form.   On the other hand, observations reveal that few low-mass galaxies could have formed all of their stars early in the universe at $z > 7$, with considerable evidence for star formation occurring in the last few Gyrs in dwarf spheroidals (e.g., Mateo 1998).  The traditional approach to studying low-mass galaxies is to examine those in the Local Group (LG). It is now well established that LG dwarf elliptical and dwarf spheroidal galaxies have varying star formation histories, with metal-poor populations as old as classical halo globular galaxies but also with evidence for recent star formation (see e.g., Mateo 1998). There are also some low-mass LG galaxies, such as Sagittarius, that contain surprisingly metal-rich populations given their luminosities (e.g., Ibata, Gilmore \\& Irwin 1995). Because LG dE galaxies are close we can resolve their stellar populations, and thus we can learn much more about them than we do dwarfs in more dense, but distant environments.  Because of this we know a great deal concerning LG dwarf properties including their internal kinematics and star formation histories. Many of the lowest mass galaxies in the LG, such as Draco and Ursa Minor, have very high inferred central M$_{\\rm tot}$/L ratios (Kleyna et al. 2002) and apparently contain the densest dark matter halos of all known galaxies, in qualitative agreement with the original CDM predictions (e.g., Lake 1990). Low-mass galaxies in clusters, however, have different kinematic and spatial properties than these LG systems (e.g.,  Conselice, Gallagher, \\& Wyse 2001,2003), suggesting they might have a different formation scenario.  Observationally, dwarf galaxies are typically faint (M$_{\\rm B} > -18$), with low surface brightnesses ($\\mu_{\\rm B} > 23$ mag arcsec$^{2}$).  Since dwarfs are so common, they are in every sense {\\rm normal} galaxies.  The most common type among dwarfs are dwarf ellipticals/spheroids, which dominate the number density of galaxy clusters down to M$_{\\rm B} = -11$ (Ferguson \\& Bingelli 1994; Trentham et al. 2002).  Based on studies of luminosity functions in clusters, there are also more dwarf ellipticals per giant in denser regions than in the field.  This implies that clusters of galaxies cannot form through simple mergers of galaxy groups. Some dwarf systems must form within the cluster environment. The nature of this over-density may be the result of initial conditions, or `non-standard' galaxy formation. That is, dwarfs may have formed after the cluster was in place.  There is now evidence for this, the implications of which can explain several galaxy cluster phenomenon.  On the other hand, there is also evidence that some dwarf ellipticals/spheroidals in clusters are dominated by old stellar populations.  In this paper we review the current observations of dwarfs in clusters and attempt to interpret these systems in terms of known properties of Local Group dwarf ellipticals/spheroidals, and in the milieu of theoretical ideas concerning low mass galaxy formation in a cosmological context. ",
        "conclusions": ""
    },
    "0512/astro-ph0512342_arXiv.txt": {
        "abstract": "I describe electromagnetic model  of gamma ray bursts and  contrast its main properties  and predictions with hydrodynamic fireball  model and its magnetohydrodynamical extension. The  electromagnetic model assumes that  rotational energy of a relativistic,  stellar-mass central source (black-hole--accretion disk system or fast rotating neutron star) is converted into magnetic energy  through unipolar dynamo mechanism, propagated to large distances in a form of relativistic,  subsonic, Poynting flux-dominated wind and is  dissipated directly into emitting particles through current-driven instabilities. Thus, there is no conversion back and forth between internal and bulk energies as in the case of  fireball model. Collimating effects of magnetic hoop stresses lead to strongly non-spherical expansion and formation of jets. Long and short GRBs may develop in a qualitatively similar way, except that in case of long burst ejecta expansion has a relatively short, non-relativistic, strongly dissipative stage inside the star. Electromagnetic and fireball models (as well as strongly and weakly magnetized fireballs) lead to different  early  afterglow dynamics, before deceleration time. Finally, I discuss the models in view of  latest observational data in the Swift era. ",
        "introduction": "Gamma Ray Bursts (GRBs) are conventionally divided into two classes, short-hard and long-soft, distinguished by their duration (with a division near $\\sim 2$ sec) and spectrum hardness \\citep{kmf+93}. Detection of Type Ic supernovae nearly coincidence with long GRBs unambiguously linked them with deaths of massive stars \\citep{smg+03,Hjorth}. Studies of the host galaxies of long GRBs, which turned out to be actively star-forming, further strengthens this association \\citep{Djorgovski}. Recent progress in observations of short bursts showed that on one hand they show qualitatively similar afterglow behavior  (but without any supernovae signature) while on the other hand their energetics was two  to four orders of magnitude smaller and they are preferentially (at the moment of writing three out of four) associated with older stellar population \\citep{Gehr05,Prochaska05,Villasenor,Covi05a,Rett05,fox}. These indirect evidences are consistent with formation of short GRBs in compact star mergers (double neutron stars or black holes--neutron star binaries) and formation of a black hole \\citep[\\eg][]{ross03,aloy05}. ",
        "conclusions": "In this contribution I outlined the underlying assumptions for the ``electromagnetic hypothesis'' for ultra-relativistic GRB outflows. The most striking implications of the electromagnetic hypothesis is that particle acceleration in the sources is due to direct dissipation of electromagnetic energy rather than shocks and that the outflows are cold, electromagnetically dominated flows, at least until they become strongly dissipative.   One of the major drawback of the model is that magnetic dissipation and particle acceleration are very complicated processes, depending crucially on the kinetic and geometric properties of the plasma. This situation may be contrasted with the shock acceleration schemes, where a qualitatively correct result for the spectrum of accelerated particles, {\\it a kinetic property}, can be obtained from simple {\\it  macroscopic } considerations (jump conditions). Example of the Solar corona shows that despite being complicated magnetic dissipation is an effective mean of particle acceleration.  I have discussed possible observational tests of the hypothesis. In particular, interpretation of early afterglow features as being  due to prompt emission seen at large angles, $\\theta \\geq 1/\\Gamma$, allows to  measure radius at which prompt emission has been produced. Large prompt emission radii, $\\sim 6 \\times 10^{15}$ cm seem to be inconsistent with the fireball model, but close to prediction of the electromagnetic model. Internal relativistic motion of ''fundamental emitters\" assumed within EMM  may also explain X-ray flares during early afterglow phases (without a need for long source activity). An important implication of the electromagnetic model is that supernova explosions may be magnetically driven as well \\citep{lw71,bk71,Wheeler,prog03}.  Over the years I have benefited from discussion with many colleagues, too numerous to be named here. In preparing this contribution I am grateful to Roger Blandford, Tomas Janka, Davide Lazzati, Ehud Nakar, Maurice van Putten  and  Stephan Rosswog for  discussions and comments. I am also indebted to Robert Mochkovitch for shearing  his unpublished results."
    },
    "0512/astro-ph0512032_arXiv.txt": {
        "abstract": "Observations and theoretical work suggest that globular clusters may be born with initially very large binary fractions.  We present first results from our newly modified Monte-Carlo cluster evolution code, which treats binary interactions exactly via direct $N$-body integration.  It is shown that binary scattering interactions generate significantly less energy than predicted by the recipes that have been used in the past to model them in approximate cluster evolution methods.  The new result that the cores of globular clusters in the long-lived binary-burning phase are smaller than previously predicted weakens the agreement with observations, thus implying that more than simply stellar dynamics is at work in shaping the globular clusters we observe today. ",
        "introduction": "Observations, in combination with theoretical work, suggest that although the currently observed binary fractions in the cores of globular clusters may be small ($\\sim 10\\%$), the initial cluster binary fraction may have been significantly larger ($\\sim 100\\%$) \\citep{2005MNRAS.358..572I}.  It has been understood theoretically for many years that primordial binaries in star clusters act as an energy source (through super-elastic scattering encounters with stars and other binaries), with a binary fraction of a few percent being enough to postpone deep core collapse for many relaxation times \\citep[see][for discussion and references]{2003ApJ...593..772F}. In addition to playing a large part in the global evolution of globular clusters, dynamical interactions of binaries also strongly affect the formation and evolution of stellar and binary exotica (e.g.\\ low-mass X-ray binaries, blue stragglers).  In previous studies using approximate methods like Monte-Carlo or Fokker-Planck to simulate the evolution of star clusters, binary interactions had generally been treated using recipes culled from the results of large numbers of numerical scattering experiments---although in one case direct integration of binary interactions was performed for equal-mass stars \\citep{2003MNRAS.343..781G}.  The recipes are typically known only for equal-mass binary interactions, thus prohibiting the use of a mass function in the cluster's initial conditions.  In order to model realistic clusters, which contain a wide range of masses, one must numerically integrate each binary interaction in order to resolve it properly.  In this article we present first results from our newly modified Monte-Carlo code, in which we have included for the first time exact integration of dynamical interactions of binaries with arbitrary mass members.  A detailed description of all new modifications (which include physical stellar collisions and improvements to the core Monte-Carlo technique), and results, will be reported in a forthcoming paper. ",
        "conclusions": "Although our models for evolving globular clusters are somewhat simplified---since they include the effects of two-body relaxation and binary interactions, but ignore the effects of stellar evolution---we can still reach some conclusions by comparing with observations. In \\citet{2003ApJ...593..772F} we found that the values of cluster half-mass radius to core radius, $r_h/r_c$, predicted by models using recipes for binary interactions were in reasonably good agreement with observations.  Our new models, which treat binary interactions exactly, predict a significantly lower core energy generation rate, and a consequently smaller core, yielding values of $r_h/r_c$ up to an order of magnitude larger.  This disagreement suggests that processes other than simply stellar dynamics---such as stellar evolution---may play a very important role in shaping the clusters we observe today."
    },
    "0512/astro-ph0512562_arXiv.txt": {
        "abstract": "{} {We intend to spatially and spectrally resolve the \\OI emission region in two nearby Herbig stars.} {We present high-resolution (\\res = 80,000) VLT/UVES echelle spectra of the \\OI 6300\\r{A} line in the Herbig Ae/Be stars \\object{HD~97048} and \\object{HD~100546}. Apart from the spectral signature, also the spatial extent of the \\OI emission region is investigated. For both stars, we have obtained spectra with the slit positioned at different position angles on the sky.} {The \\OI emission region of HD~100546 appears to be coinciding with the dust disk, its major axis located at 150$\\pm$11\\deg~east of north. The SE part of the disk moves towards the observer, while the NW side is redshifted. The \\OI emission region rotates counterclockwise around the central star. For HD~97048, the position angle of the emission region is 160$\\pm$19\\deg~east of north, which is the first determination of this angle in the literature. The southern parts of the disk are blueshifted, the northern side moves away from us. Our data support the idea that a gap is present at 10~AU in the disk of HD~100546. Such a gap is likely planet-induced. We estimate the mass and orbital radius of this hypothetical companion responsible for this gap to be $20~M_\\jup$ and 6.5~AU respectively.} {Based on temporal changes in the \\OI line profile, we conclude that inhomogeneities are present in the \\OI emission region of HD~100546. These ``clumps'' could be in resonance with the suggested companion, orbiting the central star in about 11~yr. If confirmed, these observations could point to the existence of an object straddling the line between giant planet and brown dwarf in a system as young as 10 million years.} ",
        "introduction": "HD~97048 and HD~100546 are two of the most famous Herbig Ae/Be stars. The shape of the mid-IR spectral energy distribution (SED) suggests that their circumstellar disks have a flared geometry \\citep{dominik03}. For HD~100546, broad-band images ---sensitive to scattered light--- of the disk structure have been obtained in the optical and near-IR \\citep{pantin00,augereau01,grady01}. Furthermore, the disk around HD~100546 has been spatially resolved at UV--to--mm wavelengths in the continuum as well as in spectral features emanating from the disk \\citep[e.g.][]{wilner03,liu03,vanboekel04,leinert04,grady05}. The disk's appearance and position on the sky are hence relatively well-established for this source: the inclination of the system is $i$$\\approx$50\\deg\\ and the position angle of the major axis is approximately 150\\deg\\ east of north \\citep[e.g.][]{liu03}. The disk of HD~97048 on the other hand has not been observed in scattered light, although the mid-IR (around 10~$\\mu$m) observations of \\citet{vanboekel04} show that the source is resolved in both the continuum and the PAH bands.  Herbig Ae/Be stars are by definition emission line stars. In a significant fraction of these objects, also forbidden transition lines are observed. In a previous paper \\citep[][hereafter AVD05]{ackeoi}, we have studied the \\OI lines at 6300 and 6363\\r{A}. We have shown that for HD~97048 and HD~100546, the spectral line profiles are in agreement with the theoretical profiles emanating from excited neutral oxygen atoms in the surface of a flared, rotating passive disk. We have demonstrated that these atoms cannot be thermally excited, and suggested the excitation may be due to the photodissociation of OH and H$_2$O molecules in the surface layers of the disk by the UV radiation field of the central star.  Since both sources are relatively nearby and the emission region of the \\OI 6300\\r{A} line is expected to be spatially extended, we have tried to obtain spatially resolved high-resolution spectra of the targets around this line. We present the data set and reduction method in the following Section. The analysis of the data and the confrontation with the model are described in Sects. \\ref{analysis} and \\ref{model}. In Sect.~\\ref{planet} we discuss new evidence for the presence of a giant planet around HD~100546, based on the results of the analysis. The mass and orbital radius estimates for this object ($M_\\pl = 20~M_\\jup$ and $R = 6.5$~AU respectively) agree strikingly well with previous studies concerning the gap in HD~100546's disk \\citep{bouwman03,grady05} and the spiral-arm structure in the outer disk \\citep{quillen05}. The final conclusions are summarized in Sect.~\\ref{conclusions}. ",
        "conclusions": "}  In this paper we have provided further evidence that the \\OI 6300\\r{A} emission in the disks around HD~97048 and HD~100546 emanates from the disk surface, as was first suggested in AVD05. Thanks to the combination of high spectral resolution and spatial information on the \\OI emission region contained in the UVES spectra, we are able to determine the the blue- and redshifted part of the disk in the two targets. Although the emission region was not resolved in the sense that the FWHM in the feature is larger than the continuum FWHM, the observations clearly show a variation in spatial peak position in the feature. The dependence of that variation on the slit position angle agrees qualitatively with our \\OI 6300\\r{A} emitting flared-disk model. Based on the observed intensity-versus-radius curve of HD~97048 and HD~100546, we find that the \\OI emission region in our model is more extended than the emission region in the targets.  The major axis of the well-studied circumstellar dust disk around HD~100546 has a known position angle of 150\\deg~east of north. We have independently determined the same PA from our data. Note that two different diagnostics (the scattering dust particles and the oxygen gas responsible for the \\OI line) lead to the same angle. This indicates that the \\OI emission region is linked to the geometrical distribution of the dust, in agreement with the conclusions of AVD05. For HD~97048 we find a major axis PA of 160\\deg~east of north, which is the first determination of this angle in the literature.  Furthermore, we have obtained knowledge on the rotational direction of both disks. The disk of HD~97048 is oriented approximately north-south with the south part rotating towards us and the north part redshifted. The disk of HD~100546 is positioned following a NW--SE line, with the SE part blueshifted and the NW part moving away. For the latter object we have shown that the rotation occurs counterclockwise around the central star. The data for HD~97048 do not allow us to make a similar statement.  For HD~100546, we find evidence for the presence of a gap in the circumstellar disk at about 10~AU, as was first suggested by \\citet{bouwman03} and later confirmed by \\citet{grady05}. The radial intensity distribution of the \\OI emission region indicates the presence of excess emission at $R > 10$~AU in comparison to HD~97048. We suggest that this additional emission emanates from the region on the far side of the gap in the dust disk, where a wall has formed. The gap is likely induced by a massive planet. Based on simple assumptions, we derive a diameter for the gap of $\\sim$4~AU and an orbital radius of 6.5~AU for the planet. The mass of this object is estimated to be $20~M_\\jup$. The companion's orbital period ($\\sim$11~yr) is synchronized with the temporal changes observed in the \\OI line profile. We suggest that the deviations from axisymmetry in HD~100546's \\OI emission region are also planet-induced. The \\OI line profile variations in HD~97048, which displays no evidence for the presence of a massive planet, are far less pronounced. The derived mass and orbital elements of the companion of HD~100546 are in agreement with the conclusions of \\citet{quillen05}. The object cannot be directly detected in HST images, and is massive enough to produce the large-scale spiral-arm structure observed in the outer disk of HD~100546. The estimates derived in the present paper can serve future efforts to detect the companion orbiting HD~100546 by direct imaging with ground-based Extremely Large Telescopes.  The \\OI 6300\\AA\\ emission line in HD~100546 may prove to be a valuable tool to probe the companion's orbital parameters. We will set up a long-term monitoring campaign on the Swiss Euler 1.2m telescope in La Silla (ESO) to study the variability of the line in more detail. The results of this study can provide further refinements in our knowledge of the giant planet's orbital period and hence distance to the star. Monitoring the temporal changes of the \\OI line may also lead to a first estimate of the eccentricity of the planet's orbit.  The observations of HD~100546 mentioned in the present paper all point to the presence of a massive object orbiting this target. Although the estimated mass of the companion suggests that it is a brown dwarf, it appears to have a formation history very alike that of a planet. First, the object is expected to be close to HD~100546, which is surrounded by a massive disk. Second, based on the mass estimate ($20~M_\\jup$) the companion would appear in the ``brown dwarf desert'': a lack of binary companions in the mass range between $10~M_\\jup$ and late-M-type stars \\citep[e.g.][and references therein]{mccarthy04}. \\citet{armitage02} propose that brown-dwarf companions form contemporaneously with the primary. If the disk mass is sufficiently high compared to the brown dwarf mass, disk-companion interactions make the brown dwarf migrate inward. This occurs rather rapidly ($\\sim 1 \\times 10^5$~yr for a massive disk). Given that the observed mass of the disk around HD~100546 is at present still $\\sim 0.3~M_\\odot$ \\citep[e.g.][]{ackesubmm}, the initial disk mass should have been high enough to ensure rapid inward migration of the companion. Furthermore, there does not seem to be an underabundance of single stars throughout the brown dwarf mass range. Single stars form due to self-gravitational contraction. Given the discrepancy between the presence of single and long-lived companion brown dwarfs, is unlikely that brown dwarfs in disk-dominated environments are formed following the same mechanism of self-gravitation.  On the other hand, a $20~M_\\jup$ object in a massive disk may have a similar formation history as a giant planet. Due to gravitational instabilities in the disk, the object can gain mass and grow \\citep[e.g.][]{boss05}. If this mechanism is indeed at work in massive protoplanetary disks, it appears logical that it favors low-mass, planet-like objects over somewhat higher-mass brown dwarfs. Together with the fast removal of initially formed brown dwarf companions, this would explain the brown dwarf desert which is observed in multiple systems. The HD~100546 system is at least 10 million years old \\citep{vandenancker98}. The presence of a massive object is therefore more logical under the assumption of a formation which occured in the disk (long) after the formation of the central star. We conclude that the object orbiting HD~100546 is a rare specimen of a brown dwarf companion with a formation history which is similar to that of a regular planet.  The disk of HD~100546 is peculiar and extraordinary in many ways. The longevity of the disk is striking in comparison with other pre-main-sequence stars. Most disks dissipate within a few $10^5$~yr, while the disk around HD~100546 is still clearly present. Also the presence of relatively strong accretion is exceptional in such an old system \\citep[Balmer line emission and accretion events detected in the UV, e.g.][]{deleuil04,grady05}. The presence of a planet in the inner parts of the disk would cause a flow of gas which is ``funneled'' towards the central star. Furthermore, the solid-state features observed in the near-to-mid-IR suggest a highly processed dust component, quite similar to the spectra of comets in our solar system \\citep[see e.g. the {\\em Infrared Space Observatory} data described by][]{malfait98b}. The degree of crystallinity of the dust disk makes HD~100546 an outlier in the group of Herbig Ae/Be stars \\citep{ackeiso}. The grain size distribution deduced from the (sub-)mm slope of the spectral energy distribution indicates the majority of the emitting dust grains have small sizes ($\\ll$1~mm), while the dust particles in a lot of younger systems underwent substantial grain growth \\citep[$>$1~mm,][]{ackesubmm}. Now that it becomes more and more clear that a $20~M_\\jup$ companion is present, it is likely that the characteristics which make HD~100546 special all find their roots in the interactions between the disk and this companion. Such interactions clearly maintain the disk. Based on the presence of small and highly crystalline dust particles, we suggest however that the observed disk is not pristine, but rather more evolved. The interactions of the companion and the larger grains and planetesimals in the disk ---which are expected to have formed in an earlier stage--- could tear apart the latter. The small grains which emanate from such destructive interferences then replenish the small grain population in the disk. The occurance of events like these can provide a satisfying solution to the unusual longevity of the disk in some young stellar systems.  A final pecularity which is worth mentioning in this context concerns the unusual photospheric abundances of HD~100546. HD~100546 appears to be a so-called {\\em $\\lambda$~Boo star}: the photosphere displays a deficit in metals (Mg, Si, Fe and Cr), while nitrogen and oxygen have approximately solar abundances \\citep{ackelamboo}. HD~100546 is the only Herbig star in the sample of \\citeauthor{ackelamboo} which clearly displays this behaviour. {\\em Selective accretion} of metal-depleted gas is commonly inferred to explain this selective depletion pattern. The metals condense into dust particles at much higher temperatures than the CNO elements. The radiation pressure on dust grains prevents the latter to fall in, while the gas-state elements are accreted. The gas and dust should however be decoupled in order to prevent the dust to be dragged along onto the star. The presence of a giant planet in the inner parts of the disk around HD~100546 could create the specific conditions needed to produce the selective accretion and consequent photospheric depletion pattern. If most metals are locked inside dust particles (or even planetesimals), the gas that is ``funneled'' by the planet from the disk onto the central star is metal poor. Intriguingly, a deficit of metals in the stellar photosphere is opposite to what is typical of most giant planet systems found to date, in which the stars are primarily metal rich.  Many observational characteristics of HD~100546 can be explained by the presence of a giant, nearby planet. If the argumentation in the present paper is confirmed by future observations, the companion of HD~100546 is the youngest exo-planet discovered until today. The constraints deduced from a study of such a young planet+disk system will prove to be extremely valuable in both the modeling of disk-planet interactions, as well as in improving our comprehension of disk evolution and planet formation."
    },
    "0512/astro-ph0512081_arXiv.txt": {
        "abstract": "The bright GRB~050408 was localized by HETE-II near local midnight, enabling an impressive ground-based followup effort as well as space-based followup from {\\it Swift}.  The {\\it Swift} data from the X-Ray Telescope (XRT) and our own optical photometry and spectrum of the afterglow provide the cornerstone for our analysis. Under the traditional assumption that the visible waveband was above the peak synchrotron frequency and below the cooling frequency, the optical photometry from 0.03 to 5.03 days show an afterglow decay corresponding to an electron energy index of $p_{\\rm lc} = 2.05 \\pm 0.04$, without a jet break as suggested by others.  A break is seen in the X-ray data at early times (at $\\sim$12600 sec after the GRB).  The spectral slope of the optical spectrum is consistent with $p_{\\rm lc}$ assuming a host-galaxy extinction of $A_{V} = 1.18$~mag.  The optical-NIR broadband spectrum is also consistent with $p = 2.05$, but prefers $A_{V} = 0.57$~mag.  The X-ray afterglow shows a break at $1.26 \\times 10^{4}$~sec, which may be the result of a refreshed shock.  This burst stands out in that the optical and X-ray data suggest a large \\ion{H}{1} column density of $\\mnhi \\approx 10^{22} \\cm{-2}$; it is very likely a damped Lyman $\\alpha$ system and so the faintness of the host galaxy ($M_{V} > -18$~mag) is noteworthy. Moreover, we detect extraordinarily strong \\ion{Ti}{2} absorption lines with a column density through the GRB host that exceeds the largest values observed for the Milky Way by an order of magnitude.  Furthermore, the \\ion{Ti}{2} equivalent width is in the top 1\\% of \\ion{Mg}{2} absorption-selected QSOs.  This suggests that the large-scale environment of GRB~050408 has significantly lower Ti depletion than the Milky Way and a large velocity width ($\\delta v > 200$\\kms). ",
        "introduction": "\\label{s:intro}  Leading up to the launch of {\\it Swift} \\citep{Gehrels04}, the astronomical community prepared for massive, multi-wavelength studies of GRBs expected from the satellite.  Not long after the launch of {\\it Swift}, HETE-II \\citep{Sakamoto05} triggered (H3711) GRB~050408 at 16:22:50.93 on 2005 April 8 ({\\sc UT} dates will be used throughout this paper).  Soon after its detection, {\\it Swift} triggered a Target of Opportunity on the GRB \\citep{Wells05}.  Later, a fading optical afterglow was detected \\citep{deUgartePostigo05} and a redshift of $z \\approx 1.236$ was obtained through host galaxy emission lines and afterglow absorption features \\citep{Berger05,Prochaska05:050408}.  Radio observations were also obtained but no transient was found \\citep{Soderberg05}. The X-ray afterglow \\citep{Wells05} was observed over several epochs with Swift, leading to an initial inference of a break \\citep{Godet05}, that was later retracted \\citep{Capalbi05}. Finally with all the XRT data, the Swift team suggested a jet break at $t_{\\rm break} = (1.2 \\pm 0.5) \\times 10^5$\\, sec after the GRB trigger \\citep{Covino05}.  We present light curves of the optical, infrared, and X-ray afterglows in Sections~\\ref{s:opt} and \\ref{s:xray}.  A detailed analysis of these afterglows is presented in Section~\\ref{s:aglow}.  An analysis of the optical and X-ray afterglow spectra is presented in Sections~\\ref{s:xray} and \\ref{s:oabs}.  From the absorption in these spectra we are able to place lower limits on the metallicity and the hydrogen column of the host galaxy.  Throughout the paper, the concordance cosmology of $\\Omega_\\lambda = 0.71$, $\\Omega_m = 0.29$, and $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$ is used. Though all measurements reported herein are consistent with our preliminary reports in the GCN (GRB Coordinates Network) Circulars, these measurements supersede those in the Circulars. ",
        "conclusions": "GRB~050408 is a particularly interesting object showing both the consistency of predicted models and showing new and extreme cases of physical phenomena.  In particular, we have shown:  \\begin{itemize}  \\item The synchrotron electrons had energy index of $p \\approx 2$, the lower limit of physically acceptable systems \\citep{Meszaros97,Sari98}. This is supported directly by the optical-NIR afterglow decay and the X-ray spectrum.  There is also indirect support (assuming particular models) from the optical spectrum, the optical-NIR broadband spectrum, and the X-ray afterglow decay.  \\item The X-ray afterglow shows a break at $1.26 \\times 10^{4}$ sec after the burst.  This break is not attributed to a jet break.  One possible explanation is continued energy injection.   \\item The hydrogen column is very large ($\\mnhi \\approx 10^{22}$~cm$^{-2}$).  The optical spectrum also showed one of the most extreme Ti-absorption systems observed.  The combination of these facts suggest that there is an incredibly low amount of Ti depletion in the environment of GRB~050408.  This has been noted for other GRBs, suggesting that low Ti depletion is linked to GRB environments, possibly due to high-mass star formation, the environments of newly formed supernova and GRB remnants, or dust destruction from the GRB.  \\item The large velocities associated with the absorption lines are not easily explained by the kinematics of the host galaxy.  For a systemic velocity of $v \\approx 150$ \\kms, a large mass (and possibly a special geometry) is needed.  However, we have shown that the host of GRB~050408 is faint $M_{V} > -18$, comparable to the LMC.  This suggests that the velocities originate close to the progenitor, either from a wind from the Wolf-Rayet progenitor star or older supernova explosions close to the progenitor.  \\end{itemize}"
    },
    "0512/astro-ph0512048_arXiv.txt": {
        "abstract": "Responses to questions, comments and criticism of our recent paper ``General Relativity Resolves..\" \\cite{CT} are provided. It is emphasized that our model is entirely natural to describe the dynamics of an axially symmetric galaxy and that our solution, albeit idealized, contains the essence of the problem. The discontinuity of the metric derivative on the symmetry plane is necessarily interpreted as the effect of the mathematically idealized discontinuity of the gradient of the density and is shown to be naturally connected to the distributed volume density via the Gauss divergence theorem. We present arguments to the effect that for our approximate weak field model, we can choose the physically satisfactory mass distribution without an accompanying singular mass surface layer. To support this contention, we modify our solution slightly by removing the discontinuity with a region of continuous density gradient overlapping the $z=0$ plane. The alternative of invoking a surface layer leads to the presence of a negative mass surface layer approaching the numerical value of the positive mass continuous region. This is in contradiction with the assumed stationarity of the model. We find that a test particle behaves normally as it approaches the $z=0$ plane, the acceleration being towards the direction of this plane. This is in contradiction to the negative mass layer hypothesis as negative mass would repel the test particle. Thus, further support is added to the integrity of our original model. ",
        "introduction": "Recently, \\cite{CT} we presented a paper illustrating that Newtonian dynamics is inadequate to describe galactic dynamics. We showed that general relativity, the preferred theory of gravity, is required for the task even though the fields are weak and the motion is non-relativistic. This is because in such a gravitationally bound dynamical problem with an extended matter distribution, non-linearities cannot be neglected. We showed that general relativity allows the modeling of the essentially flat galactic rotation curves without exotic dark matter. In the process, we determined the mass density of the luminous threshold based upon data in the radial direction as given in \\cite{kent}. In the short period since our presentation, we have received a large volume of correspondence with interesting questions, comments, suggestions and criticism (three such already posted \\cite{korz}, \\cite{VL} \\cite{garf}) . \\footnote{ It should be noted that in private communications, two colleagues independently alerted us to the same line of reasoning as in \\cite{korz} prior to that posting. } The essential thrust of \\cite{korz} was the claim that our particular model contained a singular disk of mass in the symmetry plane of the galaxy and thereby, vitiated our solution as a proper model for a galaxy. Subsequently, various individuals have used this argument to claim that our work is flawed. More recently \\cite{VL}, it was suggested that the symmetry plane was the seat of exotic matter with one option being a negative mass sheet. Most recently, it was argued in \\cite{garf} that the standard iterative perturbation scheme accounts for non-linearities and hence the galactic dynamics must be determined by Newtonian theory to lowest order. We will discuss these papers further in what follows.  In this paper, we respond to the criticisms. In the process, some new interesting insights emerge.  What must be emphasized is this: while it is certainly useful to have raised such criticisms, it is important to recognize that in the multitude of comments and correspondence that followed, \\textit{no one to our knowledge has found valid reason to fault our central thesis, that general relativity, the preferred theory of gravity, exhibits essential non-linearities in sources of galactic scope}. While in subsequent studies, one solution to the model may be preferred over another, the essential point is that a new route to astrophysical dynamics is opened up by the recognition that general relativity, long accepted as the key to cosmology, also comes into play with significance for the major building blocks \\textit{within} cosmology.  A frequent criticism of our work is the evidence that is presented for the existence of vast amounts of exotic dark matter in larger than galactic scales such as in the scale of clusters of galaxies. Two points must be made in this regard. Firstly, while some individuals have read into our work that we have claimed to have \\textit{proved} that such large exotic dark matter reservoirs do not exist, this is not the case. Thus far our work applies to the galactic scale. Secondly, we have pointed to the fact that it would be interesting to extend the \\textit{general relativistic} approach to the other relevant areas of astrophysics to determine whether or not exotic dark matter is truly required in those larger scale domains. For example, for the dynamics of clusters of galaxies, the virial theorem is used. This is based on Newtonian gravity theory. It would be of interest to introduce a general relativistic virial theorem for comparison. It is only after possible effects of general relativity are explored that we can be confident about the viability or non-viability of exotic dark matter in nature.  In Section 2, we review the essential structure and equations that were developed in \\cite{CT}. In Section 3, we reply to certain issues that had been raised and in Section 4, we consider the problem of matter distribution. We end with concluding remarks in Section 5. ",
        "conclusions": "In \\cite{CT}, we had modeled a galaxy as an axially symmetric pressureless stationary rotating fluid within the framework of general relativity to order $G$. We had shown that the dynamics was driven by one linear and one non-linear equation as opposed to the linear equation of Newtonian gravity. A framework of solutions with separated variables was established as a sequence of Bessel functions. This enabled us to choose appropriate parameters to fit the flat galactic rotation curves without invoking massive halos of exotic dark matter as is required using Newtonian gravity. The masses were concentrated primarily within the disk configuration. The mass values were found to be between those of Newtonian gravity and those predicted by the MOND model.   The separated variable approach led to exponential dependence in the transverse $z$ variable and reflection symmetry implied a discontinuity in density gradient at the symmetry plane. Critics had claimed that this necessitated an accompanying singular mass layer on this plane.  In this paper, we approached the problem in two ways: firstly, we took the normal route of specifying the source as a singularity-free density distribution with density gradient discontinuity and saw that this led to the solution that we had used in \\cite{CT}. The key was the observation that the field equations only dealt with $z$ derivatives in a form that led to unique limits as one approached the symmetry plane from above or below. This enabled us to specify the values of the functions on the plane as the limit as the plane was approached, from above or below.  Secondly, we bridged the region of the symmetry plane where the density gradient was discontinuous by using a new solution there, one that was symmetric and smooth, in fact infinitely differentiable. We showed that this solution could be metric and metric derivative matched to the original satisfactory exponential fall-off exterior solution.  We also considered the consequences of \\textit{not} demanding a natural non-singular density distribution as was invoked in \\cite{korz} \\cite{VL} \\cite{AW}.  In this case, we found that the mass of the layer was necessarily connected to the mass of the continuum by the Gauss divergence theorem. In turn, this implied a negative mass in the layer that numerically approximated the positive mass of the continuum. Such an enormous negative mass would contradict the assumed stationarity of the model. Finally, we considered the motion of a test mass approaching the $z=0$ plane, thoroughly searching the domain for any sign of repulsion as would be implied by the supposed vast store of negative mass in the $z=0$ plane (assuming that the matter in the plane could somehow actually hold itself together). None was found in contradiction to the layer assumption.  In future work, it will be of interest to refine the process further with finer resolution of sequence solutions and with applications to more galaxies. It would also be of interest to explore the untapped area of non-separable solutions with their singular cores. These might be particularly useful should it be thoroughly established that galactic cores harbour physical singularities, be they clothed or naked.      \\vskip 0.25in  {\\small {\\bf Acknowledgments:} We are grateful to W.B. Bonnor, F.D.A. Hartwick, D.N. Vollick and many colleagues for their valuable comments, suggestions, questions and criticism. This work was supported in part by a grant from the Natural Sciences and Engineering Research Council of Canada.}   \\clearpage {\\small"
    },
    "0512/astro-ph0512612_arXiv.txt": {
        "abstract": "Three models of a flat universe of coupled matter and dark energies with different low-redshift parameterizations of the dark energy equation of state are considered.  The dark energy is assumed to vary with time like the trace of the energy-momentum tensor of cosmic matter.  In the radiation-dominated era the models reduce to standard cosmology.  In the matter-dominated era they are, for modern values of the cosmological parameters, consistent with data from SNe Ia searches and with the data of \\citet{GUR1999} for angular sizes of ultra compact radio sources. We find that the angular size-redshift tests for our models offer a higher statistical confidence than that based on SNe Ia data. A comparison of our results with a recent revised analysis of angular size-redshift legacy data is made,and the implications of our models with optimized relativistic beaming in the radio sources is discussed. In particular we find that relativistic beaming implies a Lorentz factor less than 6,in agreement with its values for powerful Active Galactic Nuclei. ",
        "introduction": "There is now substantial observational evidence\\citep{PERA2003} that favors the existence of a smooth exotic cosmic component of energy of negative pressure.  Going at times under the name of a cosmological constant or quintessence or,at other times,dark energy,which we will adopt here,its true nature remains obscure.  The unexpected faintness of high redshift type Ia supernovae (SNe Ia)suggests that the universe is accelerating today\\citep{RIE1998,PERL1999},relentlessly driven by dark energy. When the SNe Ia results are combined with observations of the amplitudes of primordial fluctuations in the cosmic microwave background radiation the overall picture seems to be one of an accelerating flat universe.  Since the standard flat universe, despite its well-known shortcomings,has long been favored on aesthetic and theoretical grounds\\citep{KT1990},the hope has arisen that the injection of dark energy will cure its ills, particularly in regards of its age of the universe problem.  Thus a major industry of investigating the constraints imposed by continuously updated astrophysical observations on the dark energy in refined versions of the standard model has flourished in recent times. The present paper is one more contribution in this direction.  When it is assumed that the dark energy,viewed in general to be time-dependent,does not interact with matter,the energy equations for nonrelativistic pressureless matter and dark energy decouple leading to conservation of matter and to the dark energy equation $d\\rho_{de}/dz = 3(1+z)^{-1}(1+p_{de}/\\rho_{de})\\rho_{de}$,where $\\rho_{de}(z)$ and $p_{de}(z)$ are the dark energy density and pressure respectively and $z$ is the redshift. In this case a solvable cosmological model is obtained if a specific variation of $\\rho_{de}$ is invoked\\citep{MOZ1999,AMMR2002,ARMOH2005}or a definite parameterization for the equation of state $w(z)\\equiv p_{de}/\\rho_{de}$ is suggested \\citep{CPOL2001,BACCI2003,ALAMSA2003,DR2004,PADMAN2003,CORAS2004,ALAM2004,JOHR2004,JOHR2005}.  Alternatively if one assumes that the dark energy interacts with matter\\citep{MOZ1999},the energy equations for both are coupled and one needs a definite variation for $\\rho_{de}$,in addition to specifying its equation of state. In this case the conservation of the energy-momentum tensor of the field equations,which holds when matter and the dark energy are noninteracting,is replaced by the conservation of the sum of this tensor and an extra appended tensorial piece representing the time-dependent dark energy. Here we follow this line: Specifically we assume (a) $\\rho_{de} \\sim T$,where $T = \\rho -3p$ is the trace of the energy-momentum tensor of cosmic matter of density $\\rho$ and pressure $p$,and (b) a one parameter form for $w(z)$.  A variation $\\Lambda\\sim T$ was introduced by \\citet{M2001,M2003} for the cosmological constant $\\Lambda$,the motivation being to identify the cosmological constant with a Lorentz-invariant scalar representing a form of quintessence. This cosmology is reminiscent of  similar earlier attempts at identifying $\\Lambda$ with the Ricci scalar\\citep{ALMOT1996a,ALMOT1996b,AMMR1997}. The postulate $\\Lambda \\sim T$ is interesting because it implies that the cosmological constant vanishes in the radiation-dominated cosmic era of flat cosmology so that the successful standard primordial nucleosynthesis predictions are unaltered. In the matter-dominated era the postulate reduces to $\\Lambda \\propto H^2$ where $H$ is Hubble's parameter. The cosmological constant variation $\\Lambda \\propto H^2$ itself was widely discussed in the literature \\citep{F1987,CLW1992,LC1994,AW1994,WE1995,AR1997,OC1998,V2001}. In particular \\citet{CLW1992} have pointed out that it follows from dimensional arguments consistent with quantum gravity. Since such arguments do not depend on the cosmological constant equation of state $ p_\\Lambda = -\\rho_\\Lambda$ it is legitimate to regard them as equally valid for dark energy with $w(z) \\neq -1 $.  Extending this postulate to a dark energy with an equation of state of negative pressure we take $\\rho_{de} = \\kappa T$ where $\\kappa$ is a dimensionless constant. A consequence of this is that the matter density parameter $\\Omega_m$ is constant in the model. We take it to be $ 1/3$. This is because a matter density parameter around 0.30 seems to be favored by observations indicating that the dark energy accounts for 2/3 of cosmic matter \\citep{Turner2002a,Turner2002b}. In fact \\citet{Turner2002c} has strongly argued a case for $\\Omega_m = 0.33 \\pm 0.035 $ from measurements of the physical properties of clusters,CMB anisotropies and the power spectrum of mass inhomogeneities.  For the dark energy equation of state we consider 3 models with the one-parameter forms: (1)$w_{de}=w\\equiv const$,(2)$ w_{de}= -1+ wz$,and (3)$w_{de}= -1 + w\\frac{z}{1+z}$, where $w$ is constant.  Model (1), \\emph{viz}, $w_{de} \\equiv w=const <0$,is a generalization of the cosmological constant case $w_{de} =-1$. Strictly speaking a constant $w_{de}$ is valid for the cosmological constant only. Yet models of cosmic evolution driven by nonrelativistic matter and a quintessence component $X$,an exotic fluid with an arbitrary equation of state $ p_X = w_X\\rho_X $ ($w_X \\geq -1 $), have been widely studied \\citep{RAP1988,CHIB1997,TUW1997,SPER1997,FRIEM1998,CALDW1998,EFST1999,Turner2002d}. In a number of these models (particularly those with tracking solutions),both the dark energy density parameter $\\Omega_{de}(\\equiv 8\\pi G\\rho_{de}/3H^2)$ and $w_{de}$ vary so slowly with redshift \\citep{ZLAT1999,STE1999,EFST1999} as to justify the approximate use of an effective equation of state parameter $w_{eff}\\sim\\frac{\\int{w_{de}(z)\\Omega_{de}(z)dz}}{\\int{\\Omega_{de}(z)dz}}$ \\citep{WAN2000,ZHU2004}. More generally,the absence of robust fundamental physics-based dark energy models and the difficulty to observe a time dependence of $w_{de}$ from CMBR\\citep{AUS2003}or from fits to luminosity distances \\citep{PICL2003},admits the possibility of a $w_{de}$ which is constant in some specified range,and which arises as a model-independent approximation to the dark energy equation of state \\citep{KNST2003,CEP2004}.(The cosmology with a dark energy $ \\sim a^{-2}$ and decoupled from ordinary matter so that $w_{de} = -1/3$ has been recently discussed by one of us \\citep{AMMR2002} and by \\citet{ARMOH2005}).  On the other hand parameterizations (2) and (3) are special cases of the two-parameter forms: $w_{de} = w_{de}(0) + wz$,and $w_{de} = w_{de}(0)+ w\\frac{z}{1+z}$ which were proposed by \\citet{HUTU2001} and \\citet{WEAL2002}, and  by \\citet{LIN2003}respectively, and recently studied,together with the case $w_{de}=w\\equiv const$, by\\citet{DR2004}. The form $w_{de} = w_{de}(0)+wz$ diverges at very high redshifts whereas this difficulty is avoided in the model $w_{de} = w_{de}(0)+ w\\frac{z}{1+z}$, where $w_{de} \\rightarrow w_{de}(0)+w$ as $z\\rightarrow \\infty$. But, as argued by \\citet{Riess2004},a safer strategy, which we follow here,is to regard these parameterizations as only valid for low-z ($z\\ll$ decoupling redshift $z_{dec}$) and describing the late behavior of dark energy. This was done by \\citet{DR2004} who studied these parameterizations and found,on taking as prior $\\Omega_m = 0.3$ in a flat universe,that a constant $w_{de} \\equiv w_{de}(0) = -1$ is preferred by the fit to the gold data for type Ia supernovae \\citep{Riess2004,TON2003,BARR2004}.  Also it is already known that the cosmological constant scenario remains consistent with tight constraints from new cosmic microwave background and galaxy clustering data\\citep{ALESS2004}. Quite generally observations seem to require dark energy with present values $w_{de} \\sim -1 $ and $ \\Omega_{de} \\sim 0.7$ \\citep{PERA2003}. With this in mind, and noting that recent SNeIa observations from HST do not indicate a rapid variation of $w_{de}(z)$ away from its cosmological constant value,we pursue,for simplicity,the following approach: we consider the  preceding three $w_{de}(z)$ parameterizations and set in them,$\\tilde{a}b\\ initio$, $w_{de}(0)= -1$. Then we investigate the constraints on the dark energy equation of state from recent supernova data and observations of the angular sizes of ultra compact radio sources.  In section 2 we present the basic equations of the models. In sections 3 and 4 we examine the constraints on them from supernova and angular sizes data respectively. In discussing angular sizes we compare our results with those from a recent work by \\citet{JJ2006},and also consider the implications of our angular size-redshift relations in the presence of relativistic beaming of the radio sources. Section 5 winds up the paper with a discussion of the results and some concluding remarks. ",
        "conclusions": "The model's angular size-redshift curves drawn in Figure 11 fit \\citet{GUR1999} data equally well. This is a reflection of the small values and span of $\\chi_{min}^2$ as seen in Table 5 for the results : for the four models $\\chi_{min}^2$ varies only from 4.47 to 4.75. By comparison in the supernova test the $\\chi_{min}^2$ values are considerably larger and lie in the range 15.4 to 19.67 so that one has more faith in the agreement of the models with the angular size data.  The predictions of the models for the critical minimum redshift in the angular size-redshift relation give $z_m$ values in the narrow interval [1.65,1.72] compared to the standard model result $z_m = 1.25$. The constant-$w_{de} \\neq -1/2$ and the third models have the same $z_m = 1.65$ whereas the coasting universe and model 2 have $z_m 1.72$ (the same as that of \\citet{JDAL2003})and $z_m = 1.71$ respectively. Thus the minimal redshift cannot,by itself,effectively discriminate between these models.  We have also compared the predictions of our models with those of the recent work of \\citet{JJ2006}in which they revisit the old \\citet{PRES1985}catalogue of ultra compact radio sources and reconsider an angular size-redshift data set in the light of modern preferences of the  cosmological parameters. We find that their results do not agree with our models. In a sense this is not surprising since the underlying premises of the two works are different:\\citet{JJ2006}approach is based on using one simple potential to test the hypothesis that vacuum energy is constant. The present models are phenomenological and based on a time-dependent dark energy coupled to matter. Hence there is no overlap between the two approaches\\citep{JJ2007}.  Finally,we tested our angular size-redshift relations in the presence of  relativistic beaming. Relativistic beaming in the FRW $\\theta -z$ relation was investigated by \\citet{DAB1995}. Here we used a simple variation of their method in a two-jet model of radio sources with the advancing (towards us)jet axis close to the line of sight. In this picture consistency of the models with \\citet{GUR1999} data is found for a beaming Lorentz factor $\\gamma < 6$. Such an upper bound is consistent with values of the Lorentz factor for powerful AGN jets \\citep{AEL1990,PEAC1999}.  To conclude, we have presented dark energy models that are in reasonable agreement with the supernova data of \\citet{BARR2004}and in good agreement with the \\citet{GUR1999} compact radio source angular size versus redshift binned data. The three models that we have studied are simplified versions of ones recently considered by \\citet{DR2004} in a different context. But as remarked by these authors comparing models for the equation of state of dark energy will remain something of a mug's game until there exists substantially more data at higher redshifts.  We thank R.G. Vishwakarma for sending \\citet{GUR1999} complete data sets on angular sizes of compact radio sources,and J.C.Jackson for a useful correspondence. We also acknowledge the financial support of the Directorate of Scientific Research of the University of Khartoum. One of us (AMMAR) thanks Professor K.R. Sreenivasan for hospitality at the Abdus Salam International Center for Theoretical Physics, Trieste, Italy, where part of this work was done , and the Swedish Agency for International Development (Sida) for financial support. AMMAR also thanks Dr. Abu Bakr Mustafa for his hospitality at ComputerMan College for Computer Studies, Khartoum, where also part of this work was done."
    },
    "0512/astro-ph0512606_arXiv.txt": {
        "abstract": "A revision of Stod\\'o{\\l}kiewicz's Monte Carlo code is used to simulate the evolution of million body star clusters. The new method treats each {\\it superstar} as a single star and follows the evolution and motion of all individual stellar objects. A survey of the evolution of $N$--body systems influenced by the tidal field of a parent galaxy and by stellar evolution is presented. The process of energy generation is realized by means of appropriately modified versions of Spitzer's and Mikkola's formulae for the interaction cross section between binaries and field stars and binaries themselves. The results presented are in good agreement with theoretical expectations and the results of other methods. During the evolution, the initial mass function (IMF) changes significantly. The local mass function (LMF) around the half--mass radius closely resembles the actual global mass function (GMF). At the late stages of evolution the mass of the evolved stars inside the core can be as high as $97\\%$ of the total mass in this region. For the whole system, the evolved stars can compose up to $67\\%$ of the total mass. The evolution of cluster anisotropy strongly depends on initial cluster concentration, IMF and the strength of the tidal field. The results presented are the first step in the direction of simulating the evolution of real globular clusters by means of the Monte Carlo method. ",
        "introduction": "Dynamical modeling of real, large stelar systems, like globular clusters or galactic nuclei, and understanding their evolution still is a great challenge both for the theory and hardware/software. Basically, there are two approaches. Direct $N$--body method, which requires extremely large hardware requirements and very sophisticated software, and statistical modeling based on the Fokker--Planck and other approximations, which suffers from the poor understanding of the validity of assumptions.  On the side of $N$--body simulations recent years brought a significant progress in both hardware and software. First, parallel (even massively parallel) computing has opened a route to gain performance at relatively low cost and little technological advancement. Secondly, the already successful GRAPE special purpose computers (Sugimoto et al. 1990, Makino et al. 1997, Makino \\& Taiji 1998, Makino 2002, Makino \\& Hut 2003) have been developed in their present generation (GRAPE--6) and aim at the 100 Tflops speed. NBODY6++ (Spurzem 1999), the successor of Aarseth's NBODY6 code has been ported on massively parallel, general purpose computers (CRAY T3E and PC clusters) and was used for astrophysical problems related to the dissolution of globular clusters (Baumgardt 2001, Baumgardt, Hut \\& Heggie. 2002) and the decay of massive black hole binaries in galactic nuclei after a merger (Milosavljevic \\& Merritt 2001, Hemsendorf, Sigurdsson \\& Spurzem 2002). Despite such progresses in hardware and software there is still impossible to model directly evolution of real globular cluster ($N \\sim 10^6$) and galactic nuclei ($N \\sim 10^9$). Recent work by Baumgardt \\et (2002) and Baumgardt \\& Makino (2002) has pushed the limits of present direct modeling to about $10^5$ using both NBODY6++ and the GRAPE--6 special purpose hardware.  On the side of Fokker--Planck method with finite differences and an anisotropic gaseous model, recent years brought great improvements. Models can be used now to simulate more realistic stellar systems. They can tackle: anisotropy, rotation, a tidal boundary, tidal shocking by galactic disk and bulge, mass spectrum, stellar evolution and dynamical and primordial binaries (Louis \\& Spurzem 1991, Giersz \\& Spurzem 1994, Takahashi 1995, 1996, 1997, Spurzem 1996, Takahashi \\& Portegies Zwart 1998, 2000 hereafter TPZ, Drukier \\et 1999, Einsel \\& Spurzem 1999, Takahashi \\& Lee 2000, Kim et al. 2002, Kim, Lee \\& Spurzem 2004, Fiestas, Spurzem \\& Kim 2005). Unfortunately, the Fokker--Planck approach suffers, among other things, from the uncertainty of differential cross--sections of many processes which are important during cluster evolution. It can not supply detailed information about the formation and movement of all objects present in clusters. Additionally, the anisotropic gaseous model assumes a certain form of heat conductivity and closure relations between the third order moments.  There is an elegant alternative to generate models of star clusters, which can correctly reproduce the stochastic features of real star clusters, but without really integrating all orbits directly as in an $N$--body simulation. They rely on the Fokker--Planck approximation and (hitherto) spherical symmetry, but their data structure is very similar to an $N$--body model. These so--called Monte Carlo models were first presented by H\\'enon (1971, 1975), Spitzer (1975) and later improved by Stod\\'o{\\l}kiewicz (1982, 1985, 1986) and in further work by Giersz (1997, 1998, 2001), and recently by Rasio and his collaborators (Joshi, Rasio \\& Portegies Zwart 2000, Watters, Joshi \\& Rasio 2000, Joshi, Nave \\& Rasio 2001, Fregeau et al. 2003, G\\\"urkan, Freitag \\& Rasio 2004), and Freitag (Freitag 2000, Freitag \\& Benz 2001, 2002). The basic idea is to have pseudo--particles, which orbital parameters are given in a smooth, self--consistent potential. However, their orbital motion is not explicitly followed; to model interactions with other particles like two--body relaxation by distant encounters or strong interactions between binaries and other objects, a position of the particle in its orbit, and further free parameters of the individual encounter, are picked from an appropriate distribution by using random numbers. The Monte Carlo scheme takes full advantage of the established physical knowledge about the secular evolution of (spherical) star clusters as inferred from continuum model simulations. Additionally, it describes in a proper way the graininess of the gravitational field and the stochasticity of real $N$--body systems and provides, in a manner as detailed as in direct $N$--body simulations, information about the movement of any objects in the system. This does not include any additional physical approximations or assumptions which are common in Fokker--Planck and gas models (e.g. for conductivity). Because of this, the Monte Carlo scheme can be regarded as a method which lies between direct $N$--body and Fokker--Planck models and combines most of their advantages. A hybrid variant of the Monte Carlo technique combined with a gaseous model has been proposed by Spurzem \\& Giersz (1996), and applied to systems with a large number of primordial binaries by Giersz \\& Spurzem (2000) and  Giersz \\& Spurzem (2003) and for tidally limited systems by Spurzem \\et (2005). The hybrid method uses a Monte Carlo model for binaries or any other object for which a statistical description, as used by the gaseous model, is not appropriate due to small numbers of objects or unknown analytic cross sections for interaction processes. The method is particularly useful for investigating the evolution of large stellar systems with a realistic fraction of primordial binaries, but could also be used in the future to include, for example, the build up of massive stars and blue stragglers by stellar collisions.  Very detailed observations of globular clusters which are available at present or will become available in future (e.g. the Hubble Space Telescope and the new very large terrestrial telescopes) have extended and will extend our knowledge about their stellar content, internal dynamics and the influence of the environment on cluster evolution (Janes 1991, Djorgovski \\& Meylan 1993, Smith \\& Brodie 1993, Hut \\& Makino 1996, Meylan \\& Heggie 1997). This data covers luminosity functions and derived mass functions, color--magnitude diagrams, population and kinematical analysis, including binaries and compact stellar evolution remnants, detailed two--dimensional proper motion and radial velocity data, and tidal tails spanning over arcs several degrees wide (Koch et al. 2004). A wealth of information on \"peculiar\" objects in globular clusters (blue stragglers, X--ray sources (high-- and low--luminosity), millisecond pulsars, etc.) suggests a very close interplay between stellar evolution, binary evolution and dynamical interactions. This interplay is far from being understood. Moreover, recent observations suggest that the primordial binary fraction in a globular cluster can be as high as $15\\%$ -- $38\\%$ (Rubenstein \\& Bailyn 1997), and possible existence of intermediate--mass black holes in the centers of some globular clusters (Gebhardt \\et 2000). With all these observational data such as King at al. (1998), Piotto \\& Zoccali (1999), Rubenstein \\& Bailyn (1999), Ibata et al. (1999), Piotto et al. (1999), Grillmair et al. (1999), Shara et al. (1998), Odenkirchen et al. (2001), Hansen et al. (2002), Richer et al. (2002) (to mention only a few papers), easily reproducible reliable modelling becomes more important than before. Monte Carlo codes provide all the necessary flexibility to disentangle the mutual interactions between all physical processes which are important during globular cluster evolution and to perform in a reasonable time, detailed simulations of real globular clusters.  The ultimate aim of the project described here is to build a Monte Carlo code (in the framework of the MODEST international collaborations, {\\bf http:/www.manybody.org/modest}), which will be able to simulate the evolution of real globular clusters, as closely as possible. In this paper (the third in the series) the Monte Carlo code (which was discussed in detail in Papers I and II) is used to simulate the evolution of stellar systems, which consist of comparable number of stars and have mass comparable to real globular clusters. The results of simulations will be compared with those of Chernoff \\& Weinberg (1990, hereafter CW), Vesperini \\& Heggie (1997, hereafter VH), Aarseth \\& Heggie (1998, hereafter AH), Baumgardt \\& Makino (2002) and Lamers \\et (2005).  The plan of the paper is as follows. In Section 2, a short description of processes implemented into the Monte Carlo code will be presented. In Section 3, the initial conditions will be discussed and results of the simulation will be shown. And finally, in Section 4 the conclusions will be presented. ",
        "conclusions": "This paper is a continuation of Paper I and II, in which it was shown that the Monte Carlo method is a robust scheme to study, in an effective way, the evolution of very large $N$--body systems. The Monte Carlo method describes in a proper way the graininess of the gravitational field and the stochasticity of real $N$--body systems. It provides, in almost as much detail as $N$--body simulations, information about the movement of any object in the system. In that respect the Monte Carlo scheme can be regarded as a method which lies between direct $N$--body and Fokker--Planck models and combines most advantages of both these methods. This is the first important step in the direction of simulating the evolution of a real globular cluster.  It was shown that the results obtained in this paper are in qualitative agreement with these presented by CW, VH, AH, TPZ, JNR, Paper II, Baumgardt \\& Makino (2002) and Lamers \\et 2005. Particularly good agreement is obtained with VH's $N$--body simulations and what is not surprising, with results of Monte Carlo simulations presented in Paper II. All models survive the phase of rapid mass loss and then undergo core collapse and then subsequent post--collapse expansion (except model W3235) in a manner similar to isolated models. The expansion phase is eventually reversed when tidal limitation becomes important. As in isolated models, mass segregation substantially slows down by the end of core collapse. Mass loss connected with stellar evolution dominates the initial phase of cluster evolution. Then the rate of stellar evolution substantially slows down and escape due to tidal stripping takes over. During this phase of evolution the rate of mass loss is nearly constant, and higher for shallower mass functions. Energy carried away by stellar evolution events dominates the energy loss due to tidal stripping, even though the tidal mass loss is higher. For the first time, because of the large number of particles in simulations ($1,000,000$) which results in substantial reduction of statistical fluctuations, it was possible to compute in an accurate way the evolution of density profiles. The development of power--low density profile is clearly visible and agrees very well with theoretical predictions. The observed power--low index is equal to about $-2.2$. The strongly concentrated model (W7235), shows a modest initial build up of anisotropy in the outer parts of the system. As tidal stripping exposes the inner parts of the system, the anisotropy gradually decreases and eventually becomes slightly negative. Model W5235 from the very beginning develops in the outer parts of the system negative anisotropy. The cluster is not concentrated enough to prevent removal of stars which are preferentially on radial orbits. The negative anisotropy stays negative until the time of cluster disruption, when it becomes slightly positive (during cluster disruption most stars are on radial orbits). Because of mass segregation, and due to evaporation across the tidal boundary, which preferentially removes low mass stars, the mean mass in the cluster increases with time. During the core collapse the rate of increase of the mean mass is highest in the central parts of the system (mass segregation). After the core bounce there is a substantial increase in the mean mass in the middle and outer parts of the system (tidal stripping), and more modest increase in the inner parts of the system, which is mainly connected with binary activity. The fraction of evolved stars is increasing during the cluster evolution. This fraction is larger for systems with shorter relaxation time and stronger influence of the tidal field of a parent galaxy. The mass ratio of evolved stars can be as high as nearly $65\\%$ at the outer parts of the system up to nearly $98\\%$ in the core. At the time of $15$ Gyr these ratios are $90\\%$ and $10\\%$, respectively. The presence of a substantial number of practically invisible stars has very important consequences for the interpretation of observational data and it has to be taken into account when the global globular cluster parameters are drown from the observations. Because of stellar evolution, mass segregation and evaporation of stars the IMF is quickly forgotten and impossible to recover from the observational data. The actual GMF closely resembles the LMF for the middle parts of the system (close to half--mass radius). These results are in an excellent agreement with $N$--body simulations presented by VH.  In order to perform simulations of real globular clusters the description of some processes, already included in the code has to be improved, and several additional physical processes have to be added to the code. Stellar and binary evolution, more accurate treatment for energy generation by binaries, particularly in binary--binary interaction, and proper treatment of the escape process in the presence of a tidal field are still waiting for improvement. One of the population synthesis codes developed by Hurley (Hurley \\et 2000, 2002), Portegies Zwart (Portegies Zwart \\& Verbunt 1966) and Belczy\\'nski (Belczy\\'nski, Kalogera \\& Bulik 2002) can be used to follow more accurately single star and binary evolution. The implementation of techniques used in Aarseth's $NBODY$ codes for direct integrations of a few body subsystems can cure the problem of energy generation in 3--body and 4--body interactions (Giersz \\& Spurzem 2003). The treatment of escapers proposed by Spurzem \\et (2005) can be implemented to solve the problem of escapers in the tidal field (problem which is inherently connected with gaseous, Fokker--Planck and Monte Carlo codes). The tidal shock heating of the cluster due to passages through the Galactic disk, interaction with the bulge, shock--induced relaxation, primordial binaries, physical collisions between single stars and binaries are some of the processes, which are waiting for implementation into the code. The inclusion of all these processes does not pose a fundamental theoretical challenge, but is rather complicated from the technical point of view. The international collaboration called MODEST was setup a few years ago to solve problems with merging all available codes (hydrodynamical, stelar evolution and dynamical) into a code which can deal in detail with the evolution of real globular clusters (see {\\bf http://www.manybody.org/modest}). Inclusion into the Monte Carlo code of as much as possible physical processes will allow to perform detailed comparison between simulations and observed properties of globular clusters, and will also help to understand the conditions of globular cluster formation and explain how peculiar objects observed in clusters can be formed. These types of simulations will also help us to introduce, in a proper way, into future $N$--body simulations all the necessary processes required to simulate the evolution of real globular clusters on a star--by--star basis from their birth to their death. \\bigskip \\bigskip  {\\parindent=0pt {\\bf Acknowledgments} I would like to thank Douglas C. Heggie and Rainer Spurzem for stimulating discussions, comments and suggestions. This work was partly supported by the Polish National Committee for Scientific Research under grant 1 P03D 002 27.}"
    },
    "0512/astro-ph0512430_arXiv.txt": {
        "abstract": "We report the discovery of the first trans-neptunian object, designated 2004~XR190, with a nearly-cirular orbit beyond the 2:1 mean-motion resonance. Fitting an orbit to 23 astrometric observations spread out over 12 months yields an orbit of $a=57.2\\pm0.4$, $e=0.08\\pm0.04$, and $i=46.6^{\\circ}$. All viable orbits have perihelia distances $q>49$~AU. The very high orbital inclination of this extended scattered disk object might be explained by several models, but its existence again points to a large as-yet undiscovered population of transneptunian objects with large orbital perihelia and inclination. ",
        "introduction": "Over the last decade, serious observational effort has gone into detecting Trans-Neptunian Objects (TNOs), but despite increasing resources dedicated to the problem there is still a steady stream of surprises. So little light is reflected from distant TNOs that it is the very inner edge (within 50~AU) of the Kuiper Belt region that dominates detections in observational surveys (those with more than 10 include \\citet{JLC96, TJL01, Larsen01, Glad01, Allen02, Elliot05, Petit05, Allen05}). The majority of Kuiper belt detections come from these flux-limited surveys near the ecliptic plane. This results in a bias against objects which are distant (since reflected flux $\\propto d^{-4}$), on highly-inclined orbits (which spend very little time in the ecliptic plane), or intrinsically rare (like the very largest TNOs). These biases are gradually being overcome, and new dynamical classes within the Kuiper belt are being discovered.  These sub-populations of the Kuiper Belt (see \\citet{Glad05} for a recent review) preserve a record of the dynamical processes which governed the formation of the giant planets. The low-eccentricity `classical belt' appears to decline rapidly at the location of the 2:1 mean-motion resonance \\citep{Allen01, TruBro01} Other mean-motion resonances are occupied both inside and outside the 2:1, although outside the 2:1 objects are only present at high eccentricity (with the possible exception of 2004~XR190). The scattered disk population of TNOs has perihelia $q<40$~AU but large semimajor axes $a$; this is a decaying population that has presumably been flung to large-$a$ via scattering with Neptune \\citep{dl97}. Finally, the extended scattered disk (which eventually merges into the inner Oort Cloud) consists of TNOs on stable orbits, pointing to some process capable of lowering the orbital eccentricities of vast numbers of scattered disk objects \\citep{Glad02,Sedna04}.  In this Letter we report the discovery of a new extended scattered-disk object with a nearly-circular orbit beyond the edge of the classical belt and also with one of the highest orbital inclinations known. ",
        "conclusions": "We have presented the discovery of an unusual TNO, 2004~XR190. With a perihelion at $\\sim50$~AU, and a low-eccentricity orbit, it is the closest to a `distant cold Kuiper belt' object which has been detected so far. However, with an inclination of $47^\\circ$, it has clearly been dynamically perturbed at some point in its lifetime.  A plausible explanation of the origin of 2004~XR190's high inclination and low eccentricity is the action of the Kozai effect during a past residence inside the 5:2 or 8:3 mean-motion resonances of Neptune. If Neptune migrated outwards, dropping the TNO out of resonance, this could aid in freezing the $e$/$i$ combination observed today. The modification of its orbit could also be produced by now-absent bodies (rogue planets or passing stars), but producing all of the features present in the transneptunian region is problematic for all of the above models."
    },
    "0512/astro-ph0512095_arXiv.txt": {
        "abstract": "High Mass X-ray Binary Pulsars (HMXBP), in which the companion star is a source of supersonic stellar wind, provide a laboratory to probe the velocity and density profile of such winds. Here, we have measured the variation of the absorption column density along with other spectral parameters over the binary orbit for two HMXBP in elliptical orbits, as observed with the Rossi X-ray Timing Explorer (RXTE) and the BeppoSAX satellites. A spherically symmetric wind profile was used as a model to compare the observed column density variations. In 4U 1538-52, we find the model corroborating the observations; whereas in GX 301-2, the stellar wind appears to be very clumpy and a smooth symmetric wind model seems to be inadequate in explaining the variation in column density. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512576_arXiv.txt": {
        "abstract": "Understanding the physical nature of the dark energy which appears to drive the accelerated expansion of the unvierse is one of the key problems in physics and cosmology today. This important problem is best studied using a variety of mutually complementary approaches. Daly and Djorgovski (2003, 2004) proposed a model independent approach to determine a number of important physical parameters of the dark energy as functions of redshift  directly from the data.  Here, we expand this method to include the determinations of its potential and kinetic energy as functions of redshift. We show that the dark energy potential and kinetic energy may be written as combinations of the first and second derivatives of the coordinate distance with respect to redshift. We expand the data set to include new supernova measurements, and now use a total of 248 coordinate distances that span the redshift range from zero to 1.79.  First and second derivatives of the coordinate distance are obtained as functions of redshift, and these are combined to determine the potential and kinetic energy  of the dark energy as functions of redshift.  An update on the redshift behavior of the dimensionless expansion rate $E(z)$, the acceleration rate $q(z)$, and the dark energy pressure $p(z)$, energy density $f(z)$, and equation of state $w(z)$ is also presented.   We find that the standard $\\Omega_{0m}  = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ model is in an excellent agreement with the data.  We also show tentative evidence that the Cardassian and Chaplygin gas models in a spatially flat universe do not fit the data as well. ",
        "introduction": "Introduction}   The acceleration of the universe at the present epoch can be studied using a variety of techniques such as type Ia supernovae \\cite{R98, P99, T03, K03, B04, R04, A05}, and studies of cosmic microwave background radiation (e.g., \\cite{S03}) combined with large scale structure studies (e.g. \\cite{US05,AS05}), which indicate that at the present epoch the universe is expanding at an accelerting rate. Studies of powerful FRII radio galaxies also indicate that the universe is accelerting \\cite{GDW00, DG02, PDMR03}. To study this in detail, it is important to determine the redshift behavior of the expansion and acceleration rates of the universe, and the properties and redshift evolution of the driver(s) of these rates.  In this vein, \\cite{dd03,dd04} suggest a model-indepedent method of studying the redshift evolution of the expansion and acceleration rates of the universe, and show that the acceleration of the universe can be written in terms of the first and second derivatives of the coordinate distance with respect to redshift. These derivatives are obtained directly from the data on coordinate distances using a statistically robust numerical technique. The acceleration thus derived requires very few assumptions, relying only upon the assumptions that the universe is homogeneous and isotropic on large scales, and has zero space curvature. Indeed, this approach allows a determination of the acceleration of the universe that is independent of a theory of gravity and of the properties and redshift evolution of the drivers of the expansion and acceleration, such as dark energy, dark matter, or other components. Complementary  model-independent approaches based on an integral rather than a differential technique are discussed by \\cite{HT99, HT01, SRSS00, T02, SSSA03, HS03, WT04, WT05}.   Understanding the properties of the physical driver of the acceleration of the universe is of fundamental importance. The driver, commonly referred to as the dark energy, and its properties are parameterized in quantities such as the potential and kinetic energy, and the pressure, energy density, and equation of state. These quantities can be expressed as combinations of the first and second derivative of the coordinate distance, as discussed by \\cite{dd04}.  Here, we show that the dark energy potential and kinetic energy may be determined as functions of redshift by appropriately combining the first and second derivatives of the coordinate distance to sources at different redshift.  This expansion of the method is described in section \\ref{sec:PK}. The method is applied to an enlarged data sample which includes the 71 new Legacy supernovae coordinate distances, presented in section \\ref{sec:newdata}, and the coordinate distances listed in \\cite{dd04}.  The work presented here on the dark energy potential, $V(z)$, and kinetic, $K(z)$, energy is complementary to the work of \\cite {SRSS00}, \\cite{SVJ05} and \\cite{SLP05}. Saini et al. \\cite{SRSS00} derive equations for $V(z)$ and $K(z)$ and use a fitting function for the luminosity distance to obtain values and uncertainties for the parameters of the fitting function; these are then used to obtain $V(z)$, $K(z)$, and the equation of state parameter $w(z)$. Simon et al. \\cite{SVJ05} consider the reconstruction of the dark energy potential based on an expansion of the dark energy potential in terms of Chebyshev polynomials, while \\cite{SLP05} consider the expansion of the quintessence potential $V$ as a power series of the quintessence field $\\phi$. \\cite{SRSS00}, \\cite{SVJ05} and \\cite{SLP05} find that the results are consistent with those expected if a cosmological constant is driving the acceleration of the universe at the present epoch. ",
        "conclusions": ""
    },
    "0512/astro-ph0512210_arXiv.txt": {
        "abstract": "We present a sample of 33 damped Lyman alpha systems (DLAs) discovered in the Sloan Digital Sky Survey (SDSS) whose absorption redshifts ($z_{\\rm abs}$) are within 6000 \\kms\\ of the QSO's systemic redshift ($z_{\\rm sys}$). Our sample is based on 731  $2.5 < z_{\\rm sys} <$ 4.5 non-broad-absorption-line (non-BAL) QSOs from Data Release 3 (DR3) of the SDSS.  We estimate that our search is $\\approx$100\\% complete for absorbers with \\nhi\\ $\\geq 2 \\times 10^{20}$ \\cm2. The derived number density of DLAs per unit redshift, $n(z)$, within $\\Delta v < 6000$ \\kms\\ is higher (3.5 $\\sigma$ significance) by almost a factor of 2 than that of intervening absorbers observed in the SDSS DR3, i.e. there is evidence for an overdensity of galaxies near the QSOs.  This provides a physical motivation for excluding DLAs at small velocity separations in surveys of intervening `field' DLAs.  In addition, we find that the overdensity of proximate DLAs is independent of the radio-loudness of the QSO, consistent with the environments of radio-loud and radio-quiet QSOs being similar. ",
        "introduction": "\\lya\\ absorption lines in the spectra of high redshift QSOs trace the distribution of neutral hydrogen along the lines of sight to the background sources.  Studying the properties of these intervening absorption systems can provide insights into a range of cosmic structures, ranging from the intergalactic medium to protogalaxies. The absorbers with the highest column densities give rise to lines with characteristic Lorentzian profiles. Specifically, damped \\lya\\ (DLA) systems, defined as having \\nhi\\ $\\ge 2 \\times 10^{20}$ \\cm2, dominate the mass density of HI in the universe and are believed to be associated with the progenitors of present day galaxies. Studies of DLAs have furnished a wealth of information about the density and chemical enrichment of galaxy-scale clouds of neutral hydrogen at early epochs (see for example the reviews by Pettini 2004; Wolfe, Gawiser \\& Prochaska 2005).  Nearly all existing studies of DLAs have been confined to absorbers lying beyond a certain velocity from the systemic redshift of the QSO ($z_{\\rm sys}$).  The velocity cut imposed is somewhat subjective and survey dependent, but is usually 3000 -- 6000 \\kms\\ blueward of $z_{\\rm sys}$. The motivation for excluding  proximate DLAs ($\\Delta v < 6000$ \\kms, PDLAs) is two-fold.  First, studies of the \\lya\\ forest have revealed the existence of a `proximity effect', whereby the high UV flux from the QSO causes additional ionisation of the diffuse HI clouds within a few Mpc of the QSO (e.g. Bajtlik, Duncan \\& Ostriker 1988; Lu, Wolfe \\& Turnshek 1991). Even high column density systems whose interstellar media (ISM) are usually considered to be self-shielded can be affected by proximity to a powerful ionising source such as a QSO.  For example, it has been noted that PDLAs seem to preferentially exhibit \\lya\\ emission superimposed on the \\lya\\ absorption trough (e.g. Ellison et al. 2002 and references therein).  More recently, Adelberger et al. (2005) have detected \\lya\\ fluorescence in a DLA separated by 380 (proper) \\hkpc\\ from a $z = 2.84$ QSO.   Excluding absorbers which may be affected by the QSO's ionising radiation simplifies the calculations of its total hydrogen content and chemical abundances.  Second, studies concerned with the statistical properties of \\textit{intervening} galaxies could be affected by the inclusion of material associated with the QSO.  Intrinsic absorption could, for example, be associated with BAL-like outflows or with the QSO host galaxy.  However, the one $z_{\\rm abs} \\sim z_{\\rm sys}$ PDLA with measured metal abundances is known to have a metallicity $\\sim$ 15\\% of the solar value (Meyer, Welty \\& York 1989; Lu et al. 1996). This value is typical of intervening DLAs but inconsistent with the typically solar or super-solar metallicities of BAL-like outflows and intrinsic systems (Barlow \\& Sargent 1997; Petitjean, Rauch \\& Carswell 1994; D'Odorico et al. 2004).  The PDLAs also lack the strong CIV or NV absorption and complex absorption profiles characteristic of BALs (see also M$\\o$ller, Warren \\& Fynbo 1998). Taken together, this evidence implies that PDLAs are more likely to be related to gas associated with galaxies rather than AGN. The differences in redshift between QSOs and PDLAs are typically too large for the absorbing material to be part of the QSO host galaxy suggesting that PDLAs may have an origin similar to the intervening systems, but are simply located closer to the QSO.  In this paper, we consider whether there may be further justification for excluding PDLAs from surveys by investigating whether they are drawn from the same `random' population as the intervening absorbers. It has been known for more than 30 years that QSOs are often associated with galaxy evolution (e.g. Bahcall, Schmidt \\& Gunn 1969) so it is plausible that PDLAs may represent a subset of early galaxies in a distinct environment.  For example, it is well-established that overdensities of galaxies are often seen in the vicinity of QSOs (e.g. Yee \\& Green 1987; Ellingson, Green \\& Yee 1991; Yamada et al. 1997; Hall \\& Green 1998; Hall, Green \\& Cohen 1998; Sanchez \\& Gonzalez-Serrano 1999; McLure \\& Dunlop 2001; Wold et al. 2001; Finn, Impey \\& Hooper 2001; Clements 2000; Sanchez \\& Gonzalez-Serrano 2002). Our objective is therefore to determine whether DLAs located at velocities $<$ 6000 \\kms\\ from the background QSO may similarly inhabit distinct environments compared with intervening `field' absorbers. As a first approach, we compare the number density of PDLAs to that of traditionally selected $z_{\\rm abs} << z_{sys}$ DLAs. A number density of PDLAs consistent with that of intervening DLAs would imply no distinction in their respective environments. Conversely, an excess of DLAs near the QSO's systemic redshift could be interpreted as probing galaxy overdensities analogous to those found by imaging surveys.  Overdensities of PDLAs would provide an empirically motivated cut-off for intervening DLA searches whose goal is to study an unbiased sample of absorbers.  Since the majority of DLA surveys exclude PDLAs, constructing a statistical sample from the literature is challenging.  A handful of PDLAs have been reported (e.g. M$\\o$ller \\& Warren 1993; Pettini et al. 1995; M$\\o$ller et al. 1998; Ellison et al. 2002) but they have not been generally included in published lists of DLAs. Recently, Ellison et al. (2002) searched for PDLAs in the Complete Optical and Radio Absorption Line System (CORALS)  sample (see Ellison et al. 2001 for full sample definition). Combining their sample with PDLAs discovered in previous surveys of radio loud and radio quiet QSOs (RLQs and RQQs respectively) Ellison et al. (2002) found 4 PDLAs within $\\Delta$v $<$ 3000 \\kms\\ of 96 radio-loud QSOs, and just 1 within $\\Delta$v $<$ 3000 \\kms\\ of 49 radio-quiet QSOs (P\\'{e}roux et al. 2001). Given the redshift path ($\\Delta z$) covered by the RQQ sample, the detection of one PDLA is consistent with the known number density per unit redshift ($n(z)$) of intervening absorbers. However, the detection of four PDLAs associated with RLQs implies that the number density of PDLAs is $>$ 4 times that of the intervening population.  This apparent overdensity around RLQs at $z>2$ is significant only at the 1.5 $\\sigma$ level due to the relatively small number of QSOs, and Ellison et al. (2002) noted that better statistics are required to confirm the overdensity.  With the advent of the Sloan Digital Sky Survey (SDSS), the spectra of large numbers of QSOs ($\\sim$ 10$^5$, eventually) are becoming available. Here we search for PDLAs in the spectra of 731 QSOs with apparent magnitude $i \\sol 20.5$ from the SDSS Data Release 3 (DR3, Abazajian et al. 2005), in order to answer the following two questions:  \\begin{itemize} \\item Is the density of DLAs in the vicinity of QSOs significantly higher than that along the line of sight? \\item Do the environments of radio-quiet and radio-loud QSOs differ? \\end{itemize}  This work improves on that of Ellison et al. (2002) in several ways.  First, our sample is an order of magnitude larger than the CORALS survey.  With these improved statistics we can investigate the number density in a range of velocity bins from $z_{\\rm sys}$. Since the data quality of the SDSS spectra is somewhat heterogeneous, we have carried out completeness tests for our sample; these are described in \\S3.  Finally, we use a Monte Carlo simulation to assess the significance of the measured overdensity of PDLAs relative to intervening systems.  \\medskip  We have adopted the `consensus' cosmology model with $H_0=70$ \\kms\\ Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7, $\\Omega_{M}$ = 0.3.  \\begin{table} \\centering \\caption{Metal-line rest-frame wavelengths} \\begin{tabular}{lc} \\hline Line&Wavelength (\\AA)\\\\ \\hline O I    &1302.17\\\\ Si IV  &1393.76\\\\ Si IV  &1402.77\\\\ Si II  &1526.71\\\\ C IV   &1548.20\\\\ C IV   &1550.78\\\\ Fe II  &1608.45\\\\ Al II  &1670.79\\\\ Al III &1854.72\\\\ \\hline \\end{tabular} \\normalsize \\end{table}   \\small \\begin{table*} \\caption{DLAs within 6000 \\kms\\ of the redshift of SDSS DR3 QSOs} \\begin{tabular}{rrrrrccrrl} \\hline\\hline \\multicolumn{1}{c}{RA}&\\multicolumn{1}{c}{Dec}&\\multicolumn{1}{c}{$i$}&\\multicolumn{1}{c}{z$_{\\rm sys}$}&\\multicolumn{1}{c}{S:N}&$S_{1.4GHz}$ mJy&\\nhi&\\multicolumn{1}{c}{z$_{\\rm abs}$}&\\multicolumn{1}{c}{$\\Delta v$}&\\multicolumn{1}{c}{Metal lines}\\\\ \\multicolumn{1}{c}{J2000}&\\multicolumn{1}{c}{J2000}&&&&(and log$_{10} R$)&$10^{20}$ \\cm2&&\\multicolumn{1}{c}{\\kms}&\\\\ \\multicolumn{1}{c}{(1)}&\\multicolumn{1}{c}{(2)}&\\multicolumn{1}{c}{(3)}&\\multicolumn{1}{c}{(4)}&\\multicolumn{1}{c}{(5)}&(6)&(7)&\\multicolumn{1}{c}{(8)}&\\multicolumn{1}{c}{(9)}&\\multicolumn{1}{c}{(10)}\\\\ \\hline \\rule[0mm]{0mm}{4mm} 01 40 49.18&$-$08 39 42.5&17.7&3.713& 7.5&    -&4.5 &3.696&1085&OI SiII [BAL]                           \\\\ \\rule[0mm]{0mm}{4mm} 07 39 38.85&+30 59 51.2&20.4&3.399& 9.2& 6.69 (2.9)&2.0 &3.353&3160&SiII CIV                 \\\\ \\rule[0mm]{0mm}{4mm} 08 05 53.02&+30 29 37.3&19.9&3.432& 9.9&    -&2.5 &3.429&200 &OI SiII CIV                  \\\\ \\rule[0mm]{0mm}{4mm} 08 11 14.32&+39 36 33.2&20.1&3.092& 6.9&    -&7.0 &3.037&4060&                        \\\\ \\rule[0mm]{0mm}{4mm} 08 26 38.59&+51 52 33.2&17.2&2.850& 5.9&    -&5.5 &2.833&1330&OI SiIV  SiII CIV  FeII AlII      \\\\ \\rule[0mm]{0mm}{4mm} 08 44 51.72&+05 18 27.8&20.0&4.465& 7.2&    -&3.5 &4.376&4925&                        \\\\ \\rule[0mm]{0mm}{4mm} 09 09 30.42&+07 00 50.7&20.3&3.273& 5.5&    -&3.0 &3.219&3815&                        \\\\ \\rule[0mm]{0mm}{4mm} 09 40 35.93&+50 03 08.7&20.1&3.567& 5.8&    -&2.5 &3.500&4435&                        \\\\ \\rule[0mm]{0mm}{4mm} 10 11 22.59&+47 00 42.2&19.0&2.928& 9.1&29.56 (2.5)&2.5 &2.909&1490&OI SiII           \\\\ \\rule[0mm]{0mm}{4mm} 10 26 19.09&+61 36 28.9&18.6&3.849& 5.3&    -&3.0 &3.785&3985&OI SiIV  SiII CIV  AlII       \\\\ \\rule[0mm]{0mm}{4mm} 10 45 01.44&+50 40 45.9&20.0&3.998& 5.1&    -&6.0 &3.949&2955&                        \\\\ \\rule[0mm]{0mm}{4mm} 10 48 20.93&+50 32 54.2&20.2&3.889&12.0&    -&15.0&3.815&4545&OI SiIV  CIV              \\\\ \\rule[0mm]{0mm}{4mm} 11 12 24.18&+00 46 30.3&19.7&4.035&14.0&    -&19.0&3.958&4620&OI AlII                     \\\\ \\rule[0mm]{0mm}{4mm} 11 26 48.62&+05 56 28.1&19.6&3.191&17.1&    -&7.0 &3.147&3165&                        \\\\ \\rule[0mm]{0mm}{4mm} 12 27 23.65&+01 48 06.0&20.4&2.881&10.1&    -&2.0 &2.876&425 &                        \\\\ \\rule[0mm]{0mm}{4mm} 12 42 04.27&+62 57 12.1&19.6&3.321& 5.7& 0.88 (1.4)&2.0 &3.276&3140&SiIV SiII CIV  FeII AlII         \\\\ \\rule[0mm]{0mm}{4mm} 12 48 30.64&+49 14 00.2&20.4&3.079& 7.7&    -&4.0 &3.032&3460&                        \\\\ \\rule[0mm]{0mm}{4mm} 12 57 59.22&$-$01 11 30.2&18.7&4.112& 8.1&    -&2.0 &4.022&5330&OI SiIV  CIV              \\\\ \\rule[0mm]{0mm}{4mm} 13 25 21.27&+52 45 13.1&19.8&3.953& 6.3&    -&2.5 &3.861&5625&OI                      \\\\ \\rule[0mm]{0mm}{4mm} 13 41 00.13&+58 07 24.2&19.4&3.500& 7.3&    -&7.0 &3.417&5585&all those listed in Table 1                   \\\\ \\rule[0mm]{0mm}{4mm} 13 46 37.94&+56 49 15.6&19.7&3.463& 6.7&    -&3.0 &3.430&2225&OI                      \\\\ \\rule[0mm]{0mm}{4mm} 14 00 29.01&+41 12 43.4&19.0&2.540& 5.9&    -&45.0&2.537&255 &OI SiII CIV  FeII AlII           \\\\ \\rule[0mm]{0mm}{4mm} 14 20 41.96&+42 22 57.0&19.6&3.378& 8.2&    -&2.0 &3.336&2925&OI                      \\\\ \\rule[0mm]{0mm}{4mm} 14 51 20.65&+39 13 50.4&20.4&3.376& 8.3&    -&2.0 &3.307&4770&SiII CIV                 \\\\ \\rule[0mm]{0mm}{4mm} 15 12 54.37&$-$00 56 36.6&20.2&4.457& 5.6&    -&8.0 &4.400&3150&OI SiIV  SiII CIV  FeII       \\\\ \\rule[0mm]{0mm}{4mm} 15 52 48.01&+56 03 28.9&20.0&3.604&12.4&    -&2.0 &3.556&3145&                        \\\\ \\rule[0mm]{0mm}{4mm} 16 08 13.86&+37 47 27.3&20.3&3.996& 8.5&    -&2.0 &3.952&2655&OI SiIV                  \\\\ \\rule[0mm]{0mm}{4mm} 16 11 19.56&+44 11 44.0&20.0&4.024& 5.1&    -&5.0 &3.982&2520&SiIV                     \\\\ \\rule[0mm]{0mm}{4mm} 17 04 21.85&+62 47 41.5&20.1&2.980& 6.0&    -&6.0 &2.953&2040&OI SiIV                     \\\\ \\rule[0mm]{0mm}{4mm} 17 18 00.20&+62 13 25.6&19.8&3.670& 7.8&    -&3.5 &3.618&3360&OI SiII FeII AlII               \\\\ \\rule[0mm]{0mm}{4mm} 17 20 07.20&+60 28 23.8&20.4&4.425& 8.6& 5.06 (2.2)&5.0 &4.326&5525&                        \\\\ \\rule[0mm]{0mm}{4mm} 21 00 25.03&$-$06 41 46.0&18.1&3.138& 9.0&    -&9.0 &3.092&3350&all those listed in Table 1                \\\\ \\rule[0mm]{0mm}{4mm} 21 22 07.36&$-$00 14 45.7&19.0&4.072& 5.2&    -&2.5 &4.001&4230&OI CIV                   \\\\ \\hline \\end{tabular}  \\normalsize The columns give: (1-4) SDSS RA, Dec, i mag and emission-line redshift; (5) S:N in the continuum redward of the \\lya\\ emission; (6) FIRST 1.4-GHz radio flux density (and log$_{10} R$ where R is the radio-loudness parameter (see \\S2); log$_{10} R >$ 1.0 for a RLQ); (7) fitted H I column density of the DLA; (8) DLA redshift; (9) DLA velocity, relative to $z_{\\rm sys}$; (10) metal lines associated with the DLA  \\end{table*}   \\normalsize ",
        "conclusions": "We have searched the spectra of 731 SDSS DR3 2.5 $< z_{\\rm sys} <$ 4.5 QSOs for DLAs (\\nhi\\ $> 2 \\times 10^{20}$ \\cm2) and found 33 within $\\Delta v <$ 6000 \\kms\\ (and 12 within $\\Delta v <$ 3000 \\kms). Our search increases the number of PDLAs reported in the literature within $\\Delta  v <$ 3000 \\kms\\ (the conventional definition of $z_{\\rm abs} \\sim z_{\\rm sys}$ DLAs) from 5 to 16.   \\begin{figure} \\centering \\psfig{file=fig5.ps,width=8.5cm,height=6.17cm,angle=270} \\caption{\\label{v_hist} $n(z)_{DLA}$ found in SDSS as a function of $\\Delta v$.  The bins are 1000 \\kms\\ wide.  The dashed line indicates the expected number density for intervening DLAs $n(z)_{DLA}$.} \\end{figure}  The number density of DLAs within $\\Delta v <$ 6000 \\kms\\ of the redshift of the QSOs is approximately double that observed for samples of intervening DLAs (a) found in the same sample, the SDSS (Prochaska et al. 2005) and (b) for the relation derived by Storrie-Lombardi \\& Wolfe (2000).  The excess is significant at the 3.5 $\\sigma$ level. There is no significant difference between the number density of PDLAs towards RLQs and RQQs. A similar excess of PDLAs is found within $\\Delta v <$ 3000 \\kms\\ of QSOs; $n(z)_{PDLA}$ / $n(z)_{DLA}$ = 1.7 for all RQQs and RLQs combined.  We speculate that  the overdensities of DLAs in the vicinity of QSOs may trace galaxies in the same clusters or superclusters as the QSOs. The similar excesses around radio-quiet and radio-loud quasars are consistent with their inhabiting similar environments, as suggested by recent observations (e.g. Finn, Impey \\& Hooper 2001; Wold et al. 2001).  This emphasises the importance of imposing a velocity cut when studying intervening DLAs, not only to avoid intrinsic AGN absorption, and any proximity effect, but also to avoid objects in the cluster or supercluster hosting the QSO.  Moreover, we recommend that, based on Fig. \\ref{v_hist}, $\\Delta v$ $\\sim$ 6000 \\kms\\ is a suitable velocity cut.  A number of PDLAs in the DR3 exhibit strong metal lines, despite the moderate resolution and S:N ratios of the spectra. With high-resolution spectroscopy, it should be possible to determine whether the metallicities of PDLAs are enhanced compared with intervening DLAs, as may be expected if they are associated with the QSO environment.  For example, Ellison \\& Lopez (2001) and Lopez \\& Ellison (2003) found unusual relative abundances in a pair and a triplet of DLAs whose velocity separations were $<10$ $000$ \\kms.  Many QSOs in the SDSS are bright enough to follow-up with an echelle spectrograph on an 8-m telescope, such as UVES on the VLT.  Investigating the abundances of PDLAs will be an interesting, and so far unexplored, extension of this work.  \\vspace{5mm} $Acknowledgements$. We are grateful to Jason X. Prochaska for providing the statistics for the intervening DLAs discovered in Data Release 3 prior to publication, and to Patrick B. Hall for valuable advice on the SDSS database. Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0512/astro-ph0512026_arXiv.txt": {
        "abstract": "{This work presents the results from a systematic search for evidence of temporal changes (i.e., non-stationarity)  associated with spectral variations in \\object{3C~390.3}, using data from a two-year intensive RXTE monitoring campaign of this broad-line radio galaxy. In order to exploit the potential information contained in a time series more efficiently, we adopt a multi-technique approach, making use of linear and non-linear techniques. All the methods show suggestive evidences for non-stationarity in the temporal properties of 3C~390.3 between 1999 and 2000, in the sense that the characteristic time-scale of variability decreases as the energy spectrum of the source softens. However, only the non-linear, ``scaling index method\" is able to show conclusively that the temporal characteristics of the source do vary, although the physical interpretation of this result is not clear at the moment. Our results indicate that the variability properties of 3C~390.3 may vary with time, in the same way as they do in Galactic black holes in the hard state, strengthening the analogy between the X-ray variability properties of the two types of object. This is the first time that such a behavior is detected in an AGN X-ray light curve. Further work is needed in order to investigate whether this is a common behavior in AGN, just like in the Galactic binaries, or not. ",
        "introduction": "Among the basic properties characterizing an Active Galactic Nucleus (AGN) the flux variability, especially at X-rays, is one of the most common.  Previous multi-wavelength variability studies (e.g., Edelson et al. 1996; Nandra et al. 1998) have shown that AGN are variable in every observable wave band, but the X-ray flux exhibits variability of larger amplitude and on time-scales shorter than any other energy band, indicating that the X-ray emission originates in the innermost regions of the central engine.  However, even though the discovery of X-ray variability dates back more than two decades, its origin is still poorly understood. Nevertheless, it is widely acknowledged that the variability properties of AGN are an important means of probing the physical conditions in their emission regions.  The reason is that using energy spectra alone it is often impossible to discriminate between competing physical models and thus the complementary information obtained from the temporal analysis is crucial to break the spectral degeneracy.  Over the years, several variability models have been proposed, involving one or a combination of the fundamental components of an AGN: accretion disk, corona, and relativistic jet. These models can be divided into two main categories: 1) intrinsically linear models, as the shot noise model (e.g., Terrel 1972), where the light curve is the result of the superposition of similar shots or flares produced by many independent active regions, and 2) non-linear models, which require some coupling between emitting regions triggering avalanche effects, as the self-organized criticality model (e.g., Mineshige et al.1994), or models assuming that the variability is caused by variations in the accretion rate propagating inwards (e.g. Lyubarskii 1997).  The ``rms-flux relation'' recently discovered by Uttley \\& McHardy (2001), along with detection of non-linearity in different AGN light curves (i.e., in 3C~345 by Vio et al. 1992; in 3C~390.3 by Leighly \\& O'Brien 1997; in NGC~4051 by Green et al. 1999; in Ark~564 by Gliozzi et al. 2002), favors non-linear variability models, but there is no general consensus on the nature of the X-ray variability yet.  Because of their brightness, the temporal and spectral properties of Galactic black holes (GBHs) are much better known and can be used to infer information on their more powerful, extragalactic analogs, the AGN. It is well established that GBHs undergo state transitions (see McClintock \\& Remillard 2003 for a recent review), switching between ``low'', ``intermediate\", ``high'', and ``very high\" states, which are all unambiguously characterized by a specific combination of energy spectrum and power spectral density (PSD). Importantly, even when in the same state, the timing properties of GBHs may vary with time, often accompanying changes of the spectral properties. For example, Pottschmidt et al. (2003) showed that changes in the PSD of Cyg X-1 are associated with changes in the slope of the energy spectrum. Specifically, during the ``low/hard state'', the PSD characteristic frequencies increase, as the spectral slope becomes steeper.  \\begin{figure*}[th] \\psfig{figure=f1.ps,height=10cm,width=16.5cm,% bbllx=32pt,bblly=6pt,bburx=530pt,bbury=290pt,angle=0,clip=} \\caption{X-ray light curves of the background-subtracted count rate in the 2--20 keV band (top panel) and of the X-ray color 7--20 keV/2--5 keV (bottom panel) from RXTE PCA observations of 3C 390.3.  Time bins are 5760s ($\\sim$ 1 RXTE orbit). Filled circles represent data points from the first monitoring campaign from 1999 January 8 to 2000 February 29, whereas open diamonds are data points from the second monitoring campaign from 2000 March 3 to 2001 February 23. The dashed lines in the top panel indicate the intervals used for the time-resolved spectral analysis (see $\\S$9).} \\label{figure:lc} \\end{figure*}  Transitions between ``different'' states have not been observed yet in AGN. The reason may be associated with the fact that their flux variability is characterized by much longer time-scales. Nevertheless, a correspondence between GBH spectral states and some AGN classes has been hypothesized. Specifically, Broad-Line Seyfert 1 galaxies have been associated with the GBH low/hard state, and Narrow-Line Seyfert 1 galaxies with the GBH high/soft state, (for a recent discussion on this topic see, for example, McHardy et al. 2004).  Although genuine spectral transitions in AGN are unlikely to be detected due to the longer time-scales involved, it is perhaps possible to detect changes in the temporal properties (i.e., loss of stationarity) associated with spectral changes, similarly to what observed in GBHs within the same spectral state. To this end, of crucial importance are the numerous monitoring campaigns carried out by \\rxte\\ in the last decade, which have produced data sets suitable for searching for non-stationarity in the AGN light curves, where long-term spectral changes proved to be ubiquitous.  In this work, we present the results from a systematic search for evidence of temporal changes associated with significant spectral variations in the case of the broad-line radio galaxy (BLRG) \\object{3C~390.3}. We use data from a two-year intensive \\rxte\\ monitoring campaign.  Previous X-rays studies with \\rosat, and \\rxte\\ (Leighly \\& O'Brien 1997, Gliozzi et al. 2003) demonstrated that \\object{3C~390.3} is one of the most variable AGN on time-scales of days and months. Our motivation stems from the results of a long monitoring campaign carried out with \\rosat\\ HRI, which showed evidence for possible non-stationarity (and non-linearity) in the soft X-ray light curve of \\object{3C~390.3} (Leighly \\& O'Brien 1997). ",
        "conclusions": "The main results and conclusions of this work can be summarized as follows: \\begin{itemize}  \\item We have studied in detail the variability characteristics of the two-year long, \\rxte\\ monitoring light curve of \\3c. Linear methods, like PSD, SF and PDF analysis show suggestive evidence of a change in the characteristic time-scale of the system between 1999 and 2000.  \\item However, only the results from the scaling index method show a clear and significant evidence of a change in the intrinsic properties of the system that cannot be explained by red or Poisson noise effects.  \\item It is not easy to interpret the SIM results at this moment. To the extend that the scaling index is related with the timing properties of the variability process,  one possibility is a  change of the characteristic time-scale between 1999 and 2000, in agreement with the suggestive results from the PSD, SF and PDF analyses.  \\item Our results reinforce the similarity between the X-ray variability properties of AGN and GBHs. In analogy with GBH phenomenology, the loss of stationarity in 3C~390.3 seems to correspond to an increase in the frequency break as the spectrum softens, as observed in Cyg X-1, when the system is in its low/hard spectral state.  \\item  Since the SIM is not yet a standard technique, further work is needed in order to investigate its properties and its possible relation with more frequently used (and better understood) statistical analysis tools like the power spectrum. Nevertheless, we believe that the present results demonstrate the benefits from the application of the SIM (and possibly other non-linear methods) in the analysis of the AGN X-ray light curves.  \\item  The \\rxte\\ archive is populated with numerous, long AGN light curves. This is a good opportunity to start a systematic study of these data, for a search of spectrally related variations of the timing properties of these systems. Since the outcome of any Monte Carlo simulation based on an assumed ``model-PSD'' is inevitably dependent on the assumptions made, we believe that a thorough analysis of a large number of high-quality AGN light curves is the most direct way to examine whether the present results are indicative of the AGN behavior in general, or just a peculiarity associated with \\3c. We plan to explore this issue in a future work.  \\end{itemize}"
    },
    "0512/astro-ph0512356_arXiv.txt": {
        "abstract": "We present VLT/ISAAC near-infrared imaging of the host galaxies of 15 low luminosity quasars at 1 $<$ z $<$ 2. This work complements our studies to trace the cosmological evolution of the host galaxies of high luminosity quasars. The radio-loud (RLQ) and radio-quiet (RQQ) quasars have similar distribution of redshift and luminosity, and together the high and low luminosity quasars cover a large range of the quasar luminosity function. Both RLQ and RQQ hosts resemble massive inactive ellipticals undergoing passive evolution. However, RLQ hosts are systematically more luminous than RQQ hosts, as also found for the high luminosity quasars. The difference in the host luminosity remains the same from z = 2 to z = 0. For the entire set of quasars, we find a correlation between the nuclear and the host luminosities, albeit with a large scatter. The correlation is less apparent for the RQQs than for the RLQs. ",
        "introduction": "\\label{intro}  Low redshift (z $\\leq$ 0.5) quasars are predominantly hosted by massive, bulge-dominated galaxies \\citep{bahcall97,dunlop03,pagani03}. This is consistent with the fact that nearby massive spheroids host inactive supermassive black holes (BH) \\citep{ferrarese02}, and suggests that episodic quasar activity may be common in galaxies and that the nuclear power depends on the mass of the galaxy \\citep{kauffmann03}.  At low redshift, the BH mass is related to the luminosity and velocity dispersion of the bulge \\citep{marconi03,bettoni03,haring04}. Furthermore, the strong cosmological evolution of quasars is similar to the BH mass accretion rate and the evolution of the cosmic star formation history \\citep{madau98,barger01,yu02}. Therefore, determining the properties of quasar hosts close to the peak of quasar activity is crucial to investigate the fundamental link between the formation and evolution of massive galaxies and the nuclear activity.  The detection of the host galaxies of high redshift quasars is very challenging because the host galaxy rapidly becomes very faint compared to the nucleus. To cope with this, high spatial resolution and S/N, and a well defined PSF are needed. All these requirements can be fulfilled by ground-based 8m class telescopes. We recently carried out a systematic VLT/ISAAC imaging study of 17 quasars (10 radio-loud quasars [RLQ] and seven radio-quiet quasars [RQQ]) at 1 $<$ z $<$ 2 to characterize their host galaxies \\citep{falomo04}. The evolution of both RLQ and RQQ hosts until z $\\sim$2 is consistent with that of massive ellipticals undergoing passive evolution. There is no significant decrease in the host mass as would be expected in hierarchical formation models \\citep{kauffmann00}. RLQ hosts are more luminous by $\\sim$0.6 mag than RQQ hosts at all redshifts.  The quasars in \\citep{falomo04} belong to the bright end of quasar luminosity function. Here we present imaging of quasars that have on average lower luminosity by $\\sim$2 mag, to study the dependence of host properties on nuclear luminosity. We use H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.3 and $\\Omega_\\Lambda$ = 0.7. ",
        "conclusions": ""
    },
    "0512/astro-ph0512483_arXiv.txt": {
        "abstract": "We present the results of 37 nights of CCD unfiltered photometry of nova V2574 Oph (2004) from 2004 and 2005. We find two periods of 0.14164 d $\\approx$ 3.40 h and 0.14773 d $\\approx$ 3.55 h in the 2005 data. The 2004 data show variability on a similar timescale, but no coherent periodicity was found.  We suggest that the longer periodicity is the orbital period of the underlying binary system and that the shorter period represents a negative superhump. The 3.40 h period is about 4$\\%$ shorter than the orbital period and obeys the relation between superhump period deficit and binary period. The detection of superhumps in the light curve is evidence of the presence of a precessing accretion disk in this binary system shortly after the nova outburst. From the maximum magnitude -- rate of decline relation, we estimate the decay rate $t_2$ = 17 $\\pm$ 4 d and a maximum absolute visual magnitude of $M_{Vmax}$ = --7.7 $\\pm$ 1.7 mag. ",
        "introduction": "Novae are a subclass of cataclysmic variables (CVs) which contain a white dwarf and a solar-like companion that fills its Roche Lobe. Typically, the white dwarf accretes mass from an accretion disk that surrounds it. After of order ten thousand years, the white dwarf accumulates enough material for a thermonuclear-runaway event, which results in the observed nova outburst (e.g., Warner 1995). It is believed that the nova eruption disrupts the accretion disk, which is reformed several weeks or months later (e.g., Retter 2004).  Nova V2574 Oph ($\\alpha_{2000.0}=17^h38^m45^s, \\delta_{2000.0}= -23^o28'18''.5$) was discovered by Akira Takao (Kitakyushu, Japan) at V $\\approx$ 10.2 on 2004 Apr. 14.80 UT (Yamaoka 2004). Mason et al. (2004) reported that echelle spectra (390 -- 900nm; resolution 48000) of V2574 Oph, obtained on 2004 Apr. 17.32 and 18.37 UT at La Silla with the 2.2-m telescope, were dominated by H$\\alpha$, Na D, and Ca II emission lines, which are flanked by double P-Cyg profiles. From the minima of the P-Cyg absorptions, the measured expansion velocities were estimated as 400 and 1050 km s$^{-1}$ for H$\\alpha$ and 400 and 1000 km s$^{-1}$ for Na D.  Rudy et al. (2004) observed V2574 Oph 73 days after maximum and found that the nova was still in an `Oxygen I' phase where the O I lines at 0.8446 and 1.1287 microns that are fluorescently excited by Lyman $\\beta$ have strengths comparable to H$\\alpha$. The optical region showed numerous multiplets of Fe II which characterize this class as novae. The low expansion velocities of the ejecta mentioned above seem to rule out the possibility that V2574 Oph is a hybrid nova, which evolved to the He/N class from the Fe II type (Della Valle \\& Livio 1998). The overall appearance of the spectra indicated a substantial evolution with respect to early reports and suggested that this is a slow nova caught at maximum light (Bond \\& Walter 2004).  So far, there are about 50 novae with known orbital periods (Warner 2002). Typical nova periods range from about 2 to 9 h. Finding the orbital period of a nova yields an estimate of the secondary mass (e.g., Smith \\& Dhillon 1998). In addition, detecting several periodicities in novae can help in classifying the system into different groups of CVs such as magnetic systems, intermediate polars, and/or permanent superhump systems (e.g., Diaz \\& Steiner 1989; Baptista et al. 1993; Retter, Leibowitz \\& Ofek 1997; Patterson et al. 1997; Skillman et al. 1997; Patterson \\& Warner 1998; Retter \\& Leibowitz 1998; Retter, Leibowitz \\& Kovo-Kariti 1998; Patterson 1999; Retter, Leibowitz \\& Naylor 1999; Skillman et al 1999; Retter \\& Naylor 2000; Patterson 2001; Woudt \\& Warner 2001; Lipkin et al. 2001; Patterson et al. 2002; Warner 2002; Woudt \\& Warner 2002; Retter et al. 2003; Ak, Retter \\& Liu 2005, Balman et al. 2005; see also Retter et al. 2002; Retter, Richards \\& Wu 2005). This yields valuable information on the magnetic field of the white dwarf and reveals the presence or absence of the accretion disk.  We have an ongoing program to observe novae with small telescopes to search for periodicities in their optical light curves. In this paper, we present extensive photometric observation of V2574 Oph, which suggests the presence of orbital and superhump periods.   \\begin{table}\\caption{The observations time table} \\begin{center} \\begin{tabular}{ccccc}  UT       & Time of Start    & Run Time  & Points  & Comments      \t\\\\ (yymmdd) & (HJD --2453000)  &  (Hours)  & number                        \\\\ \\\\ 040531   & 157.07290        & 7.8       &   219     \t&  \t        \\\\ 040622   & 179.00080        & 8.5       &   235\t        & \t\t\\\\ 040623   & 180.03790        & 7.4       &   205\t\t& \t\t\\\\ 040804   & 222.00030        & 5.2       &   386\t\t& \t\t\\\\ 040805   & 223.01050        & 5.1       &   371\t\t& \t\t\\\\ 040807   & 225.00060        & 5.2       &   286\t\t& \t\t\\\\ 040808   & 226.03010        & 4.0       &   177\t\t&\t\t\\\\ 040809   & 227.01290        & 5.0       &   219\t\t& \t\t\\\\ 040810   & 227.99230        & 4.9       &   221\t\t& \t\t\\\\ 040812   & 229.98880        & 5.3       &   181\t\t& \t\t\\\\ 040813   & 231.00440        & 4.9       &   199\t\t& \t\t\\\\ 040814   & 232.01670        & 5.0       &   203\t\t& \t\t\\\\ 040815   & 233.01580        & 4.9       &   199\t\t& \t\t\\\\ 040816   & 234.00970        & 5.0       &   206\t\t& \t\t\\\\ 040817   & 235.01410        & 4.6       &   189\t\t& \t\t\\\\ 040818   & 236.03860        & 2.8       &   127\t        & \t\t\\\\ 040820   & 238.00240        & 4.5       &   191\t\t& \t\t\\\\ 040822   & 240.06770        & 2.9       &   129\t\t& \t\t\\\\ 040826   & 244.07070        & 3.1       &   142\t\t& \t\t\\\\ 040827   & 244.98720        & 4.2       &   190\t\t& \t\t\\\\ 040828   & 245.96630        & 4.7       &   209\t\t& \t\t\\\\ 040829   & 247.01970        & 3.8       &   161\t\t& \t\t\\\\ 040830   & 248.01820        & 3.9       &   152\t        & \t        \\\\ 050507   & 498.14124 \t    & 6.3       &   176\t\t& \t\t\\\\ 050508   & 499.14860        & 6.2       &   166\t\t& \t\t\\\\ 050509   & 500.14133  \t    & 6.4       &   177\t\t& \t\t\\\\ 050510   & 501.14419  \t    & 6.1       &   170\t\t& \t\t\\\\ 050511   & 502.13795  \t    & 6.4       &   176\t\t& \t\t\\\\ 050512   & 503.14005  \t    & 6.4       &   175\t\t& \t\t\\\\ 050513   & 504.14247  \t    & 6.5       &   176\t\t& \t\t\\\\ 050514   & 505.14854  \t    & 6.3       &   169\t\t& \t\t\\\\ 050516   & 507.13301  \t    & 6.3       &   177\t\t& \t\t\\\\ 050607   & 529.08677        & 7.1       &   176         & clouds\t\\\\ 050610   & 532.08986        & 6.9       &   165         & clouds\t\\\\ 050611   & 533.08452        & 4.9       &   109 \t& clouds\t\\\\ 050613   & 535.09411        & 6.6       &   187    \t& \t\t\\\\ 050614   & 536.06776        & 7.9       &   223    \t& \t\t\\\\ \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "(1) By using the maximum visual magnitude of V = 9.5 $\\pm$ 0.2 we measured the decay time $t_2$ = 17 $\\pm$ 4 d from the long-term light curve. This makes V2574 Oph a fast nova. We estimated a visual absolute magnitude in maximum of $M_{Vmax}$ = --7.7 $\\pm$ 1.7 mag.  (2) We found two periods of 0.14164 d $\\approx$ 3.40 h and 0.14773 d $\\approx$ 3.55 h in the 2005 data.  (3) We interpret the longer 3.55 h period as the orbital period of the binary system and the 3.40 h period as a negative superhump period.  (4) More observations (including radial velocity studies) are required to confirm our results and to follow the evolution of the periodicities in time."
    },
    "0512/astro-ph0512489_arXiv.txt": {
        "abstract": "The Gamma-Ray Burst (GRB) prompt emission is believed to be from highly relativistic electrons accelerated in relativistic shocks. From the GRB high-energy power-law spectral indices $\\beta$ observed by the Burst and Transient Source Experiment (BATSE) Large Area Detectors (LAD), we determine the spectral index, $p$, of the electrons' energy distribution. Both the theoretical calculations and numerical simulations of the particle acceleration in relativistic shocks show that $p$ has a universal value $\\approx 2.2-2.3$. We show that the observed distribution of $p$ during GRBs is not consistent with a $\\delta$-function distribution or an universal $p$ value, with the width of the distribution $\\ge$ 0.54. The distributions of $p$ during X-ray afterglows are also investigated and found to be inconsistent with a $\\delta$-function distribution. The $p$-distributions in blazars and pulsar wind nebulae are also broad, inconsistent with a $\\delta$-function distribution. ",
        "introduction": "GRBs are observed to have non-thermal spectra during its prompt emission phase (Band et al. 1993). It is widely believed that the synchrotron radiation and/or the inverse Compton scattering are the likely emission mechanism(s) for GRB's prompt hard X-ray and $\\gamma$ ray emission. The electrons accounting for these emissions are thought to be accelerated in relativistic shocks in GRBs. According to the shock diffusive acceleration model, particles are accelerated when they repeatedly cross a shock front, and the competition between the particle's energy again and escape probability per shock crossing cycle leads to a power-law spectrum for the particles: \\begin{equation} N(\\gamma) d\\gamma \\propto \\gamma^{-p} d\\gamma\\,\\,,  \\end{equation} where $\\gamma$ is the Lorentz factor of the particle (e.g., Blandford \\& Ostriker 1978). For non-relativistic shocks, the value of $p$ depends on the the compression ratio of the flow stream across the shock; while in relativistic or ultra-relativistic shocks, which is most likely the case in GRBs, analytical and numerical studies show that $p$ has an  ``universal'' value, $\\approx 2.2-2.3$ (Kirk et al. 2000, Achterberg et al. 2001, Bednarz \\& Ostrowski 1998, Lemonine \\& Pelletier 2003).  The work reported in this paper is to investigate the ``universality'' of the power-law index $p$ for GRBs, which we calculate directly from the high-energy ($0.1 - 2$ MeV) photon spectrum of GRBs (Preece et al. 1998, 2000), assuming the spectrum is from synchrotron or synchrotron self-inverse-Compton emission of the power-law distributed highly relativistic electrons, using the relations between $p$ and the spectral index , $\\beta$, of the high-energy power-law photon spectrum.  In \\S2, we describe the GRB spectral data set used and the process of determining the parent $p$-distribution. In \\S3, we examine the contributions from the spectral fit procedure and the time averaging effect to the dispersion of the parent distribution of $p$. The $p$-distributions derived from BeppoSAX GRBs and from HETE-2 GRBs, X-ray flashes and X-ray rich GRBs are presented in \\S4. We determine the $p$-distributions for X-ray afterglows in \\S5 and for blazars and pulsar wind nebulae in \\S6. A summary and discussions are given in \\S7. ",
        "conclusions": "Motivated by theoretical calculations and numerical simulations showing that the shock-accelerated electrons in relativistic shocks have a power-law distribution with an universal index $p \\simeq 2.2-2.3$, we have determined the values of $p$ from $\\gamma$-ray and X-ray spectra  for a number of relativistic sources such as GRBs (prompt emissions and afterglows), blazars and pulsar wind nebulae.  The maximum likelihood estimate of the width of the parent distribution for GRB prompt emission is found to be quite broad, $\\sigma_p=0.51\\pm 0.02$; the probability that the distribution is consistent with a $\\delta$-function is extremely small, and therefore this result does not support that there is an universal $p$.  We have considered the systematic errors in photon index due to spectra fit and time averaging of spectra and their contributions to the scatter in $p$ distribution. We have shown that those contributions are very small for GRBs and can not explain the scatter in $p$ distribution.  For X-ray afterglows of GRBs, the $p$-distribution of the BeppoSAX sample can not rule out a possibility that the parent distribution is a $\\delta$-function distribution; however, a larger sample of Swift afterglows is inconsistent with a $\\delta$-function parent distribution. We point out that the smaller width of parent distribution for the BeppoSAX sample is due to its larger measurement error in $\\beta$.  Analysis of 44 blazar spectra and 9 pulsar wind nebulae shows that the distributions of $p$ for blazars and pulsar wind nebulae (PWNe) are also broad, not consistent with a $\\delta$-function distribution.  Possible situations in which the ``universality'' of $p$ could break are: (i) The shock is Mildly relativistic (cf. Kirk et al. 2000); (ii) The magnetic field is oblique to the shock normal (Baring 2005); (iii) The nature and strength of the downstream magnetic turbulence are varying (Ostrowski \\& Bednarz 2002, Niemiec \\& Ostrowski 2004). A non-Fermi acceleration in a collisionless plasma shock was studied by Hededal et al. (2004), in which electrons are accelerated and decelerated instantaneously and locally, by the electric and magnetic fields of the current channels formed through the Weibel two-stream instability. It is not known whether an ``universality'' of $p$ could hold for this mechanism. The ``universality'' of $p$ might not happen in non-shock accelerations; for instance, in an alternative model for GRBs (Lyutikov \\& Blandford 2003), the energy is carried outward via magnetic field or Poynting flux. The particles accounting for the $\\gamma$-ray emissions are accelerated by magnetic field reconnection which may also produce a power-law spectra of accelerated particles with a variable $p$ (however, this is still poorly understood)."
    },
    "0512/astro-ph0512506_arXiv.txt": {
        "abstract": "We discuss an extended set of Tree+SPH simulations of the formation of clusters of galaxies, with the goal of investigating the interplay between numerical resolution effects and star--formation/feedback processes.  Our simulations were all carried out in a concordance $\\Lambda$CDM cosmology and include radiative cooling, star formation, and energy feedback from galactic winds.  The simulated clusters span the mass range $\\mvir \\simeq (0.1-2.3)\\times 10^{15}\\msun$, with mass resolution varying by several decades. At the highest achieved resolution, the mass of gas particles is $m_{\\rm gas}\\simeq 1.5\\times 10^7\\msun$, which allows us to resolve the virial region of a Virgo--like cluster with more than 2 million gas particles and with at least as many dark--matter (DM) particles.  Our resolution study confirms that, in the absence of an efficient feedback mechanism, runaway cooling leads to about 35 per cent of baryons in clusters to be locked up in long lived stars at our highest resolution, with no evidence of convergence.  However, including feedback causes the fraction of cooled baryons to converge at about 15 per cent already at modest resolution, which is much closer to the typical values inferred from observational data.  Feedback also stabilizes other gas--related quantities, such as radial profiles of entropy, gas density and temperature, against variations due to changes in resolution. Besides effects of mass resolution, we also investigate the influence of the gravitational force softening length, and that of numerical heating of the gas induced by two-body encounters between DM and lighter gas particles. We also show that simulations where more DM than gas particles are used, such that $m_{\\rm gas}\\simeq m_{\\rm DM}$, show a significantly enhanced efficiency of star formation at $z\\magcir 3$, but they accurately reproduce at $z=0$ the fraction of cooled gas and the thermodynamical properties of the intra--cluster gas.  Our results are important for establishing and delineating the regime of numerical reliability of the present generation of hydrodynamical simulations of galaxy clusters. ",
        "introduction": "\\label{sec:intro}  Within the hierarchy of cosmic structures, clusters of galaxies mark an interesting transition between the large--scale regime, where the dynamics is dominated by gravity, and the small--scale regime, where complex astrophysical processes, such as gas cooling, star formation and energetic feedback processes, control the formation and evolution of galaxies \\citep[e.g.][ for recent reviews]{2002ARA&A..40..539R,2005RvMP...77..207V}. If gravity were the only player at the scales relevant for galaxy clusters, then the overall thermal content of baryons would be completely determined by the processes of adiabatic compression and shock heating within the dark--matter (DM) dominated potential wells \\citep[e.g.][]{1986MNRAS.222..323K,2001ApJ...546...63T}. Since gravity has no characteristic scales, this scenario predicts that galaxy clusters and groups of different mass should appear as scaled versions of each other, with only weak residual trends due to a variation of the halo concentration with mass. However, a number of observational facts highlight that the evolution of cosmic baryons within clusters must be influenced by physical processes other than gravity.  First, scaling relations between X--ray observables demonstrate that groups and poor clusters have a lower content of diffuse hot gas than rich clusters, and that they lie on a comparatively higher adiabat \\citep[e.g.][]{1998ApJ...504...27M,1999MNRAS.305..631A,2002A&A...391..841E,2003MNRAS.343..331P,2004MNRAS.350.1511O}. Second, the bulk of the gas in central cluster regions lies at a temperature which is never observed to fall below 1/3--1/2 of the virial temperature, despite the fact that the cooling time of this gas is much shorter than the Hubble time \\citep[e.g.][]{2001A&A...365L.104P,2001ApJ...560..194M,2002A&A...382..804B}. Third, the fraction of baryons in the stellar phase within clusters is consistently rather small, around $\\sim 10$ per cent \\citep[e.g.][]{2001MNRAS.326.1228B}, with a tendency for poor clusters and groups to have a somewhat higher amount of stars \\citep[e.g.][]{2003ApJ...591..749L,2004AJ....128.1078R}.  These observational evidences indicate the existence of a delicate interplay between the physical processes which determine the evolution of the inter--galactic (IGM) and intra--cluster (ICM) media. Including hydrodynamical effects beyond ordinary gas dynamics, such as gas cooling, star formation and associated feedback processes, in simulation models is however met with substantial difficulty.  While analytical and semi--analytical calculations can provide very useful guidelines for the expected effects \\citep[e.g.][]{2001ApJ...546...63T,2001MNRAS.325..497B,2002MNRAS.330..329B,2003ApJ...593..272V,2005ApJ...619...60L}, we nevertheless rely on such simulation models, because the non-linearity of the physics and the geometrical complexity of cluster dynamics can usually only be captured in full by direct simulation.  This makes it extremely important to obtain an in-depth understanding of the robustness of the numerical methods employed, of the assumptions behind them, and of the sensitivity of numerical predictions for observable quantities on the adopted numerical parameters and approximations.  One important validation test is to compare different simulation codes and different schemes for solving hydrodynamics with each other. For a single code, additional control tests that scrutinize the robustness of the predictions against resolution and different implementations of sub--resolution processes are highly desirable as well. An example for the importance of such tests is represented by radiative cooling, which is well known to exhibit a runaway character: cooling leads to an increase of the gas density, which, in turn, increases the cooling efficiency. Consistently, hydrodynamical cosmological simulations have shown that the fraction of gas that cools down and becomes available for star formation is much larger than indicated by observations \\citep[e.g.][]{1993ApJ...412..455K,1998ApJ...507...16S,2000MNRAS.317.1029P,2000ApJ...536..623L,2002MNRAS.335..762Y}, a tendency that increases with resolution \\citep[e.g.][]{2001MNRAS.326.1228B,2003MNRAS.342.1025T}. Solving this problem apparently requires the introduction of suitable feedback physics which reduces the fraction of cooled gas and stabilizes its value against numerical resolution.  So far, detailed comparisons between different cosmological hydrodynamical codes have been restricted to the case of non--radiative simulations \\citep[e.g.][]{1994ApJ...430...83K,1999ApJ...525..554F,2005ApJS..160....1O}. While these comparisons have shown a reasonable level of agreement, some sizeable differences have been found in the profiles of thermodynamic quantities (e.g., entropy) when simulations of galaxy clusters performed with Eulerian and Lagrangian codes were compared \\citep{1999ApJ...525..554F}.  A limited number of studies have been presented so far which were aimed at discussing the effect of resolution and different implementations of cooling/star--formation physics within a single code. For instance, \\cite{2002MNRAS.330..113K} compared different implementations of star formation. They concluded that different prescriptions provide broadly consistent results, although the amount of cooled gas is always much higher than observed. They also verified that this overcooling can be partly ameliorated by resorting to either kinetic or thermal feedback.  Besides checking the influence of different parameterizations of physical effects, it is crucial to have purely numerical effects well under control, such as those originating from mass resolution and force softening.  Detailed studies have been presented about the determination of an optimum softening for purely collisionless simulations \\citep[e.g.][]{2003MNRAS.338...14P,2005astro.ph..7237Z}, usually based on a compromise between the desire of obtaining an accurate estimate of the accelerations together with high spatial resolution, and at the same time, the need to suppress two-body relaxation and the formation of bound particle pairs if too small a softening is chosen \\citep[e.g.][]{1992MNRAS.257...11T}.  The situation is more complicated for hydrodynamical simulations, where energy can be spuriously transferred from the collisionless to the collisional component, thereby affecting the evolution of the gas.  As \\cite{1997MNRAS.288..545S} have shown, this can happen when the gas particles are substantially lighter than the dark matter particles, such that the former receive a systematic energy transfer in two--body encounters, leading to artificial heating of the gas.  As a result, a lower--limit exists for the required mass resolution in order to provide a correct description of gas cooling within halos of a given mass.    In a series of papers \\citep[e.g.,][]{2004MNRAS.348.1078B,2004ApJ...606L..97D,2004ApJ...607L..83M,2004MNRAS.354..111E,2005MNRAS.356.1477D,2005A&A...431..405C,2005astro.ph..9024E} we have presented results on the properties of galaxy clusters, extracted from a large--scale simulation and performed with the {\\small GADGET-2} code \\citep{SP01.1,2005astro.ph..5010S}, including the star formation and feedback model of \\cite{2003MNRAS.339..289S,2003MNRAS.339..312S}. This model is formulated as a sub--resolution model to account for the multiphase nature of the interstellar medium (ISM), and contains a phenomenological model of energy feedback from galactic winds triggered by supernova (SN) explosions (see Section \\ref{sec:sims}). The present paper is specifically aimed at discussing the numerical robustness of the results we obtained, by analysing an extended set of re--simulations of galaxy clusters, spanning a fairly large range both in cluster mass and numerical resolution.  More specifically, we will primarily address the following two questions: {\\em (a)} How does numerical resolution affect the properties of the diffuse baryons and the distribution of star formation within clusters?  {\\em (b)} How large is the impact of artificial heating and how can numerical parameters (i.e., force softening and mass--ratio between gas and DM particles) be chosen to minimize its effect?  The plan of the paper is as follows. In Section~2, we provide the details of the simulated clusters. We analyze in Section~3 the effects of changing resolution over a fairly large range, up to a factor 45 in particle masses. In this section we will also discuss how the adopted feedback from galactic winds affects the resolution dependence of measured properties for the simulated clusters. In Section~4, we discuss the role of different sources of numerical heating. Finally, we will summarize our results and draw our main conclusions in Section~5. ",
        "conclusions": "\\label{sec:conc}  In this paper we have presented results from a large set of hydrodynamical simulations of galaxy clusters, carried out with the Tree+SPH code {\\small GADGET-2}. Our simulations include radiative cooling, star formation and energy feedback by a phenomenological model for galactic winds.  The main target of our analysis has been the study of the stability of simulation results with respect to numerical parameters, such as mass resolution or gravitatiobal softening length. We also considered different sources of numerical heating, and their interplay with the complex physical effects included. As such, our analysis also represents a validation study of our previous results \\citep[e.g.,][]{2004MNRAS.348.1078B}, which were based on a large cosmological box simulated a relatively low resolution.  Our simulated clusters span more than one order of magnitude in collapsed mass and several decades in mass resolution. At the highest resolution, the mass of the gas particles is $m_{\\rm gas}\\simeq 1.5\\times 10^7 h^{-1}{\\rm M}_\\odot$, which allows us to resolve the virial region of a Virgo--like cluster with more than 2 million gas particles and at least as many dark--matter (DM) particles. Our main results are concerned with the effects of resolution on the properties of the stellar populations and on the intra--cluster medium of the simulated galaxy clusters. They can be summarized as follows.  \\begin{description} \\item[(a)] In the absence of an efficient energy feedback, the fraction of cooled baryons steadely increases with resolution, reaching $\\simeq 35$ per cent at the highest achieved resolution (VR runs in Table \\ref{tab:res}), with no indication for convergence. \\item[(b)] Including feedback from galactic winds has the effect of stabilizing the stellar fraction inside clusters. Assuming a high efficiency for the supernova driven winds of order unity, we find that the fraction of cooled baryons converges already at a relatively modest resolution, with an indication to decrease very slightly at the highest resolution. This arises as a consequence of the self--regulation property of star formation and feedback. While improving the resolution increases the star formation efficiency at very high redshift, this at the same time also provides a significant contribution to gas pre--heating which reduces star formation later on.  The fraction of cooled baryons within \\rvir~ lies in the range 12-18 per cent, with a decreasing trend with cluster mass. These values increase by about 15 per cent when the normalization of the power--spectrum is raised from $\\sigma_8=0.8$ to 0.9. \\item[(c)] The feedback provides the necessary continued energy supply to keep gas particles of comparately low entropy and short cooling times in the hot phase of clusters, while without feedback these particles would cool down and drop out of the cluster atmosphere. As a result, the central gas density is higher in runs with feedback than in the runs with no galactic winds. The temperature profiles are shallower in the central cluster regions, while isentropic cores are much less pronounced, thus alleviating the discrepancy with the observed properties of the intra--cluster medium. \\end{description}  A further series of tests presented in this paper concerns the effect of numerical heating. The main results from these tests can be summarized as follows.  \\begin{description} \\item[(a)] Our Plummer--equivalent force softening of $\\epsilon_{\\rm Pl}\\simeq 10\\hk$ at a mass resolution of $m_{\\rm DM}\\simeq 5\\times 10^9\\msun$, scaled to other particle masses as $\\epsilon_{\\rm Pl}\\propto m_{\\rm DM}^{1/3}$, represents a reasonable compromise between the need of preventing a spurious numerical heating of the gas, and the desire to resolve galaxies in the largest possible number of small DM halos. \\item[(b)] {Increasing the mass resolution in the DM component while keeping the gas mass resolution fixed,} helps to reduce numerical gas heating by two--body encounters \\citep{1997MNRAS.288..545S}. We find that decreasing $m_{\\rm DM}$ has a negligible effect on the fraction of cooled gas at $z=0$ and on the radial profiles of gas properties. However, it has a non--negligible effect on the star formation history, which occurs with enhanced efficiency at $z\\magcir 3$. This is the result both of a non--negligible numerical heating in small halos when gas and DM particles have rather different masses, and also of a better resolved dark matter content in the first generation of halos. \\end{description}  In sum, our analysis helps to establish and delineate the regime of numerical reliability of the present generation of hydrodynamical SPH simulations of galaxy clusters. Even though quite technical in nature, it is clear that systematic tests like the ones discussed here are required to elucidate the often delicate interplay betweeen numerical effects and the physical model that is studied. Only when these effects are understood, the predictive power of numerical experiments can be fully exploited.  At the same time, this understanding is also required to successfully push to new generations of simulations with yet higher resolution. A reassuring aspect of the results discussed here is that even simulation models with complex and highly non-linear physical models for star formation and feedback can produce surprisingly robust results. This clearly is an encouragement for attempts to improve the fidelity with which the physics is represented in future simulation work."
    },
    "0512/astro-ph0512056_arXiv.txt": {
        "abstract": "We present a simple method, based on the deformation of spherically symmetric potentials, to construct explicit axisymmetric and triaxial MOND density-potential pairs. General guidelines to the choice of suitable deformations, so that the resulting density distribution is nowhere negative, are presented. This flexible method offers for the first time the possibility to study the MOND gravitational field for sufficiently general and realistic density distributions without resorting to sophisticated numerical codes.  The technique is illustrated by constructing the MOND density-potential pair for a triaxial galaxy model that, in the absence of deformation, reduces to the Hernquist sphere.  Such analytical solutions are also relevant to test and validate numerical codes. Here we present a new numerical potential solver designed to solve the MOND field equation for arbitrary density distributions: the code is tested with excellent results against the analytic MOND triaxial Hernquist model and the MOND razor-thin Kuzmin disk, and a simple application is finally presented. ",
        "introduction": "Milgrom (1983) proposed that the failure of galactic rotation curves to decline in Keplerian fashion outside the galaxies' luminous body arises not because galaxies are embedded in massive dark halos, but because Newton's law of gravity has to be modified for fields that generate accelerations smaller than some characteristic value $\\az$. Subsequently, in order to solve basic problems presented by this phenomenological formulation of the theory (now known as Modified Newtonian Dynamics or MOND), such as conservation of linear momentum (Felten 1984), Bekenstein \\& Milgrom (1984) substituted the heuristic 1983 model with the MOND non-relativistic field equation \\begin{equation} \\nabla\\cdot\\left[\\mu\\left({\\Vert\\nabla\\phi\\Vert\\over\\az}\\right) \\nabla\\phi\\right] = 4\\pi G \\rho, \\label{eqMOND} \\end{equation} where $\\Vert ...\\Vert$ is the standard Euclidean norm, $\\phi$ is the gravitational potential produced by the density distribution $\\rho$, and $\\nabla\\phi\\to 0$ for $\\Vert\\xv\\Vert\\to\\infty$.  As stressed by Bekenstein \\& Milgrom (1984), the equation above is obtained from a variational principle applied to a Lagrangian with all the required symmetries, so the standard conservation laws are obeyed.  Thus, equation~(\\ref{eqMOND}) plays in MOND the same role as the Poisson equation \\begin{equation} \\nabla^2\\phiN=4\\pi G\\rho \\label{eqPoisson} \\end{equation} in Newtonian gravity, and the MOND gravitational field $\\gv$ experienced by a {\\it test} particle is \\begin{equation} \\gv=-\\nabla\\phi. \\label{eqgv} \\end{equation}  As well known, in the regime of intermediate accelerations the function $\\mu$ is not fully constrained by theory or observations, while in the asymptotic regimes \\begin{equation} \\mu(t)\\sim\\cases{t&for $t\\ll 1$,\\cr 1&for $t\\gg 1$. } \\label{eqmuasym} \\end{equation} Throughout the paper we conform to the standard assumption \\begin{equation} \\mu (t)={t\\over\\sqrt{1+t^2}} \\label{eqmu} \\end{equation} (see however Famaey \\& Binney 2005).  From equation~(\\ref{eqmuasym}) it follows that equation~(\\ref{eqMOND}) reduces to the Poisson equation when $\\Vert \\nabla \\phi \\Vert \\gg \\az$, while the limit equation \\begin{equation} \\nabla\\cdot\\left({\\Vert\\nabla\\phi\\Vert}\\nabla\\phi\\right) = 4\\pi G \\az \\rho, \\label{eqdMOND} \\end{equation} obtained by assuming $\\mu(t)=t$ in equation~(\\ref{eqMOND}), describes systems for which (or regions of space where) $\\Vert\\nabla\\phi\\Vert \\ll\\az$, i.e. systems for which the MOND predictions differ most from the Newtonian ones.  As a consequence equation~(\\ref{eqdMOND}), characterizing the so-called `deep MOND regime' (hereafter dMOND), is of particular relevance in MOND investigations.  Nowadays a considerable body of observational data seems to support MOND well beyond its originally intended field of application (see, e.g., Milgrom~2002; Sanders \\& McGaugh~2002), making this theory an interesting alternative to the Cold Dark Matter paradigm. It is thus natural to study in detail MOND predictions, in particular focusing on dMOND systems, i.e. systems that should be dark matter dominated if Newtonian gravity holds.  Potential problems of the theory have already been pointed out by various authors (see, e.g. The \\& White 1988; Buote et al.~2002; Sanders~2003; Ciotti \\& Binney~2004; Knebe \\& Gibson~2004; Zhao et al.~2005), but further analyses are needed to reach firmer conclusions.  Unfortunately, MOND investigations have been considerably slowed down by the lack of aspherical density-potential pairs to test theory predictions in cases more realistic than those described by spherical symmetry: the search for a sufficiently general method to construct aspherical MOND solutions is the subject of this paper.  In MOND, the main difficulty to obtain exact aspherical density-potential pairs (or to build robust numerical solvers) originates from the non-linear nature of the theory, which makes impossible a straightforward use of the analytical and numerical techniques available for the Poisson equation (such as integral transforms or expansion in orthogonal functions).  In addition, a simple relation between the Newtonian and the MOND gravity fields in general does not exist. Although equation~(\\ref{eqPoisson}) can be used to lower the order of equation~(\\ref{eqMOND}) by eliminating the source density, it follows that $\\mu(\\Vert\\nabla\\phi\\Vert/\\az)\\nabla\\phi$ and $\\nabla\\phiN$ differ by some unknown solenoidal field\\footnote{That the solenoidal field $\\Sv$ in general cannot be arbitrarily set to zero is due to the fact that the MOND acceleration field must be derived from a potential, while equation~(\\ref{eqmug}), if correct in general, would imply that $\\nabla \\Vert \\gv \\Vert \\wedge \\gv=0$, an identity which is not necessarily true when $\\gv$ is derived from a potential.} $\\Sv={\\rm curl\\,}\\hv$. Remarkably, Brada \\& Milgrom (1995; hereafter BM95) showed that when the modulus $\\gN$ of the Newtonian gravitational field produced by $\\rho$ is a function of $\\phiN$ only, $\\Sv$ vanishes, and so the MOND acceleration $\\gv$ is related to $\\gvN\\equiv-\\nabla\\phiN$ by \\begin{equation} \\mu\\left({g\\over\\az}\\right)\\gv=\\gvN, \\label{eqmug} \\end{equation} where $g=\\Vert\\gv\\Vert$.  Equation~(\\ref{eqmug}) coincides with the original MOND formulation of Milgrom (1983) and can be solved algebraically: particularly simple cases described by this equations are those in which the density distribution is spherically or cylindrically symmetric, or stratified on homogeneous planes.  In such cases the MOND potential is not more difficult to construct than the corresponding Newtonian potential.  This is the reason why {\\it all} the (astrophysically relevant) analytical MOND density-potential pairs known are spherically symmetric: in fact, the only exact, aspherical MOND density-potential pair available is the axisymmetric razor-thin Kuzmin disk (and the derived family; BM95), for which $\\gN=\\gN(\\phiN)$.  In all the other cases one is forced to solve numerically equation~(\\ref{eqMOND}).  The importance of a general method to obtain the explicit MOND potential of density distributions with prescribed shape and stratification is then obvious, because it would allow for orbit integration, without resorting to numerical integration of equation~(\\ref{eqMOND}); in addition, analytical solutions could be used to test numerical MOND solvers in more realistic cases than spherical symmetry.  In this paper we show that such an approach can be devised. In particular, we show how a ``seed'' spherical density distribution can be deformed in an axisymmetric or triaxial density distribution with analytical potential satisfying the MOND equation, by means of a pair of very simple existence theorems and (for example) by using building blocks obtained from the Ciotti \\& Bertin (2005; hereafter CB05) method. In addition, as a complement to the analytical method, we also illustrate a new numerical MOND solver based on spectral methods, which adds to the short list of the others available (Milgrom 1986; BM95; see also Brada \\& Milgrom~1999).  The paper is organized as follows. In Section~2 we present the general method (postponing to the Appendix the proof of three technical results on which the method is based), and in Section~3 we construct families of triaxial profiles obtained by deformation of the Hernquist (1990) sphere. The original numerical code developed to compute the MOND potential of generic density distributions is then described in Section~4, where we also show how well it recovers the analytical triaxial potentials of Section~3 and the MOND Kuzmin density-potential pair. The code is finally used to investigate the field $\\Sv$ of highly flattened triaxial density distributions.  The main results of the paper and possible future applications are finally summarized in Section~5. ",
        "conclusions": "In this paper we presented a few representative axisymmetric and triaxial density models, both analytical and numerical, obeying the MOND field equation. The analytical density-potential pairs are constructed by means of a simple method based on the deformation of the potential of spherically symmetric systems. We show that in this way it is possible to build systems with radial profiles and shapes similar to those of real galaxies. Our method, although applied in the present paper by using simple deformations is easily generalizable to more complicated cases.  As a complement to the analytical method, we also presented a numerical code (based on spectral methods) to solve the MOND field equation. We tested the code against the new analytical triaxial models and the MOND Kuzmin disk, with excellent results in terms of both computational time and accuracy of the numerical solution.  As a first simple application of the code, we explored the relevance of the solenoidal field $\\Sv$ in MOND distributions. Though we confirm that this field is typically small compared to the Newtonian field of the density distribution, we also found that for some systems $\\Sv$ is certainly not negligible, at least in some regions of space.  A particularly important application of the numerical solver will be its implementation in a particle-mesh N-body code. Such an implementation is promising, because our MOND solver is based on an iterative scheme, and at each time in a N-body simulation the potential at the previous time is a very good seed for the iterative procedure that should converge efficiently.  More immediate possible applications of the presented analytical models and numerical code are the study of MOND orbits in aspherical density distributions, and the vertical motions of stars near the galactic plane.  We conclude by pointing out that the developed numerical code can also be used to solve the equation for the scalar field in the non-relativistic limit of TeVeS, the relativistic MOND theory introduced by Bekenstein~(2004; equation~53), allowing, for example, to investigate MOND gravitational lensing from non-spherical lenses.   We are grateful to James~Binney, Hongsheng~Zhao, and the anonymous referee for very useful comments.  This work was partially supported by a MIUR grant CoFin2004.    \\appendix"
    },
    "0512/gr-qc0512157_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\  Scalar-tensor theories of gravity have been studied by many theoretical physicists as natural alternatives to general relativity since the pioneering work by Brans and Dicke\\cite{1}. In particular, a theory of dilaton gravity has been of interest as an effective theory of the superstring theories at low energy scales\\cite{2}. In these theories, the gravity is mediated not only by a tensor field but also by a scalar field. Several theoretical predictions in the scalar-tensor theories have been obtained$,^{3)-7)}$ and it has been found that a wide class of the scalar-tensor theories can pass all the experimental tests in the cases of weak gravitational fields. In the cases of strong gravitational fields, however, it has been found that the scalar-tensor theories show different aspects of the gravity in contrast to general relativity. For example, it has been shown numerically that nonperturbative effects in the scalar-tensor theories increase the maximum mass of an isolated system such as a neutron star$.^{3)-5)}$ In these works, numerical methods play an important role, and a spherical exterior solution, which is a static, spherically symmetric exterior solution outside a spherical body, is matched to the numerical solution in interpreting its results. A usual astronomical object, however, has a non-spherical form and is in rotation, and not only a spherical exterior solution but also axisymmetric exterior solutions may therefore play an important role in discussing the crucial difference between general relativity and scalar-tensor theories of gravity.  Among the axisymmetric exterior solutions, ones corresponding to the series of solutions of Tomimatsu and Sato in general relativity\\cite{8} must be of significant interest. It is, however, difficult to obtain their explicit forms, and the case is similar even for solutions corresponding to Kerr's solution in general relativity. Instead of studying such the Kerr-like solutions, Tsuchida and Watanabe\\cite{7} have investigated the Weyl-like solutions, namely the solutions which are reduced to those of Weyl's series in the case of a constant scalar field. Among these Weyl-like solutions, we have a special solution, which we shall refer to as the scalar-tensor-Weyl solution. This solution is very notable for the reason that it has asymptotic flatness properties and moreover is reduced to Voorhees's solution\\cite{9}, namely a solution of Weyl's series of prolate solutions, in the case of a constant scalar field. Several geometrical properties of the scalar-tensor-Weyl solution, especially the properties of light propagation, have been investigated by analytical approaches in Ref.7. On the basis of this work, we will study light propagation and gravitational lensing on the scalar-tensor-Weyl solution by numerical approaches because the local properties of null geodesics must tell us important geometrical information of the spacetime.  In this paper, we first study the global structure of the scalar-tensor-Weyl solution. We obtain an asymptotic form of the solution near the spatial infinity in order to clarify the physical significance of three model parameters found in the solution. It may be interesting that the Schwarzschild-like coordinates can be naturally introduced in doing that. It will be shown that these parameters are related to an amplitude of the scalar field, a non-sphericality of the spacetime symmetry and a mass-like parameter in the Einstein frame. The directional-dependent properties of the spacetime singularity will be also shown. After giving the several analytical results as briefly summarized in Abstract, we numerically solve null geodesic equations and Sachs's optical scalar equations\\cite{10}. By using a technique of the conformal transformation, we then study deflection and shear of light rays on the equatorial plane in the manner such that the obtained results are independent of details of scalar-tensor theories of gravity. We numerically obtain the deflection angle and the image distortion rate as functions of the impact parameter and find a close relationship between their qualitative properties and the classification of the model parameter space according to the Weyl source-term. Several analytic results shown in the previous papers are summarized in Appendices\\cite{7,11}.  Throughout this paper, we use the system of units such that $c=G=1$.  \\ ",
        "conclusions": ""
    },
    "0512/astro-ph0512110_arXiv.txt": {
        "abstract": "We continue our numerical analysis of the morphological and energetic influence of massive stars on their ambient interstellar medium for a 35 $\\Msun$ star that evolves from the main sequence through red supergiant and Wolf-Rayet phases, until it ultimately explodes as a supernova. We find that structure formation in the circumstellar gas during the early main-sequence evolution occurs as in the 60 $\\Msun$ case but is much less pronounced because of the lower mechanical wind luminosity of the star. Since on the other hand the shell-like structure of the \\HII region is largely preserved, effects that rely on this symmetry become more important. At the end of the stellar lifetime 1\\% of the energy released as Lyman continuum radiation and stellar wind has been transferred to the circumstellar gas. From this fraction 10\\% is kinetic energy of bulk motion, 36\\% is thermal energy, and the remaining 54\\% is ionization energy of hydrogen. The sweeping up of the slow red supergiant wind by the fast Wolf-Rayet wind produces remarkable morphological structures and emission signatures, which are compared with existing observations of the Wolf-Rayet bubble S308, whose central star has probably evolved in a manner very similar to our model star. Our model reproduces the correct order of magnitude of observed X-ray luminosity, the temperature of the emitting plasma as well as the limb brightening of the intensity profile. This is remarkable, because current analytical and numerical models of Wolf-Rayet bubbles fail to consistently explain these features. A key result is that almost the entire X-ray emission in this stage comes from the shell of red supergiant wind swept up by the shocked Wolf-Rayet wind rather than from the shocked Wolf-Rayet wind itself as hitherto assumed and modeled. This offers a possible solution to what is called the ``missing wind problem'' of Wolf-Rayet bubbles. ",
        "introduction": "\\label{sec_intro}  In the first paper of this series \\citep*[subsequently referred to as Paper I]{freyer03} we studied the evolution of the circumstellar gas around an isolated 60 $\\Msun$ star by means of numerical two-dimensional radiation hydrodynamic simulations. We found that the interaction of the photoionized \\HII region with the stellar wind bubble (SWB) strongly influences the morphological evolution during the early main-sequence (MS) phase of the star. On the one hand, the results show that the dynamical interaction processes contribute to the formation of complex structures which can be found in \\HII regions. On the other hand, these processes also impact on how and to what extent the stellar energy input (wind and H-ionizing radiation) is supplied to the interstellar medium (ISM), distributed among different forms of energy, and ultimately radiated from the system. While the consideration of the stellar wind strongly enhances the kinetic energy of bulk motion present in the system, ionization energy and the associated thermal energy of warm gas are generally lowered because the stellar wind intensifies the formation of high-density structures (clumps) with shorter hydrogen recombination times and stronger cooling.  With this paper we continue our numerical analysis of the morphological and energetic influence of massive stars on their ambient ISM for a 35 $\\Msun$ star that evolves from the MS through the red supergiant (RSG) and the Wolf-Rayet (W-R) phases until it explodes as a supernova (SN). The goals of this paper are to examine the combined influence of wind and ionizing radiation on the dynamical evolution of circumstellar matter for this second set of stellar parameters, to compare and contrast the two numerical models we have obtained so far and to compare the model with observations of bubbles around W-R stars that have undergone an RSG phase. We will complete our little sequence of circumstellar gas models for a 85 $\\Msun$ star (D. Kr\\\"oger et al. 2005, in preparation) and a 15 $\\Msun$ star, representing the upper and lower end, respectively, of the stellar mass range that we are investigating.  The remainder of this paper has the following structure: In section \\ref{sec_observations} we review some recent observations of W-R stars with circumstellar nebulae that are conjectured to be ejected during the RSG phase of the star. Section \\ref{sec_numerics} briefly reflects on the numerical methods and initial conditions used to produce the results presented in this paper and describes the set of stellar parameters used as time-dependent boundary conditions that drive the evolution. The results of the model calculations along with a comparison to observations and analytical models are presented in section \\ref{sec_results}. We summarize our main results and conclusions in section \\ref{sec_conclusions}. ",
        "conclusions": "\\label{sec_conclusions}  The basic difference between the simulations presented in this paper and those of \\citeauthor{freyer03} is the choice of the stellar parameters, which here are chosen to model a 35 $\\Msun$ star that undergoes the evolution from the MS through the RSG and W-R stages. The fundamental structures which evolve are basically the same as observed in the 60 $\\Msun$ case. They are generally smaller because for most of the time the stellar wind luminosity and the Lyman continuum luminosity are lower than in the 60 $\\Msun$ case. At the end of the simulation the entire bubble structure has a radius of $\\approx$44 pc, which is some 6 pc smaller than the final bubble of the 60 $\\Msun$ case, although the 35 $\\Msun$ star lives 0.88 Myr longer than the 60 $\\Msun$ star.  Instability-driven structure formation during the early MS phase of the star (the formation of ionized fingers corrugating the ionization front and the formation of neutral spokes shadowed by dense clumps), which we found to be quite prominent in the case of the 60 $\\Msun$ calculation, is also visible in the calculations presented here, but it is much less pronounced and only short-lived, because of the lower mechanical wind luminosity of the star. The lower mechanical wind luminosity of the star reduces the thermal pressure of the hot gas in the bubble. The lower pressure in the hot bubble increases the geometrical thickness of the shell of swept-up \\HII gas, making it less sensitive to thin-shell instabilities that could trigger the formation of the morphological structures described above. Since this behavior better preserves the basic spherical structure, other morphological effects become visible, which might be prohibited in the 60 $\\Msun$ case by the strong corrugation of the bubble shell: When the swept-up \\HII shell broadens geometrically, the plasma density in the shell decreases. This rarefaction reduces the rate of Lyman continuum photons necessary to sustain the photoionization of the shell and excess photons become available, which drive the ionization front outward and photoevaporate additional material from the neutral shell of swept-up ambient medium. The inward flow of evaporated material collides with the outward flow of the dissolving swept-up \\HII shell, temporarily forming a new shell of enhanced density. Since all the plasma within the \\HII region is almost isothermal, this density fluctuation vanishes soon.  Another consequence of the reduced stellar wind luminosity (compared to the 60 $\\Msun$ case) is the fact that for most of the time the shape of the \\HII region is more or less preserved as a broad shell interior to the thin shell of swept-up ambient material rather than being so thin as in the 60 $\\Msun$ case where the \\HII region is compressed into the illuminated inner part of the outer shell. Nevertheless, the geometrical thickness of the photoionized shell shrinks at the end of the simulation when the stellar wind luminosity reaches its maximum during the final W-R stage.  The morphological impact of the low-velocity mass-loss phase (the RSG stage) is more prominent than that of the 60 $\\Msun$ star (the LBV stage), because the total mass loss during the RSG stage is higher ($\\approx$18.6~$\\Msun$) than during the LBV stage of the 60 $\\Msun$ star ($\\approx$7.3~$\\Msun$). The slowly expanding RSG wind material is subsequently swept up and accelerated by the shocked W-R wind so that Rayleigh-Taylor instabilities break it into filaments.  The final, pre-supernova structure that shows up in the 35 $\\Msun$ case after 4.945 Myr is basically comparable to that of the 60 $\\Msun$ case at its end. The entire bubble is slightly smaller, the \\HII shell at the inside of the outer neutral swept-up shell is more extended than in the 60 $\\Msun$ case. Outer shell and \\HII region are less clumpy and show less rippling as a consequence of the different evolutionary scenarios.  The reduced substructure formation during the early MS phase is also reflected in the circumstellar energetics. The decrease of ionization energy and thermal energy of warm gas (compared to the respective case without stellar wind) resulting from the formation of dense photoionized structures with short recombination times is smaller than for the 60 $\\Msun$ case. Ionization energy dominates the energy in the circumstellar gas for most of the evolution. The kinetic energy of bulk motion as well as the warm and the hot component of thermal energy stay fairly close together from $t \\approx 0.3~\\mathrm{Myr}$ until the star enters the RSG phase.  The energetic variations during the RSG stage are stronger than those during the LBV stage of the 60 $\\Msun$ star because, due to the duration of the RSG phase, the ionizing radiation of the star is switched off for a considerable period of time so that the photoionized regions recombine. In addition, the mass that is ejected during the RSG stage and accelerated during the subsequent W-R stage is about 2.5 times higher than the mass ejected during the LBV stage of the 60 $\\Msun$ star.  At the end of the 35 $\\Msun$ simulation the total energy transfer efficiency is 1\\%. This value is about the same as in the corresponding case without stellar wind, but it is 2.7 times higher than the value at the end of the 60 $\\Msun$ simulation. 54\\% of the net energy which has been added to the system is then in the form of ionization energy, 36\\% in thermal energy and 10\\% in kinetic energy of bulk motion. The corresponding values at the end of the 60 $\\Msun$ calculation are 25\\%, 40\\%, and 35\\%, respectively. This is another indication that the stellar wind plays a more prominent role in the 60 $\\Msun$ case.  Remarkable agreement of the X-ray properties is found when comparing our model calculations during the early W-R phase with observations of S308. The order of magnitude of the observed X-ray luminosity as well as the temperature of the emitting plasma and the limb brightening of the intensity profile are well reproduced. The obvious explanation that our model overcomes the ``missing wind problem'' described in {\\S} \\ref{sec_observations} is that almost the entire X-ray emission during this phase comes from the W-R shell rather than from the shocked W-R wind. Analytical models constructed so far \\citep[see e.g.][]{garcia95a} assume the W-R shell to be thin and cool so that the energy in the forward shock is completely dissipated in the low-energy wavelength range. The source of X-rays in these models is the shocked W-R wind and the efficiency of heat conduction and thermal evaporation between the hot gas and the cold W-R shell strongly influences the luminosity and the spectral shape of the X-ray emission. If heat conduction is efficient enough to cool the hot shocked W-R wind down to the observed temperature and if the W-R wind luminosity in the models is adjusted to reproduce the observed X-ray luminosity, the W-R wind luminosity is usually much lower than observed \\citep{wrigge99}.  Another factor that reduces the X-ray luminosity in our model to values roughly comparable with S308 is the assumed set of chemical abundances. Although He and N are overabundant according to the observations in the nebula, the underabundance of the other metals reduces the X-ray luminosity by a factor of $3-4$ compared to what can be expected from solar chemical composition.  A further consequence of this interpretation of our results is that the $\\Ha$ emission originates mostly from the RSG shell. This is in agreement with the finding of \\citet{chu03} that the X-ray emission is completely interior to the optical shell. Since the age of S308 (and other W-R bubbles) was hitherto derived from the expansion velocity of the optical nebula under the assumption that the nebula is part of the W-R shell, our results imply an age of S308 which is much younger (${\\approx}2 \\times 10^4~\\mathrm{yr}$) than assumed so far.  However, the match of our model data and the observations is worse for the case of NGC 6888. Besides numerical or model restrictions, differences of the mass-loss and luminosity history between the central star of NGC 6888 and our model star might be responsible for the discrepancies."
    },
    "0512/astro-ph0512326_arXiv.txt": {
        "abstract": "% Over much of the initial mass function, stars are destined to become luminous and cool red giants. They may then be able to produce dust in an atmosphere which has been elevated by strong radial pulsations, and hence drive a wind. The amount of mass that is lost in this way can be a very significant fraction of the stellar mass, and especially in the case of intermediate-mass stars it is highly enriched. The delay between a star's birth and its feedback into the environment varies from several million years for massive stars to almost the age of the Universe for the least massive red giants we see today. I here present a review on the metallicity dependence of red giant winds. I show that recent measurements not only confirm theoretical expectations, but also admonish of common misconceptions with implications for feedback at low initial metallicity. ",
        "introduction": "Before discussing their stellar winds, I will first briefly introduce the red giants themselves. Metalliciy is an important factor in determining the structure and photospheric properties of these stars. Metal-poor stars are generally more compact and have warmer photospheres because they lack opacity in their mantles. But chemical enrichment of their mantles and photospheres can influence that.  \\subsection{Asymptotic Giant Branch stars with carbon+oxygen cores}  A star with an initial, zero-age main sequence mass $M_{\\rm ZAMS}\\sim1$ to 8 M$_\\odot$ will at some point in its evolution switch from core-hydrogen burning to hydrogen-shell burning, and then via a phase of core-helium burning to hydrogen-shell burning again. During that final phase it will become increasingly luminous as the energy-producing shell deposits its waste onto the growing core, and as a reaction its mantle will expand and its surface cool: the star ascends the Asymptotic Giant Branch (AGB). With a luminosity of up to $L_{\\rm AGB-tip}\\sim6\\times10^4$ L$_\\odot$, AGB stars are powerful beacons that can be used to probe intermediate-age populations between $\\sim30$ Myr and 10 Gyr old --- i.e.\\ over much of the Universe's history. Especially at infrared wavelengths the contrast with the underlying main-sequence population and more massive, hotter stars is favourable.  An AGB star is most easily recognised if it shows regular variability in brightness with a period of the order of a year or more at an amplitude that can be several magnitudes at optical wavelengths, when it is called a ``Mira variable'' after its peculiar prototype $o$\\,Ceti (Mira). This variability is explained by radial pulsation of the photosphere as a result of an instability in the balance between the gravitation and radiation pressures, which arises around the hydrogen and/or helium recombination zone(s) where the opacity changes abruptly: the $\\kappa$ mechanism (Christy 1962). The pulsation is most vigorous if excited in the fundamental mode, when the rate at which energy is deposited into the expanding mantle and released as the mantle contracts again is highest.  On the upper slopes of the AGB, the pressure between the core and the hydrogen-burning shell steadily grows until suddenly this helium-rich interface layer ignites. The helium-burning shell eventually extinguishes itself and the star switches back to hydrogen-shell burning. This repeats itself on timescales of $10^4$ to $10^5$ yr: thermal pulses (Iben \\& Renzini 1983). The effects, which are larger for less massive mantles, are twofold: (1) the luminosity and temperature at the surface undergo an excursion either way, and (2) the convective zone which occupies the outer mantle dips into the fresh produce of nucleosynthesis, enabling these products to travel to the stellar surface where they modify the photospheric abundance pattern. The latter is called ``$3^{\\rm rd}$ dredge-up'' (because two earlier phases of surface enrichment are possible). This leads to enrichment in carbon and products from the slow neutron capture process such as the unstable element technetium, of which $^{99}$Tc with a half-life of $2\\times10^5$ yr is an unambiguous indicator of a thermal-pulsing AGB star.  The most dramatic result is obtained when the surface becomes so enriched in carbon that the carbon atoms outnumber those of oxygen: a carbon star. This is especially important because the surfaces of AGB stars are cool enough, $T_{\\rm eff}<4,000$ K, to form a molecular atmosphere. Although molecular hydrogen, H$_2$, is the most abundant molecule in the presence of dust grains, it is also the simplest and symmetric molecule; as such it is very hard to notice by its opacity or emissivity. Most of the molecular opacity is due to molecules involving either a carbon or an oxygen atom, due to their abundance and high reactivity. It is therefore not surprising that carbon-monoxide, CO, is the most common molecule after H$_2$. It also means that the molecular chemistry is dominated by which ever is left after the formation of CO: oxides such as H$_2$O (water), TiO and SiO, or carbonaceous molecules such as C$_2$, CN and C$_2$H$_2$ (acetylene). Stellar models and observations suggest that carbon stars are only produced for stars more massive than $M_{\\rm ZAMS}\\sim1.3$ to 1.5 M$_\\odot$, and that AGB stars more massive than $M_{\\rm ZAMS}\\sim4$ M$_\\odot$ do not become carbon stars because the carbon is burnt to nitrogen at the bottom of the convection zone: Hot Bottom Burning (Marigo, Girardi \\& Bressan 1999; van loon, Marshall \\& Zijlstra 2005).  \\subsection{Red Supergiants}  In the cores of massive stars, with $M_{\\rm ZAMS}>8$ M$_\\odot$, eventually carbon burning takes place. This creates an oxygen+neon core. At the lower mass end, the evolution of these stars closely resembles that of AGB stars, hence these are called ``super-AGB'' stars (Gil-Pons et al.\\ 2005). More massive stars will not undergo thermal pulses but they will continue to burn ever heavier elements after core-carbon burning is exhausted until the iron barrier induces core collapse. During core-helium burning, such a massive star may become cool enough to also develop a convective, pulsationally unstable mantle wrapped in a molecular atmosphere: a red supergiant (RSG), with a luminosity $L_{\\rm RSG}\\sim10^5$ L$_\\odot$.  \\subsection{Red Giant Branch stars with electron-degenerate helium cores}  The helium cores of low-mass stars, with $M_{\\rm ZAMS}<1.5$ to 2 M$_\\odot$, are electron-degenerate. With their cores so compact, in reaction their mantles become relatively extended, making them already very cool during the first phase of hydrogen-shell burning. As first-ascent Red Giant Branch (RGB) stars they reach luminosities up to $L_{\\rm RGB-tip}<3\\times10^3$ L$_\\odot$, until core-helium burning ignites which ends the electron degeneracy in the core. On the upper reaches of the RGB much of the photosphere is located in the molecular atmosphere, and the convective mantle may undergo radial pulsation in an overtone mode. Hence, like red supergiants, RGB stars share similarities with AGB stars. ",
        "conclusions": "Observations support the paradigm that as long as a luminous giant star is able to develop a cool, molecular atmosphere it will also pulsate and develop a dust-driven wind. The dust:gas ratio is then irrelevant for the mass-loss rate, although it does determine the wind speed. Although metal-poor stars may not reach that stage, the alternative mechanism of chromospherically-driven mass loss may in fact be more efficient for them, and total mass-loss rates may not be lower than for metal-rich red giants.  That all is not that simple is plainly demonstrated, for instance by the surprising observation of dust in metal-poor globular clusters (Origlia et al.\\ 2002; Evans et al.\\ 2003). Because the nearby Universe does not offer a direct means of observing supergiants or massive AGB stars at $Z<0.01$ Z$_\\odot$, it is essential that a complete {\\it theory} for the mass loss from red giants be established. Until then, there are plenty of ways in which {\\it observations} can further our understanding of the red giant winds. More precise measurements of the optical depth in large samples of red giants in the Magellanic Clouds and the nearest metal-poor dwarf spheroidal galaxies might resolve differences in the {\\it inferred} mass-loss rates (such as in Fig.\\ 2) between these galaxies for different scaling laws of the dust:gas ratio. Future measurements of CO in the envelopes around magellanic red giants by ALMA will provide values for the wind speed and dust:CO ratio in both carbon stars and oxygen-rich AGB stars and RSGs. This will greatly improve the accuracy of the derived mass-loss rates and test the dust formation and wind acceleration mechanisms. Measurements of the carbon and oxygen abundances in massive AGB stars as a function of metallicity are possible and needed to test theories that predict them to become carbon stars at very low metallicity. Chromospherically-driven winds need to be studied for metal-poor AGB stars and RSGs to gauge whether this would be a viable alternative to dust-driven winds. It is too early to discard the r\\^{o}le of red giants in the Early Universe."
    },
    "0512/astro-ph0512332_arXiv.txt": {
        "abstract": "We study the evolution of the fine-structure constant, $\\alpha$, induced by non-linear density perturbations in the context of the simplest class of quintessence models with a non-minimal coupling to the electromagnetic field, in which the two available free functions (potential and gauge kinetic function) are Taylor-expanded up to linear order. We show that the results obtained using the spherical infall model for an infinite wavelength inhomogeneity are inconsistent with the results of a local linearized gravity study and we argue in favour of the second approach. We also discuss recent claims that the value of $\\alpha$ inside virialised regions could be significantly different from the background one on the basis of these findings. ",
        "introduction": "One of the deepest question of modern physics is whether or not there are fundamental scalar fields in nature. For example, they are a key ingredient in the standard model of particle physics (cf. the Higgs particle, which is supposed to give mass to all other particles and make the theory gauge-invariant), but after several decades of accelerator experiments there is still no shred of experimental evidence for them.  The early universe is a much better (not to mention cheaper) laboratory for fundamental physics. Observations suggest that the recent universe is dominated by an energy component whose gravitational behaviour is quite similar to that of a cosmological constant (as first introduced by Einstein). This could of course be the right answer, but the observationally required value is so much smaller than what would be expected from particle physics that a dynamical scalar field is arguably a more likely explanation. Now, the slow-roll of this field (which is mandatory so as to yield negative pressure) and the fact that it is presently dominating the universe imply (if the minimum of the potential vanishes) that the field vacuum expectation value today must be of order $m_{Pl}$, and that its excitations are very light, with $m\\sim H_0\\sim 10^{-33}$ eV. But a further consequence of this is seldom emphasized \\cite{CARROLL}: couplings of this field lead to observable long-range forces and time-dependence of the constant of nature (with corresponding violations of the Einstein Equivalence Principle).  Measurements of various dimensionless couplings, such as the fine-structure constant $\\alpha$ (which will be the focus of this paper) or the electron to proton mass ratio \\cite{THOMPSON} are therefore unique tests of fundamental physics. Note that since the scalar field is effectively massless on solar system scales, it should in principle be easier to find new physics on astrophysical and cosmological scales. Moreover, bounds on varying couplings restrict the evolution of the scalar field and provide constraints on dark energy \\cite{PARKINSON,NUNES} that, with new datasets becoming available in the near future will be complementary to (and indeed more powerful and constraining than) those obtained by traditional means.  Let us now focus on the fine-structure constant, $\\alpha$, which among other things measures the strength of the electromagnetic interaction. The good news is then that, since the standard physics is changed in a number of key ways if there is a spacetime variation of $\\alpha$, there are many different ways in which measurements of $\\alpha$ can be made. To name just a few, locally one can use atomic clocks \\cite{Marion} or the Oklo natural nuclear reactor \\cite{Fujii,Lamoreaux}. On the other hand, on astrophysical and cosmological scales a lot of work has been done on measurements using quasar absorption systems \\cite{Webb1,Webb2,Murphy,Chand} and the cosmic microwave background \\cite{Avelino1,Avelino2,Martins1,Martins2,Rocha}.  The bad news, however, is that these different measurements probe very different environments, and therefore it is not trivial to compare and relate them. Simply comparing at face value numbers obtained at different redshifts, for example, it is at the very least too naive, and in most cases manifestly incorrect. Indeed, detailed comparisons can often only be made in a model-dependent way, meaning that one has to specify a cosmological model (crucial to define a clock in the universe, that is, a timescale) and/or a specific model for the evolution of $\\alpha$ as a function of redshift. Simply assuming, for example, that alpha grows linearly with time (so that its time derivative is constant) is not satisfactory, as one can easily see that no sensible particle physics model will ever yield such a dependence for any significant redshift range. Last but not least, new methods are being developed for measuring $\\alpha$ using emission lines \\cite{SDSS,JARLE,GERKE} as well as the electron to proton mass ratio (using absorption) \\cite{Ivanchik1,Ivanchik2,Ubachs}, so the issue of detailed comparisons between datasets will be even more important for the next generation of datasets.  Here, we discuss one specific aspect of this issue. The scalar field responsible for the variation of $\\alpha$ will (in any sensible particle physics model) couple to the matter sector. Among other things, this implies that when, in the course of the cosmological evolution, inhomogeneities grow, become non-linear and decouple from the background evolution, the same could happen to the local variations of $\\alpha$. This has been previously studied in \\cite{Barrow1,Mota1,Mota2} using a simple spherical infall model for the evolution of infinite wavelength density perturbations and a particular generalization of the Bekenstein model \\cite{Sandvik} for the evolution of $\\alpha$. (See also \\cite{Barrow2} for a discussion of large scale variations of $\\alpha$.) It was found that in the linear regime and in the matter era the variation of $\\alpha$ would follow the density contrast. Moreover, it was also claimed that this approach was valid in the non-linear regime (meaning turnaround and collapse). Here we revisit these results, in particular questioning the applicability of the spherical infall model. Note that in the particular case of the models of \\cite{Barrow1,Mota1,Mota2,Sandvik} it is enforced (purely by hand) that the evolution of $\\alpha$ is driven by a coupling to charged non-relativistic matter alone and in that case the scalar field cannot also provide the dark energy. For that reason one should see these models as toy models that are useful for computational purposes. To some extent a similar comment applies to the models that we shall consider, though the reason here is purely their extreme simplicity. However, we do expect that at least at a qualitative level our results will be representative of more realistic models.  While this paper was being written up, ref. \\cite{Shaw} appeared. This provides a more detailed and mathematically-inclined analysis of local variations in physical `constants', but does confirm our results. Our approach, while much simpler, has the advantage of making explicit the reasons why spatial variations have to be small, and why the use of the spherical collapse model of a infinite wavelength perturbation is inadequate.  We will start in Sect. \\ref{lin} with a brief description of the models that we will be using. In Sect. \\ref{nonlin} we discuss the non-linear evolution of the fine structure constant using two different approaches. Finally, we describe and discuss our results in Sect. \\ref{res}, and present our conclusions in Sect. \\ref{end}. Throughout this paper we shall use fundamental units with $\\hbar=c=G=1$. ",
        "conclusions": ""
    },
    "0512/astro-ph0512618_arXiv.txt": {
        "abstract": "Mid-infrared observations of Active Galactic Nuclei (AGN) are important for understanding of the physical conditions around the central accretion engines. {\\it Chandra} and {\\it XMM-Newton} X--ray observations of a 300 arcmin$^2$ region in the Extended Groth Strip are used to select a sample of $\\sim 150$ AGN. The {\\it Spitzer} instruments IRAC and MIPS detect 68--80\\% of these sources, which show a wide range of mid-infrared properties. About 40\\% of the sources have red power-law spectral energy distributions ($f_{\\nu} \\propto {\\nu}^{\\alpha}$, $\\alpha<0$) in the $3.6-8$\\mic IRAC bands. In these sources the central engine dominates the emission at both X--ray and IR wavelengths. Another 40\\% of the sources have blue mid-IR spectral energy distributions ($\\alpha>0$) with their infrared emission dominated by the host galaxy; the remaining 20\\% are not well-fit by a power law. Published IRAC color criteria for AGN select most of the red sources, but only some of the blue sources. As with all other known methods, selecting AGN with mid-IR colors will not produce a sample that is simultaneously complete and reliable. The IRAC SED type does not directly correspond to X--ray spectral type (hard/soft). The mid-IR properties of X--ray-detected Lyman-break, radio, submillimeter, and optically-faint sources vary widely and, for the most part, are not distinct from those of the general X--ray/infrared source population. X--ray sources emit 6--11\\% of the integrated mid--IR light, making them significant contributors to the cosmic infrared background. ",
        "introduction": "Understanding the nature of active galactic nuclei (AGN) and the galaxies that host them is important for such diverse goals as pinpointing the sources of the cosmic X--ray and infrared backgrounds and deriving the star formation history of the universe. Multi-wavelength surveys are particularly important for the study of AGN because their appearance in different wavelength regimes can be quite different. Selecting AGN at one particular wavelength is no guarantee of a complete sample. X--ray selection has the advantages of being reasonably efficient and reliable \\citep{mush04} but may miss some obscured sources \\citep{pet05}. Infrared and radio observations can identify AGN missed in the X--ray \\citep{donley05,aah05} and, for X--ray-selected AGN, can help to distinguish between the different AGN types. Infrared data are also needed to constrain the fraction of the cosmic infrared background (CIRB) emitted by AGN.  The limited spatial resolution of previous generations of X--ray and infrared observatories made cross-identification between wavelengths difficult. Recent work has benefited from the much smaller point spread functions of the {\\it Chandra X--ray Observatory}, {\\it XMM-Newton}, the {\\it Infrared Space Observatory (ISO)}, and the {\\it Spitzer Space Telescope}. \\citet{fadda02} combined {\\it Chandra}, {\\it XMM}, and {\\it ISO} data from the Lockman Hole and Hubble Deep Field North to conclude that AGN contribute $15\\pm5$\\%  of the CIRB at 15\\mice. Work with {\\it Spitzer} data to date has concentrated on the properties of the X--ray source population. \\citet{rigby04} studied Chandra Deep Field South sources detected at hard X--ray and 24\\mic wavelengths and  found, surprisingly, that X--ray-hard AGN are not infrared-brighter (as would be expected if they were embedded in the obscuring matter). \\citet{aah04} found that similarly-selected Lockman Hole sources exhibit a variety of optical/IR spectral types. About half of their sources had spectral energy distributions (SEDs) dominated by stellar emission or showed significant obscuration. \\citet{fra05} found a similar mix of spectral types among ELAIS-N1 {\\it Chandra}/{\\it Spitzer} sources and concluded that about 10--15\\% of 24\\mic sources are dominated by an AGN.  This paper combines {\\it Spitzer}, {\\it Chandra}, and {\\it XMM} observations to understand the mid-infrared properties of the X--ray sources in the `Extended Groth Strip' (EGS). The X--ray and infrared observations in this region are intermediate in depth and area between GOODS \\citep{goods} and the shallower NOAO Deep-Wide Field \\citep{ndwfs,irac_shallow} and SWIRE \\citep{swire} surveys. As such, the EGS provides a valuable probe of the properties of AGN at intermediate fluxes and an additional measure of the cosmic variance of those properties. The extensive multi-wavelength observational data available for the EGS, particularly its spectroscopic redshift survey, should eventually produce an extremely thorough characterization of the X--ray sources and allow comparisons and cross-identifications with other classes of galaxies. In this paper we focus on combining mid-IR and X--ray data to study the properties of AGN and their host galaxies. Identifications of X--ray sources in the {\\it Chandra} deep fields \\citep{horn03, bar03} suggests that almost all of the EGS X--ray sources should be AGN. The EGS X--ray sources' median redshift is expected to be $z\\sim1$ with no strong dependence on X--ray flux \\citep{bar05}. ",
        "conclusions": "Of about 150 X--ray sources within the Extended Groth Strip, more than 90\\%  have IRAC counterparts at flux densities $>1-6 \\mu$Jy, and about two-thirds are detected with MIPS at flux densities $>83\\mu$Jy. At the flux limits of the X--ray surveys, most of the sources are expected to be AGN. The ratios of X--ray to optical and IR flux are consistent with this expectation.  The infrared SEDs of the X--ray sources show a broad range of properties but reasonable agreement with predicted colors from nearby template objects. About 40\\% of the X--ray sources have a red power-law SED dominated by the AGN. The remaining 60\\% of the sources have either blue or non-power-law IRAC SEDs, indicating domination by galaxy light, PAH emission, or a mixture of galaxy and AGN light. \\citet{fra05} found that 39\\% of {\\it Chandra} sources in the ELAIS-N1 region had optical/infrared SEDs classified as `QSO' or `Seyfert~1'; if most of our red sources are also `type 1' (unobscured) AGN, then there is good agreement between the two studies. Published  IRAC color-color criteria select the EGS X--ray sources with 5--9\\% reliability (these values would likely be higher with deeper X--ray observations) and 25--75\\% completeness.  There are good correlations between IR and X--ray fluxes for the red sources, indicating that the AGN dominates in both wavelength regimes. We find only marginal evidence (a difference in IR spectral index $\\alpha$ between hard and soft sources) for agreement between AGN classifications based on IR SEDs and X--ray hardness ratios even though the amount of obscuration should affect the observed properties in both. Variation in the gas-to-dust ratio and broad ranges of the intrinsic AGN properties could account for some of the lack of correspondence between IR and X--rays. Such an explanation is necessary not only for the properties of observed X--ray sources but to account for the many non-X--ray-detected sources that have similar IR properties.  The X--ray sources detected at other wavelengths show a wide range of properties. None are completely distinct from the main X--ray/IR sample, but most of the subsets are infrared-fainter and redder. The Lyman-break/X--ray sources are bright and blue when compared to other LBSs but faint and red (consistent with being at high redshift) compared to other X--ray sources. Two sub-millimeter sources are quite different in their X--ray properties, implying that the properties of the central AGN and of the intense star formation producing the sub-mm emission are not necessarily connected. One of the sub-mm sources and one additional X--ray source are `over-luminous' in the mid-IR and we propose several possible explanations. The radio-detected X--ray sources divide into two groups based on their mid-IR and X--ray properties, but these groups do not relate to the radio emission. The optically faint sources (including 4 sources undetected in the optical) are again undistinguished in their mid-IR and X--ray properties.  The integrated infrared light from X--ray sources provides a lower bound to the fraction of the cosmic IR background originating from AGN. Our measurements of the integrated light from the {\\it Chandra} sources are about half of the predicted AGN$+$host values from \\citet{smg04}: as expected, the X--ray observations do not detect all the AGN. The amount of integrated light from the EGS X--ray sources is in surprisingly good agreement with the \\citeauthor{smg04} predictions for Compton-thin AGN: this may be a coincidence (if the X--ray sources contain the right number of thick and thin AGN) or an indication that most non-X--ray-detected AGN are Compton-thick. Disagreements between the observed fractions of IR light from X--ray sources and the models at some wavelengths may be due to under- or over-predictions of the total CIRB.   Future work on this topic will benefit from the ongoing observations in the EGS region. Statistics will be improved by the larger, more uniform sample of X--ray sources soon to become available from the {\\it Chandra} observations of the full 2\\arcdeg-long EGS. Spectroscopic redshift information from the DEEP2 survey will allow determination of luminosities and $K$-corrections such that intrinsic properties can be computed. Cross-identification with observations at other wavelengths will allow comprehensive spectral energy distributions to be constructed. Such SEDs can be compared with population synthesis models to estimate galaxy stellar masses and star formation histories, which can in turn be compared to black hole masses computed from X--ray luminosities. The large sample of galaxies observed in the EGS no doubt contains many non-X--ray-selected AGN, and a better understanding of the X--ray-selected sources may be of tremendous help in finding these `dark' AGN."
    },
    "0512/astro-ph0512104_arXiv.txt": {
        "abstract": "Gravitational lensing can be used to directly constrain the projected density profile of galaxy clusters. We discuss possible future constraints using lensing of the CMB temperature and polarization, and compare to results from using galaxy weak lensing. We model the moving lens and kinetic SZ signals that confuse the temperature CMB lensing when cluster velocities and angular momenta are unknown, and show how they degrade parameter constraints. The CMB polarization cluster lensing signal is $\\sim 1\\muK$ for massive clusters and challenging to detect; however it should be significantly cleaner than the temperature signal and may provide the most robust constraints at low noise levels. Galaxy lensing is likely to be much better for constraining cluster masses at low redshift, but for clusters at redshift $z\\agt 1$ future CMB lensing observations may be able to do better. ",
        "introduction": "The distribution of clusters of galaxies as a function of mass and redshift depends on the cosmological model, and can be modelled increasingly accurately~\\cite{Birkinshaw:1998qp,Shaw:2005dy,Battye:2003bm,Majumdar:2003mw}. Observations of clusters can therefore be used to learn about cosmology, as well as to test models for cluster formation and evolution. Observations of the thermal Sunyaev-Zel'dovich (SZ) effect~\\cite{SZ,Birkinshaw:1998qp,Battye:2003bm} are a powerful probe of the cluster gas, but do not measure the mass directly. To relate the gas properties to the total mass involves modelling potentially complicated baryonic gas physics. By contrast gravitational lensing probes the projected total mass, not just the gas, and therefore can provide direct information about cluster masses. In this paper we analyse the potential for reconstruction of parameterized cluster profiles from future observations of cluster lensing of the CMB and weak lensing of distant galaxies. For the first time we include constraints from CMB polarization, and also include a model of the moving lens effect that confuses the CMB temperature signal when the clusters have unknown velocities. The CMB polarization signal is much cleaner than the temperature at low noise levels, and may prove to be a good way to constrain cluster masses at high redshift. We perform an essentially optimal statistical analysis in the approximation that the unlensed fields can be treated as Gaussian.  Current CMB observations are not of high enough resolution or sensitivity to measure the cluster lensing signal. However, future missions aimed at detecting small levels of primordial gravitational waves via their distinct $B$-mode polarization signal will require both high sensitivity and resolution. This is because lensing by large scale structure can convert scalar $E$-modes into $B$-modes, and hence this lensing signal has to be subtracted to extract a small primordial $B$-mode signal from gravitational waves. The lensing reconstruction requires high resolution observations in order to have enough information to solve for both the primordial $B$-modes and the unknown large scale structure distribution~\\cite{Seljak:2003pn}. It is therefore of interest to see what other useful information can be gained from such future high resolution observations. The resolution required for $B$-mode cleaning is probably rather less than needed for good cluster mass constraints from CMB lensing, however it is clearly of interest to see what can be gained from observations with slightly higher arcminute-level resolution. Cluster lensing of CMB temperature and polarization that we consider here is one potentially useful possibility.  Cluster lensing also generates a shear field that is observable by looking at the shapes of galaxies lying behind the cluster. This method of cluster mass constraint is promising and  possible with current observations, though at some level has to be limited by the finite number of source galaxies available behind the lens. The seminal paper of  Kaiser \\& Squires~\\cite{Kaiser:93} describes how to do a non-parametric mass reconstruction; this technique and its variants have been applied to numerous clusters (e.g. Refs.~\\cite{Fischer:1997si,Clowe:2000iq}). Here we focus on parameterized cluster models, which enable statistical comparisons to be made between clusters and as a function of observational strategy. The likelihood techniques are based on those developed in Refs.~\\cite{Schneider:1999ch,King:2000wc}. We shall investigate how galaxy lensing compares to CMB lensing as a function of noise level, cluster redshift and galaxy number count. We use natural units where the speed of light is unity. ",
        "conclusions": "Weak lensing is a valuable and promising method for studying cluster masses. Measurements of lensing shear using observations of lensed galaxies provides tight constraints for low redshift clusters. For clusters at higher redshift than the peak of the galaxy-redshift distribution the galaxy lensing constraints become much poorer, and CMB lensing can do better. However, even with simple spherically symmetric models the temperature lensing signal can be degraded by various other second order effects. For futuristic arcminute-resolution observations at low noise levels the CMB polarization lensing signal may be much cleaner and a more robust way to measure cluster properties. Measurements of lensing by high redshift clusters is therefore something that future CMB polarization missions may wish to aim to achieve.  The results from galaxy lensing are limited by the intrinsic ellipticity dispersion of the galaxies, and the fact that there are only a finite number of sources behind the cluster. To do better one could try to find sources which have a higher number density. Possibilities include high redshift sources observed with 21cm, sources from the time of inhomogeneous reionization, and secondary doppler CMB signals from velocities after reionization~\\cite{Zaldarriaga:1998te,Pen:2003yv,Zahn:2005ap}. In addition strong lensing can be used to help constrain the central region of the cluster profile. The ultimate limit from CMB polarization will depend on how efficiently spectral information can be used to clean out confusing signals, and the extent to which cluster substructure complicates the signal from quadrupole scattering. Future work could investigate this using numerical simulations.  With the appropriate increase in resolution and sensitivity, methods for cluster CMB lensing could be extended to constrain galaxy profiles (as discussed for the CMB temperature in Ref.~\\cite{Dodelson:2003gv})."
    },
    "0512/astro-ph0512274_arXiv.txt": {
        "abstract": "Results from a study of Fe K${\\alpha}$ emission lines for a sample of six non-magnetic Cataclysmic Variables (CVs) using the high resolution X-ray data from the \\emph{Chandra} High Energy Transmission Grating (HETG) are presented.  Two of the sources, SS~Cyg and U~Gem are observed in both quiescent and outburst states whereas V603~Aql, V426~Oph, WX~Hyi and SU~UMa are observed only in quiescence. The fluorescent Fe line is prominent in V603~Aql, V426~Oph and SS~Cyg during quiescence indicating the presence of a conspicuous reflection component in these sources.  The observed equivalent width of the fluorescent Fe line is consistent with reflection from a white dwarf surface that subtends 2$\\pi$ solid angle at the X-ray source. During the outburst in SS~Cyg, the fluorescent line is red-shifted by about 2300 km s$^{-1}$.  The Fe XXV triplet at 6.7 keV is found to be dominant in all sources. The value of the G-ratio derived from the Fe XXV triplet indicates that the plasma is in collisional ionization equilibrium during the quiescent state. The Fe XXV line is significantly broadened in U~Gem and SS~Cyg during the outbursts compared to quiescence, indicating the presence of a high velocity material near the white dwarf during the outburst.  The ratio of Fe XXVI/XXV indicates a higher ionization temperature during quiescence than in outburst in U~Gem and SS~Cyg. ",
        "introduction": "Non-magnetic cataclysmic variables (CVs) are a subclass of CVs in which a white dwarf with a weak magnetic field (B$\\lesssim$10$^{4}$G) accretes material from the Roche lobe of a late type dwarf companion star \\citep{war95}. The accretion takes place via a disk around the white dwarf.  Due to the weakness of the magnetic field the disk extends to the surface of the white dwarf. The Keplerian velocity of the material in the disk is generally greater than the rotational speed of the white dwarf.  Near the inner edge of the disk the material slows down to match the white dwarf rotation.  The X-ray emission is thought to arise from this boundary layer between the white dwarf and the inner edge of the accretion disk.  Dwarf Novae (DNe) are a subclass of the non-magnetic CVs that show frequent (at the interval of weeks to months) outbursts.  According to theory, a thermal--viscous instability in the disk produces repetitive outbursts in these systems \\cite[see][for a review]{lasota01}.  On the other hand, in classical novae, another subclass of non-magnetic CVs, the eruptions are due to thermonuclear runaway of hydrogen rich material accreted on to the white dwarf. The amplitude of outburst in these systems is greater than in DNe.  A sample of DNe were studied using low-resolution (E/$\\bigtriangleup$E $\\simeq$ 5 @ 6 keV) \\emph{EXOSAT} medium energy (ME) data \\citep{mukai93} and it was found that line emission near 6.7 keV is a common feature in the hard X-ray spectra of DNe. The Fe XXV triplet is one of the most intense set of lines in the hard X-ray spectra of CVs \\citep{pandel05}.  The line emission near 6.7 keV originates from Fe K${\\alpha}$ emission that has three main components: a fluorescent line at 6.41 keV, Fe XXV line (He-like) that has four subcomponents (a resonance line `$r$' at 6.7002 keV, two intercombination lines `$i_{1}$' and `$i_{2}$' at 6.6821 and 6.6673 keV, respectively and a forbidden line `$f$' at 6.6364 keV), and Fe XXVI Ly$_{\\alpha}$ line (H-like) with two subcomponents at 6.973 and 6.952 keV.  The Fe XXV triplet ($r$, $i$, and $f$) is due to the transitions 1s2p $^{1}$P$_{1}$ $\\longrightarrow$ 1s$^{2}$ $^{1}$S$_{0}$, 1s2p $^{3}$P$_{2,1}$ $\\longrightarrow$ 1s$^{2}$ $^{1}$S$_{0}$, and 1s2s $^{3}$S$_{1}$ $\\longrightarrow$ 1s$^{2}$ $^{1}$S$_{0}$.  The two components of the Fe XXVI line are due to the transitions 2p $^{2}$P$_{3/2}$ $\\longrightarrow$ 1s $^{2}$S$_{1/2}$ and 2p $^{2}$P$_{1/2}$ $\\longrightarrow$ 1s $^{2}$S$_{1/2}$.  The Fe XXV and XXVI lines come from a plasma having a temperature of 10$^{7-8}$ K. The fluorescent line originates from relatively cold iron (Fe I--XVII) having temperatures $\\leq$10$^{6}$ K.  The low-resolution of the \\emph{EXOSAT} ME data was insufficient to distinguish between the various components of Fe K$_{\\alpha}$ emission.  Moderate resolution (E/$\\bigtriangleup$E = 50 @ 6 keV) spectroscopy with \\emph{ASCA} and \\emph{XMM-Newton} has been a little more successful in resolving the components of Fe K$_{\\alpha}$ emission, and several authors have reported Fe K${\\alpha}$ lines for individual non-magnetic systems using the \\emph{ASCA} \\citep[see][for SS~Cyg]{done97} and \\emph{XMM-Newton} data \\citep[see][for YZ~Cnc and references therein]{hakala04}. Recently, \\cite{baskill05} presented a study of 34 non-magnetic CVs using \\emph{ASCA} data.  The CCD spectrometer (SIS) allowed them to separate the three Fe K$_{\\alpha}$ lines in the spectra.  Of the 34 objects, only 4 bright sources showed the presence of a prominent fluorescent Fe line at 6.4 keV.  This feature is attributed to the reflection of X-rays from the white dwarf surface and/or the accretion disk.  High spectral resolution (E/$\\bigtriangleup$E $>$200 @ 6 keV) data obtained with \\emph{Chandra} provides an opportunity to study in detail the strength and profile of these lines. The Fe XXV and XXVI emission lines that originate from the high temperature (10$^{7-8}$ K) plasma serve as an important diagnostic tool for temperatures near the shocked regions in the boundary layer of these CVs.  The intensity ratio of Fe XXVI to Fe XXV lines provides a measure of the ionization temperature of the plasma.  Intensity ratios defined using the Fe XXV triplet can be used in principle to get information about the ionization state, temperature and density of the emitting plasma in collisional ionization equilibrium. Specifically, the ratio G, defined as ($f$+$i$)/$r$ is sensitive to the electron temperature as well as the ionization state of the plasma.  The ratio R=$f$/$i$ provides a density diagnostic.  The fluorescent line of Fe that comes from relatively colder material can provide information about the contribution of reflected hard X-rays from the illuminating region.  In this paper, we present a detailed study of Fe K${\\alpha}$ emission lines from a sample of non-magnetic CVs using the best available energy resolution X-ray data available from \\emph{Chandra} grating instruments.  The next section describes the observations and the data analysis procedure.  The results are presented in \\S 3, and are discussed in \\S 4.  The conclusions are summarized in \\S 5. ",
        "conclusions": "\\subsection{Fe XXV triplet: Plasma Diagnostics} The intensity ratios of the emission lines from He-like ion provide valuable spectral diagnostics for temperature, density and ionization state for plasma in collisional ionization equilibrium.  For Fe XXV the principal lines can be contaminated by the dielectronic satellites (DES) for plasma in collisional ionization equilibrium \\citep{oelgoetz01}.  It is not possible to resolve these lines from the principal lines at the available energy resolution but it is possible to infer their relative contribution to the observed spectrum depending on the temperature of the emitting plasma.  The DES dominate principal lines at temperatures below 3$\\times$10$^7$ K and their contribution at higher temperatures is negligible \\citep{oelgoetz01}.  It can be seen from the Fig. \\ref{fig1} and Table \\ref{tbl-2} that the resonance line is stronger compared to the other two components of Fe XXV line for all sources during quiescence, except for SU~UMa. The relative strength of the resonance line indicates that the temperature of the emitting plasma is above 3 $\\times$ 10$^{7}$ K \\citep{oelgoetz01}, where the principal lines dominate.  Therefore, we can use these lines to infer the temperature and density of the emitting region.  The G-ratio is very close to 1 (within 90\\% confidence limit; see Table \\ref{tbl-3}) for most of the sources indicating that the plasma is mainly in collisional ionization equilibrium \\citep{oelgoetz01} with electron temperature T$_{e}$ $\\geq$ 10$^{7}$ K for Fe XXV.  SS~Cyg, during one of the outbursts, shows a somewhat higher G-ratio value ($\\sim$2.4) that might indicate the presence of hybrid plasma in the system. For U~Gem during the outburst, the value of the G-ratio is unconstrained, and therefore does not allow us to comment on the ionization state and the temperature of the emitting plasma.  The mean value of the R-ratio for SS~Cyg (during both quiescence and outburst), V603~Aql, V426~Oph, and WX~Hyi (2002 July 28 observation) is close to unity and varies between 0 to 2.5. For two sources, U~Gem and WX~Hyi (2002 July 25 observation) during quiescence, the R-ratio is very high and essentially unconstrained (see Table~\\ref{tbl-3}).  On the other hand SU~UMa and U~Gem (during outburst) show very low values of the R-ratio. According to the theoretical curves presented by \\cite{bau00}, the Fe~XXV R line ratio is $\\sim$1 for low densities and rolls over to 0 at a critical density of $\\sim$1$\\times 10^{17}$ cm$^{-3}$.  Thus, the above mentioned values of the R-ratio (also see Table~\\ref{tbl-3}) for the Fe~XXV triplets do not allow us to constrain the plasma densities in the non-magnetic CVs studied here. It should be noted, however, that the Fe~XXV triplets are not fully resolved with the HEG data and hence the error estimations on individual line fluxes are not completely independent of each other. This may affect the errors on the G and R-ratios.  Better resolution data with sufficient signal to noise ratio are required to properly constrain the line ratios for these sources.  \\subsection{Fe XXVI to XXV line ratio} In Figure~\\ref{fig4}, we show the observed line ratio of Fe XXVI/Fe XXV in various sources as a function of ionization temperature.  We have used the summed flux of $r$, $i$ and $f$ lines of Fe XXV for calculating this ratio.  The solid curve represents the expected line ratio for Fe as a function of ionization temperature assuming the plasma is in collisional ionization equilibrium \\citep{mewe85}.  It is clear from the Fig.~\\ref{fig4} that the lower limit of the ionization temperature is not constrained by the HEG data for several sources, namely SU~UMa, V603~Aql, WX~Hyi (2002 July 28), and SS~Cyg during one of the outburst observations. However, it is well constrained for U~Gem, SS~Cyg during quiescence, and V426~Oph.  In general the ionization temperature for all CVs studied here is $\\lesssim$12 keV. For SS~Cyg, it has been observed that during optical outburst the hard X-ray (3--20 keV) bremsstrahlung temperature is lower by a factor of $\\sim$2 than during quiescence \\citep{done97,wheatley03,mcgowan04}. Since the plasma lines of Fe originate from the hot shocked region, the ionization temperature derived from these lines is expected to show a similar behavior.  The observed decrease in the ionization temperature during outburst is consistent with previous hard X-ray observations of the continuum in SS~Cyg.  \\subsection{Reflection and the Fluorescent Fe line} In non-magnetic CVs, the fluorescent iron line is believed to arise due to reflection of hard X-rays from the white dwarf surface or the inner edge of the accretion disk.  In the quiescent state at low accretion rates, the inner accretion disk is either absent or optically thin and hence contributes little to the observed reflection component.  Therefore, a significant contribution to the fluorescent Fe line in DNe in the quiescent state comes from the reflection off the white dwarf surface. It can be seen from the Table \\ref{tbl-2} that the fluorescent line in all sources is not very strong (with EW$\\leq$60 eV) except for V603~Aql.  This could be because only the surface of the white dwarf contributes to the fluorescent line.  If the fluorescent line is mainly due to reflection, then the strength of this line should depend on the inclination angle of the system. For a system having a large inclination angle the fluorescent line is expected to be weak, because it prevents the observation of any reflection from the inner edge of the disk \\citep{ramsay01}.  We have looked for such correlations in this sample of CVs.  In Figure \\ref{fig3} we have plotted the observed EW of fluorescent line as a function of system inclination.  The values of the inclination angle have been taken from \\cite{baskill05} and \\cite{war95}.  As the figure demonstrates there is no clear trend between the two parameters.  The systems with inclination angle above 30$^{\\circ}$ show equivalent widths that are similar within the 90\\% confidence interval with the mean value of $\\sim$50 eV.  In our sample there is only one source below 30$^{\\circ}$ (V603~Aql) that shows the highest equivalent width of 162$^{+99}_{-65}$ eV.  We need a larger number of sources with inclinations below 30$^{\\circ}$ to explore any relation between the two parameters.  However, it gives some indication that the systems with very low inclination (below 20$^{\\circ}$) may have a strong fluorescent line whereas those with an inclination above 70$^{\\circ}$ have a weak line.  On the other hand, \\cite{hakala04} reported that a dwarf nova YZ~Cnc has a low inclination angle and still lacks the 6.4 keV fluorescent emission line.  Also \\cite{baskill05} found prominent fluorescent lines in only 4 CVs out of 34 with no correlation of the line strength and the system inclination.  Thus the reason for the strength (or weakness) of this line in the non-magnetic CVs remains unclear.  The old nova V603~Aql shows a strong fluorescent line with the highest EW of 162$^{+99}_{-65}$ eV among all sources (see \\S 3.1). \\cite{mukai05} reports a similar EW for this line and attribute it to the reflection from a surface that subtends an angle of 2$\\pi$ as seen from the primary X-ray source assuming a solar abundance of Fe. The line is found to be broadened with a best fit value of line width $\\sigma \\sim$37$^{+24}_{-23}$ eV that corresponds to about 1730$^{+1130}_{-1050}$ km s$^{-1}$.  The broadening in this line can be due to thermal broadening, Doppler broadening or Compton scattering in the material responsible for reflection.  The natural width of this line is few eV \\citep{george91} and any broadening due to the thermal motion of the emitting atoms is 0.4(T/$10^{6}$)$^{1/2}$ eV, which is very small compared to the resolution of the HEG.  Hence the observed broadening can only be due to either the Doppler effect, Compton scattering or both.  V426~Oph shows a much weaker fluorescent line emission with an EW of 39$^{+22}_{-18}$ eV when compared with the previously reported value of 185$\\pm$40 eV by \\cite{baskill05} using \\emph{ASCA} data taken during 1994 September 18.  However, they found that the inclusion of a reflection continuum reduced the plasma temperatures and the EW of the 6.4 keV line.  SS~Cyg, with the highest signal-to-noise ratio, shows a prominent line at 6.4 keV with EW of 59$^{+25}_{-20}$ eV.  Previous \\emph{ASCA} and \\emph{Ginga} data of SS~Cyg show the presence of significant reflection in their spectra \\citep{done97}.  This line is also present in SS~Cyg during outbursts with a similar EW but their line centers are red-shifted by 50 eV.  This shift corresponds to a very high velocity of about 2300 km s$^{-1}$ during the outburst.  An outflowing wind with such a high velocity has been reported from SS~Cyg during an outburst \\citep{mauche04}.  Also the presence of a high velocity wind in CVs has been reported previously by \\cite{prinja95} using high resolution IUE observations. We suggest that the observed high velocity probably arises from the wind that is moving away from the system during the outburst.  \\cite{ezuka99} have shown that apart from reflection from the white dwarf surface, N$_{H}$ along the line of sight also contributes to the EW of the fluorescent iron line.  The observed value of N$_{H}$ for these sources vary in the range of 10$^{20-22}$ cm$^{-2}$ \\citep{mukai03,homer04,perna03,pandel05}.  If N$_{H}$ is the only contributor to the fluorescent Fe line, then for such a low value of N$_{H}$, the expected value of EW is $\\lesssim$10 eV assuming that the source is surrounded by a uniform absorber at a 4$\\pi$ solid angle. However, the observed best fit value of the EW varies in the range of 23--162 eV (see Table \\ref{tbl-2}) indicating the dominance of the reflection process over the absorption, resulting in the observed fluorescent Fe line.  \\subsection{Quiescence vs. Outburst States} As mentioned in \\S 2, two of the six sources, viz., U~Gem and SS~Cyg, are observed in both quiescence and outburst.  Such observations provide an opportunity to study the response of the system with changes in its intensity.  As shown in Figs. \\ref{fig1} \\& \\ref{fig2}, the three Fe emission lines have considerably different profiles during the two states. In particular, the profile of the Fe XXV triplet is significantly different during the two states.  In U~Gem during quiescence, this line has a double humped profile reflecting the dominance of the $r$ and $f$ lines over the $i$ components.  During the outburst, the line shows a broad symmetric Gaussian profile indicating the presence of high velocity material near the white dwarf. The common line width ($\\sigma$) for $r$, $i$ and $f$ lines is 55$^{+13}_{-19}$ eV, which corresponds to a velocity of 2460$^{+580}_{-850}$ km s$^{-1}$.  A ``flat-top'' profile is observed for SS~Cyg during an outburst. Such a profile has been reported for N VII line in O-type star, $\\zeta$~Puppis using \\emph{Chandra} HETG data and attributed to small velocity gradient at the larger radii in the outflowing winds in the star \\citep[see][]{cassi01,kahn01}. The failure to find a common shift for the $r$, $i$ and $f$ lines during the outburst, and the unusual result of a blue-shifted resonance component and red-shifted $i$ and $f$ components, is most likely just symptomatic of the non-gaussian shape of the line profile. This could be due to velocity-smeared bulk motion of material present either in a Keplerian orbit around the white dwarf or a wind flowing away from the system.  \\cite{mauche04} reported the presence of a high velocity wind ($v\\approx$2500 km s$^{-1}$) in SS~Cyg during an outburst observed with the Low Energy Transmission Grating (LETG) data.  On the other hand, the spectra of SS~Cyg in quiescence do not show any shift or broadening in the Fe XXV triplet.  Flat top or a double peaked line profiles have also been observed in active galactic nuclei, and interpreted as due to the material flowing away from the central object.  For example, a Seyfert-1 Galaxy Mrk~509 has been observed to show such a line profile from outflowing gas \\citep{phillips83}.  Under favorable viewing conditions one can see both the material moving away and approaching towards the observer that can produce such a line profile. However, the present data cannot resolve this issue. It is also possible that some higher-order effects like optical depth and radiation transfer may be playing a role in producing the observed line profile. In addition, the exposure time of 36 and 60 ksec (see Table~\\ref{tbl-1}) of SS~Cyg during 12 and 14 September 2000 covers $\\sim$1.5 and 2.5 orbital cycles with an orbital period of $\\sim$6.6 hrs.  The discordant wavelength shift is probably caused by the long exposure, during which the conditions in the emitting region might have changed significantly. The summed line may represent a merger of a broad range of physical conditions.  Detailed radiative transfer calculations are required to account for the higher-order effects, which is not within the scope of the present paper.  It has been observed during an optical outburst of SS~Cyg that the hard X-ray flux ($>$3 keV) is suppressed and the spectrum softens, whereas the X-ray flux is high and the spectrum is harder during the optical quiescence \\citep[][and references therein]{mcgowan04}.  The observed X-ray flux in SS~Cyg in the HEG band of 0.4--8 keV is 1.8 $\\times$ 10$^{-10}$ ergs cm$^{-2}$ s$^{-1}$ during quiescence and it is 4.0 $\\times$ 10$^{-11}$ and 4.8 $\\times$ 10$^{-11}$ ergs cm$^{-2}$ s$^{-1}$ during the outburst of 12 and 14 September 2000, respectively.  This suggests that for SS~Cyg during an optical outburst the X-ray flux is suppressed by a factor of $\\sim$3--4 supporting the previous observations.  On the other hand U~Gem shows the opposite behavior, when it goes in to optical outburst its X-ray flux in 0.4--8 keV range increases by a factor of 2--3.  \\subsection{Comparison with Magnetic CVs} Prominent Fe K$_{\\alpha}$ emission lines have been observed in magnetic CVs (MCVs).  The Fe XXV and Fe XXVI lines are believed to originate in a hot plasma in the post-shock region and the appearance of the fluorescent Fe line is attributed to reflection from the white dwarf surface or the cool pre-shock material \\citep{ezuka99}.  A study of Fe K${\\alpha}$ complex in magnetic CVs has been reported by \\cite{hellier04} using the \\emph{Chandra} HETG data.  They use a power law and three Gaussian components to model the Fe K$_{\\alpha}$ emission lines in a sample of five MCVs.  They find that the Fe XXVI line is red-shifted by 260 km s$^{-1}$ from a simultaneous fit to the spectra of five MCVs and suggested that the presence of the satellite lines can cause the small observed red-shift.  For non-magnetic CVs, we have not observed any detectable shift in the line energy of the Fe XXVI line during both quiescence and outburst suggesting an absence of any Doppler shift in these systems.  \\cite{hellier04} find that the Fe XXVI line is, in general, broadened by 1000 km s$^{-1}$ in MCVs.  For non-magnetic CVs, only U~Gem shows a significant broadening in Fe XXVI line by 1400$^{+840}_{-600}$ km s$^{-1}$ in quiescence.  \\cite{hellier04} also reported that the two intermediate polars (IPs), AO~Psc and EX~Hya, have a much stronger resonance line of the Fe XXV triplet than is observed in the polar AM~Her.  In our study, this line is also found to be strong like that observed in the two IPs in which the accretion takes place via a partial or truncated disk.  The common thread between the IPs and DNe is an accretion disk that is absent in polars.  This might provide some clue for the dependence of the Fe XXV resonance line strength on the presence of accretion disk or the strength of magnetic field.  Observations of a large number of CVs with a high resolution X-ray instrument can help to investigate this further.  \\cite{ezuka99} find that the reflection of hard X-rays from the white dwarf generally makes a significant contribution to the fluorescent iron line in MCVs having a weak line-of-sight absorber (N$_{H}<$10$^{22}$ cm$^{-2}$).  This suggests that in these systems the production mechanism of the fluorescent Fe line is independent of the magnetic nature of the system.  \\cite{hellier04} have detected a red wing of the fluorescent Fe line extending to 6.33 keV in the IP GK~Per. This broadening corresponds to a Doppler shift of up to 3700 km s$^{-1}$.  They suggest that this arises from the pre-shock material that is free-falling onto the white dwarf.  We have detected a symmetrically broadened fluorescent Fe line in V603~Aql with a broadening of $\\sim$1700 km s$^{-1}$, which is probably due to Doppler broadening or Compton scattering of material in the accretion disk. The red-shift of up to 2300 km s$^{-1}$ detected in SS~Cyg during outburst is attributed to the wind flowing away from the system unlike the free-falling material in MCVs."
    },
    "0512/astro-ph0512042_arXiv.txt": {
        "abstract": "We use the average \\EBV\\ and \\znii column densities of a sample of $z\\sim1$ Ca~{\\sc ii}~$\\lambda\\lambda 3935, 3970$ absorption line systems selected from the fourth data release of the Sloan Digital Sky Survey (SDSS)to show that on average, with conservative assumptions regarding metallicities and dust-to-gas ratios, they contain column densities of neutral hydrogen greater than the damped Lyman-$\\alpha$ (DLA) limit.  We propose that selection by \\caii absorption is an effective way of identifying high column densities of neutral hydrogen, and thus large samples of DLAs at $z_{abs}\\lsim1.3$ from the SDSS.  The number density of strong \\caii absorbers (with rest-frame equivalent width $W_{\\lambda 3935} \\geq 0.5$\\,\\AA), is $\\sim 20-30\\%$ that of DLAs, after correcting for the significant bias against their detection due to obscuration of the background quasars by dust.  On average these absorbers have \\EBV\\,$ \\gsim 0.1$\\,mag; the dustiest absorbers show depletions of refractory elements at a level of the largest depletions seen in DLAs.  For the first time we can measure the dust-to-metals ratio in a sample of absorption selected galaxies, and find values close to, or even larger than, those observed locally.  All of these properties suggest that a substantial fraction of the \\caii absorbers are more chemically evolved than typical DLAs.  There is a trend of increasing dust content with $W_{\\lambda3935}$; this trend with strong-line equivalent width is also observed in an equivalent, but much larger, sample of \\mgii absorbers.  Such a trend would result if the dustier systems are hosted by more massive, or disturbed, galaxies.  Follow-up imaging is required to provide conclusive evidence for or against these scenarios.  From consideration of the \\EBV\\ distribution in our sample, and assuming \\caii absorbers represent a subset of DLAs, we calculate that dust obscuration causes an underestimation in the number density of DLAs by at least $8 - 12$\\% at these redshifts.  Finally, the removal of Broad Absorption Line (BAL) quasars from the SDSS quasar sample increases the sensitivity of the detection of reddening by intervening absorbers.  To this end, we describe a new, automated, principal component analysis (PCA) method for identifying BAL quasars. ",
        "introduction": "Measurements of the chemical composition of galaxies play an important role in understanding galaxy formation and evolution.  The study of galaxy metallicities is closely intertwined with the question of their dust content: systems richer in metals have the greater potential to form dust grains. These grains selectively deplete metals in the interstellar medium (ISM) of the galaxies, redden their spectral energy distribution and cause an overall extinction of the light.  The uncertainty in dust composition and distribution in galaxies affects the interpretation of many extragalactic observations.  Quasar absorption line systems provide a powerful probe of metals in the ISM of galaxies at high redshift, unhampered by luminosity bias, through the absorption of the light from a background quasar by gas phase metal ions.  Notably, reddening due to dust has failed to be detected in samples of the highest column density absorption line systems, damped Lyman-$\\alpha$ systems (DLAs), at $z\\sim2-4$ \\citep{2004MNRAS.354L..31M,corals_ebv}. The relative paucity of dust particles in DLAs is backed up by observations of absorbers in radio selected quasar spectra, which indicate that extinction due to dust is a small effect \\citep[CORALS survey][]{2001A&A...379..393E}, and metal abundances show DLAs to be generally metal- and therefore dust-poor \\citep[e.g.][]{2004cmpe.conf..257P}.   \\subsection{\\caii in the local universe} This paper investigates the dust and metal properties of a new class of absorption line system, selected via the \\caii$\\lambda\\lambda3935,3970$ doublet\\footnote{Vacuum wavelengths are used throughout this paper.}  from the Sloan Digital Sky Survey (SDSS). These are the familiar K and H lines of the solar spectrum.   Ca is an $\\alpha$-capture element thought to be produced mainly by Type II supernovae.  The convenient positioning of the resonance lines of its singly-ionised ion in the violet portion of the optical spectrum has made it accessible to stellar and interstellar studies for over a century, with major Galactic surveys by \\citet{1949ApJ...109..354A} and \\citet{1972ApJ...173...43M}.  The inference of Ca abundance from \\caii absorption lines is however compromised by two considerations.  First, with an ionisation potential of 11.9\\,eV, lower than that of H~{\\sc i}, Ca$^+$ is a minor ionisation stage of Ca in the neutral ISM---most of the Ca is doubly ionised.  Second, Ca is one of most depleted elements in the ISM---typically more than 99\\% of all Ca is `hidden from view' having been incorporated into dust grains \\citep{1996ARA&A..34..279S}. The degree of depletion is expected to depend significantly on both the density of the gas (not just $N$(H~{\\sc i})) and on the presence of shocks, which may result in grain destruction, returning Ca to the gas phase. However, as stressed by \\citet{2000ApJ...544L.107W}, the dispersion in the measured gas-phase abundance of \\caii at fixed N(HI) in the Milky Way is remarkably small.  Although the precise interpretation of the detection of a significant column density of \\caii is complex, it is clear that a large column density of neutral hydrogen is implied, given the high degree of depletion and low ionisation energy.  \\caii absorption in quasar spectra due to low redshift galaxies was detected early on in the history of quasar absorption line studies; for a review of early results see \\citet{1988qsal.proc..147B}. The most comprehensive study was carried out by \\citet{1991MNRAS.249..145B} who concluded that at projected separations $\\gsim 20\\,$kpc from the centres of visible galaxies, \\caii absorption is patchy and, where it does occur, relatively weak, with the rest frame equivalent width of the strongest member of the \\caii doublet $W_{\\lambda 3935}\\lsim 0.2\\,$\\AA.  Examples of strong \\caii absorption ($W_{\\lambda 3935}\\gsim0.5\\,$\\AA) were mostly confined to sightlines with projected separations $\\lsim 10\\,$kpc. Evidence for mergers was seen in those cases with large projected separations and large column densities of \\caii \\citep{1991MNRAS.251..649B}.  Since the early 1990s the focus of quasar absorption line research moved to high redshift and \\caii was largely ignored due to both its rarity and wavelength. At $z>1$ the doublet moves into the near infrared which, at that time, was difficult to access.  Recently, very strong \\caii absorption has been detected in a small number of highly reddened, \\EBV$\\sim 1$, sightlines to quasars \\citep{2000A&A...359..457P,2002ApJ...575L..51H,2005ApJ...622L.101W}. Such detections confirm the link between high gas column densities, dust and \\caii absorption in relatively extreme circumstances when the line-of-sight is almost coincident with the centre of a galaxy.   \\subsection{Metal absorbers and DLAs}  DLAs are quasar absorption line systems with the highest column densities of neutral hydrogen, with a nominal column density limit of \\nhi $>2\\times10^{20}$cm$^{-2}$.  Although they have been studied extensively since the seminal paper by \\citet{1986ApJS...61..249W} twenty years ago, their precise nature, evolution and relation to the population of galaxies as a whole continue to be the subject of much discussion \\citep{wolfe05}. Interest in DLAs results primarily from their apparent domination of the neutral gas mass fraction of the Universe at high redshift. Such cold, neutral gas is required for later star formation.  The defining diagnostic of a DLA, the Lyman-$\\alpha$ absorption line, does not enter the optical atmospheric window until redshift $z_{\\rm abs} \\sim 1.8$. At $z\\sim 0$ DLAs can be identified through 21\\,cm emission \\citep[e.g.][]{zwaan05}. The need for space-based observations has made it difficult to obtain large samples of intermediate redshift DLAs which are crucial both to clarify the relation of DLAs to luminous galaxies by direct imaging of the DLA hosts (which is much easier at redshifts $z \\lsim 1$), and to follow the evolution of their properties over the course of time.  \\citet{2000ApJS..130....1R} proposed a useful prognostic for intermediate redshift DLAs: the equivalent widths of the strong \\mgii$\\lambda2796$ and \\feii$\\lambda2600$ absorption lines can be used to identify DLAs with a $\\sim$40\\% success rate \\citep{astro-ph/0505479}.  A complementary method to that proposed by Rao \\& Turnshek for recognising DLAs at intermediate redshift is through the detection of absorption lines which are intrinsically weak, due to low cosmic abundance, low transition probability, ionisation state or affinity for dust grains.  By choosing lines appropriately, the success rate of identification can approach 100\\%, circumventing to some extent the need for UV spectroscopy; for example, the detection of the \\znii$\\lambda\\lambda2026, 2062$ doublet virtually guarantees the DLA nature of an absorption system \\citep{1990ApJ...348...48P}.  As highlighted above, little is known of the strength of \\caii lines in DLAs; here we propose that approximately one fifth to one quarter of DLAs have \\caii lines with rest frame equivalent widths exceeding 0.5\\,\\AA\\ for the stronger member of the doublet.  The Sloan Digital Sky Survey (SDSS) quasar catalogue, with more than 45\\,000 spectra in its third data release \\citep[DR3,][]{astro-ph/0503679}, provides the ideal database for obtaining large samples of metal line systems over a wide redshift range.  The ability to detect significant numbers of DLAs using the SDSS over the entire redshift range $0\\lsim\\zabs<1.3$ would provide a much needed injection of new intermediate to low redshift DLAs for imaging and chemical evolution studies.     \\subsection{Aims of this paper}  In a recent paper \\citep[][hereafter Paper I]{2005MNRAS.361L..30W}, we assembled a sample of 31 \\caii absorption systems selected from a high signal--to--noise ratio subset of the SDSS DR3 at redshifts $0.84 < \\zabs <1.3 $ and reported the detection of a significant degree of reddening---corresponding to a colour excess \\EBV$\\, = 0.06$---of the background quasars by dust in these absorbers.  In this follow-up paper, we analyse the relative abundances of refractory elements in these intermediate redshift \\caii absorbers with a view to establishing their connection to the DLA population.  Specifically, we construct composite spectra which allow us to measure the relative abundances of Zn~{\\sc ii}, Cr~{\\sc ii}, Fe~{\\sc ii}, Mn~{\\sc ii}, Ti~{\\sc ii} and Ca~{\\sc ii}.  The Zn measurement is crucial since, among the elements considered, it is the only one which shows little affinity for dust \\citep{1996ARA&A..34..279S} and therefore gives a standard against which to compare the gas-phase abundances of the others.  We add to the sample of Paper~I six additional \\caii systems identified from the fourth data release (DR4) of the SDSS (Adelman-McCarthy et~al.  2005)\\nocite{DR4}, taking the total number of \\caii absorbers to 37.  27 of these have rest frame spectra that encompass the region of the \\znii doublet. We also re-examine in more detail the reddening results of Paper~I in light of the abundance determinations presented here.   The paper is organised as follows.  In Section 2 we describe our search criteria for identifying absorption line systems.  Strong \\mgii $\\lambda2796$ and \\feii $\\lambda2600$ absorption lines are seen in all \\caii absorbers; we compare their equivalent width distributions with those of \\mgii absorption systems at similar redshifts in SDSS quasars, paying particular consideration to the DLA selection criteria of \\citet{astro-ph/0505479}.  In Section 3 we calculate the redshift path of our \\caii survey. In Section 4 we derive column densities and relative abundances of elements detected in composite spectra of \\caii absorption line systems; we compare the abundance patterns with those seen in known DLAs and in the interstellar medium of the Milky Way.  Section 5 deals with the reddening introduced by different samples of \\caii and \\mgii absorbers on the observed spectral energy distributions (SEDs) of the background quasars.  We discuss our results in Section 6, focusing in particular on the use of \\caii in the identification of DLAs, on the metallicities and dust content of Ca~{\\sc ii}-selected DLAs at $z \\sim 1$, and on the implications of our findings for the more general issue of quasar obscuration by dust-rich intervening absorbers. Appendix A gives details of our new method for recognising quasars that show evidence of broad absorption line systems (BALs). ",
        "conclusions": "In this paper we have extended the study of intermediate redshift ($z_{\\rm abs}\\sim 1$) absorption systems selected via \\caii equivalent width which was begun by \\citet{2005MNRAS.361L..30W} by: (a) expanding the sample of \\caii absorbers to a total of 37 with new systems selected from the SDSS DR4; (b) measuring the relative abundances of refractory elements from composite spectra of a subset of 27 systems at redshifts such that the \\znii$\\lambda\\lambda 2026, 2062$ doublet is covered in the SDSS spectra; and (c) comparing the degree of reddening imposed on the spectra of background quasars by dust in intervening \\caii absorbers to that caused by Mg~{\\sc ii}-selected DLAs.  We have confirmed the detection of of significant reddening, at the level of \\EBV\\,$\\simeq 0.06$, associated with the \\caii absorption systems. This is particularly noteworthy considering that much larger samples of Mg~{\\sc ii}-selected DLAs at the same redshifts and of Lyman~$\\alpha$-selected DLAs at higher redshifts have all failed so far to show a clear reddening signal.  Only Mg~{\\sc ii}-selected DLAs with the strongest metal lines approach the reddening of the \\caii absorbers, once allowance is made for the fraction of sub-DLAs in the sample.  Given the lack of direct measurements of $N$(H~{\\sc i}) for the foreseeable future, we have considered various lines of indirect evidence which all point towards the conclusion that our absorption systems with a \\caii$\\lambda 3935$ equivalent width $W_{\\lambda 3935} \\gsim 0.68$\\,\\AA\\ are a subset of the DLA population, and those with lower equivalent widths are likely to be DLAs or lie close to the nominal DLA column density limit: \\begin{enumerate} \\item The strengths of \\mgii$\\lambda 2796$ and \\feii$\\lambda 2600$ absorption lines fulfil the DLA selection criteria of \\citet{astro-ph/0505479} in all but two cases (which are borderline---see Fig.~\\ref{fig:mgiifeii}). \\item The ratios of the refractory elements Cr, Fe, Ti and Mn to the undepleted Zn all span similar ranges to those in confirmed DLAs, with the strongest \\caii systems exhibiting some of the most pronounced degrees of depletion encountered so far (Fig.~\\ref{fig:dla}). \\item With the conservative assumption of solar metallicity, the measured \\znii column densities imply an average \\lognhi\\,$= 20.19$, just below the threshold of \\lognhi\\,$= 20.3$ considered to be the definition of a damped system. Even a moderate, and more realistic, underabundance of Zn and other Fe-peak elements in these systems would bring them within the conventional DLA samples. \\item In itself, the magnitude of reddening detected implies \\lognhi\\,$ \\ge 20.43$ adopting the conservative assumption of a Galactic gas-to-dust ratio. \\end{enumerate} Furthermore, these are likely to be underestimates for the population of \\caii absorbers as a whole, since our magnitude limited quasar survey certainly misses a greater fraction of absorbers with the largest values of \\EBV.  After correcting for our detection efficiency, including dust obscuration bias, we calculate that \\caii absorbers with $W_{\\lambda 3935} > 0.5$\\,\\AA\\ have a number density per unit redshift $n(z) \\sim 0.022$ at $\\langle z_{\\rm abs} \\rangle = 0.95$, 20--30\\% of that of Lyman~$\\alpha$-selected DLAs at the same redshifts.  This in turn implies that the regions within galaxies giving rise to strong \\caii absorption have small dimensions, with radii of 7--8\\,$h^{-1}$\\,kpc. Possibly, these are the inner regions where the highest column densities $N$(H~{\\sc i}), gas densities $n_{\\rm H}$, and/or metallicities (Zn/H) are encountered and where the characteristics of DLAs begin to approach those of the galaxies detected directly via their UV stellar light.  Even without a knowledge of their neutral hydrogen column densities we can deduce the typical dust-to-metals ratio $\\cal R_{\\rm DM}$ of Ca~{\\sc ii}-selected DLAs from their measured columns of metals (via Zn~{\\sc ii}) and dust [via \\EBV].  We find high values of $\\cal R_{\\rm DM}$, as high or higher than the typical ratio in the Milky Way, although the errors affecting our estimates are still large. This may be a further indication that \\caii absorbers are at an advanced stage of chemical evolution if the dust-to-metals ratio correlates with metallicity.  Interestingly, we find a clear trend of increasing reddening and element depletions with \\caii equivalent width; a positive correlation of \\EBV\\ with line strength is also present among the Mg~{\\sc ii}-selected DLA candidates.  Since these strong lines are generally saturated, their equivalent widths depend more on the velocity dispersion of the absorbers than on their column densities. Thus, these trends may indicate that the systems with the larger values of \\EBV \\ are either more massive or undergoing some form of disturbance.  Whatever their exact nature, it is clear that Ca~{\\sc ii}-selected systems are an interesting class of quasar absorbers which have been relatively neglected until now and yet merit considerable attention. Their relation to the DLA population at large will be clarified by measuring the \\caii lines associated with confirmed DLAs; recent improvements in detector and spectrograph efficiency at far-red and near-IR wavelengths now make the \\caii$\\lambda\\lambda3935,3970$ lines more easily accessible to high resolution spectroscopy at redshifts $z_{\\rm abs} > 1$.  Such measurements will establish the relation between $N$(H~{\\sc i}), \\caii equivalent width and dust content, and will provide firm estimates of the metallicities and gas-to-dust ratios which are lacking from the current analysis.  At $z \\sim 1$ imaging is considerably easier than for the bulk of known DLAs at $z > 2$ and it should be relatively straightforward to identify the objects responsible for the \\caii absorption. In this respect, our sample of 37 \\caii absorbers represents a major injection of new data into a field of study which up to now has been severely limited by the small number of known DLAs at $z \\sim 1$. The images will establish whether the absorption does take place in the inner regions of galaxies which have undergone significant chemical evolution, as we propose, or whether DLAs with \\caii absorption single out major mergers \\citep{1991MNRAS.251..649B}. Future papers in this series will investigate further both the relation between \\caii absorbers and DLAs, and the properties of the host galaxies of \\caii absorbers through follow-up observations.  Ultimately, the combination of absorption line spectroscopy and deep imaging will clarify the link between absorption- and emission-selected galaxies and thereby help constrain galaxy evolution models. The \\caii absorbers may well turn out to be the key to understanding how different classes of high redshift galaxies fit into a unified picture."
    },
    "0512/gr-qc0512005_arXiv.txt": {
        "abstract": "After explaining the physical origin of the quasinormal modes of perturbations in the background geometry of a black hole, I critically review the recent proposal for the quantization of the black-hole area based on the real part of quasinormal modes. As instantons due to the barriers of black-hole potentials lie at the root of a discrete set of complex quasinormal modes frequencies, it is likely that the physics of quasinormal modes can be learned from quantum theory. I propose a connection of a system of quasinormal modes of black holes with a dissipative open system, in particular, the Feshbach-Tikochinsky oscillator. This argument is supported in part by the fact that these two systems have the same group structure $SU(1,1)$ and the same group representation of Hamiltonians; thereby, their quantum states exhibit the same behavior. \\\\ \\noindent Keywords: Quasinormal modes, Black Hole, Dissipative open system ",
        "introduction": "In linear perturbation theory, a perturbed black hole emits waves that are outgoing to spatial infinity and the event horizon. The wave function with this boundary condition is called a quasinormal mode of the black hole and has a complex frequency, whose imaginary part leads to a decaying behavior. There are several motivations to study quasinormal modes of black holes. The original motivation is that quasinormal modes carry carry information of black hole, such as mass, charge, and angular momentum, and may be observed by a gravitational wave detector \\cite{vishveshwara}. Another motivation comes from a recent argument that the real part of the quasinormal mode frequency may play a role in explaining the quantization of the black-hole area \\cite{hod,kunstatter,horowitz}. The quantization of the black-hole area, which is believed to be closely related with the microscopic origin of black-hole entropy, is also explained in loop quantum gravity \\cite{dreyer}. Ideas have been suggested to understand the origin of quasinormal modes and various methods have been developed to find quasinormal modes (for review and references, see Refs. \\cite{kokkotas,natario}). Quite recently, I have suggested that the quantum theory of quasinormal modes might be related with a dissipative open system and possibly with black-hole thermodynamics \\cite{kim0,kim1}.   The purpose of this paper is {\\it not} to develop any new method for finding analytically quasinormal modes, {\\it but} to exploit the physical interpretation of quasinormal modes of black holes. In particular, I shall provide some supporting arguments for my recent proposal on the connection of quasinormal modes with a dissipative open system, the Feshbach-Tikochinsky (FT) oscillator, and possibly with the thermodynamics of black holes in a semiclassical approach. There are some open questions about quasinormal modes and thermodynamics. First, is there any connection between a system of quasinormal modes and a dissipative open system? Second, is there any connection between quasinormal modes and black-hole thermodynamics? It is known that perturbations of a system may provide some information of the system, and, in some cases, may carry all information. Then, what black-hole physics can we learn from quasinormal modes? ",
        "conclusions": "The quasinormal modes of black-hole perturbations are decaying wave functions that are outgoing to spatial infinity and the event horizon. Through a study of exactly solvable potentials as approximations for black-hole potentials, the physical origin of quasinormal modes is shown to be the single and multi-instantons of the black-hole potentials. The instanton is a periodic solution of the inverted potential in Euclidean time and corresponds to a bound state. The recent proposal that the real parts of quasinormal-mode frequencies are related with the quantization of the black-hole area may suggest deep physical implications of quasinormal modes in black-hole physics. However, the complete understanding of quasinormal modes should include not only the real parts but also the imaginary parts of the quasinormal frequencies. It will be interesting to investigate the origin of damping factors due to the discrete set of the imaginary frequencies. As instantons responsible for discrete complex frequencies are due to the quantization of periodic motions in an inverted potential, it is likely that a semiclassical approach may provide better comprehension of quasinormal modes than a classical approach.   The decaying behavior of quasinormal modes is a characteristic aspect of dissipative systems. In this paper, I proposed the quantization of quasinormal modes and a relation with the FT oscillator, a dissipative open system. The canonical Hamiltonian for each quasinormal mode has the two-mode representation of the group $SU(1,1)$. There is an amplified mode in addition to a damped mode coming from the completeness of mode solutions in quantizing a field. The imaginary frequency provides an interaction that connects the damping mode with the amplified mode. The canonical Hamiltonian of a quasinormal mode has the same group representation as the FT oscillator. In the FT oscillator, the damping constant provides the interaction between the damped and the amplified modes. The energy of the damped mode is transferred to the amplified mode; thus, the total energy is conserved. The quantum theory of the FT oscillator, which has been intensively studied, is expected to give useful information on black-hole physics via quasinormal modes. The quantum theory of the FT oscillator and/or the quasinormal mode of a black hole is nonunitary because it is a dissipative open system. An interesting model in this direction is provided by three-dimensional BTZ black holes which allow exact solutions and have AdS/CFT correspondence \\cite{myung}. It is shown there that a nonrotating BTZ black hole having quasinormal modes is nonunitary and, thus, belongs to a dissipative system. This nonunitarity is in agreement with conformal field theory \\cite{cardoso,birmingham}. On the other hand, an extreme BTZ black hole or pure AdS spacetime has only real quasinormal-mode frequencies and, thus, is unitary. This implies that quasinormal modes are related with a dissipative system and nonunitarity.  Finally, it would be very interesting to find a connection between quasinormal mode and black-hole thermodynamics, if any. From the arguments in this paper, that is highly likely. This point is under study.   {\\it Note added}. After submission of this paper, I was informed by several authors of many relevant references. Berti {\\it et al} pointed out in Refs. \\cite{berti1,berti2} that the real parts of the highly damped quasinormal modes of the Kerr black hole do {\\it not} have the asymptotic frequencies suggested by Hod. This fact, however, does not significantly change the arguments of this paper because the imaginary parts of the quasinormal modes are related with dissipative open systems in the quantization of the quasinormal modes. Padmanabhan {\\it et al} used the first Born approximation to calculate the scattering matrix of the quasinormal modes, from which the discrete complex frequencies are derived \\cite{padmanabhan1,padmanabhan2}. Setare {\\it et al} also discussed the black-hole area quantization for a non-rotating BTZ black hole, an extremal Reissner-Nordstr\\\"{o}m black hole, and an extremal Kerr black hole \\cite{setare1,setare2,setare3}."
    },
    "0512/nucl-th0512095_arXiv.txt": {
        "abstract": "s{The effects of vector-isovector terms predicted by relativistic mean field models based on effective field theory  on the equation of state and the neutrino mean free path in neutron stars have been studied. Using a procedure similar to Ref.\\,\\cite{pieka}  in treating the isovector sector, a  parameter set (G2*) that predicts a much smaller proton fraction than the standard one (G2) can be obtained. We found that the G2* has a softer equation of state compared to the G2 at higher densities and the disappearance of the anomalous behavior of the neutrino mean free path in neutron star does not depend on the proton fraction.} ",
        "introduction": "The relativistic mean field (RMF) model allows ground state properties of finite nuclei to be described with relatively high precision (for a review see, e.g., Ref.\\,\\cite{pg}). If we extrapolate this model to the high density regime, it is known that the RMF model predicts a much stiffer equation of state (EOS) compared to the Bethe Brueckner Goldston (BBG) ``data''~\\cite{pg,brock}, variational calculation of Akmal {\\it et ~al}.\\,\\cite{akmal}, Dirac Brueckner Hartree Fock (DBHF)~\\cite{li}, non relativistic Brueckner Hartree Fock (BHF) with AV14 potential plus 3BF~\\cite{baldo}, or heavy ion experimental data~\\cite{daniel}, in the region where baryon density is around $2.0\\rho_0-4.5\\rho_0$, with $\\rho_0$ the nuclear matter saturation density. This model also predicts a too large proton fraction ($Y_p$)~\\cite{parada} which leads to a too low threshold density to start the direct URCA cooling process in neutron star. Furthermore, the predicted neutrino mean free path (NMFP) shows an anomalous behavior. On the other hand, there are parameter sets of RMF model which are specifically parameterized for interstellar purpose (e.g., parameter sets of Refs.\\,\\cite{pieka,glen}). This kind of parameter sets predicts a soft EOS at high densities but the predicted finite nuclei ground state properties are quite unsatisfactory. The possible reason is because the parameter sets with soft EOS have a relatively large nucleon effective mass ($M^*$) for finite nuclei density, which has a consequence that their spin-orbit splittings are relatively narrow compared with experimental data~\\cite{anto}. This means that the  RMF model is not ``effective'' enough to cover a wide range of densities and the reason lies in the form of the nonlinear self coupling~\\cite{brock}.  In 1996 the chiral effective Lagrangian model was proposed by Furnstahl, Serot and Tang~\\cite{fst}, whose mean field treatment is hereafter called the ERMF model. This model has accurate predictions of the ground state properties of nuclei and the extrapolation to high densities yields a soft EOS which is consistent with other calculations~\\cite{brock,akmal,li,baldo}, as well as experimental data of Refs.\\,\\cite{daniel,amu}. Nevertheless, it is widely known that this model still predicts a too large $Y_p$ in the neutron star~\\cite{parada}. This is caused by the role of isovector terms, in which the corresponding parameters  are poorly constraint by insensitive isovector observables of finite nuclei.  Therefore, in this report, we adjust the isovector-vector channel of the ERMF model in order to have a reasonable neutron star $Y_p$ without changing its ground state predictions in finite nuclei by fine tuning the symmetry energy ($E_{\\rm sym}$) of the symmetric nuclear matter (SNM) to achieve a softer $E_{\\rm sym}$ at high densities. We follow the prescription of Horowitz and Piekarewicz~\\cite{pieka}  for the adjustment procedure. Besides that, we also study the anomalous behavior of the predicted NMFP in the neutron star. ",
        "conclusions": "The standard RMF model needs additional nonlinear terms in order to simultaneously produce accurate ground state predictions of finite nuclei and realistic description in the high densities regime. The ERMF model has a strong theoretical ground to meet this requirement through the language of effective field theory. The reason that the ERMF model can cover a wide range of densities originates from the fact that the model has $M^* \\sim 0.6 M$ at the nuclear saturation density, but a large $M^*$  and a soft $E_{\\rm sym}$ at high densities. A more detailed analysis can be found in Ref.\\,\\cite{iso_prc}."
    },
    "0512/astro-ph0512038_arXiv.txt": {
        "abstract": "The European {\\Gaia} astrometry mission is due for launch in 2011. {\\Gaia} will rely on the proven principles of ESA's \\textit{Hipparcos} mission to create an all-sky survey of about one billion stars throughout our Galaxy and beyond, by observing all objects down to 20th magnitude. Through its massive measurement of stellar distances, motions and multi-colour photometry it will provide fundamental data necessary for unravelling the structure, formation and evolution of the Galaxy. This paper presents the design and performance of the broad- and medium-band set of photometric filters adopted as the baseline for {\\Gaia}. The nineteen selected passbands (extending from the ultraviolet to the far-red), the criteria, and the methodology on which this choice has been based are discussed in detail. We analyse the photometric capabilities for characterizing the luminosity, temperature, gravity and chemical composition of stars. We also discuss the automatic determination of these physical parameters for the large number of observations involved, for objects located throughout the entire Hertzsprung-Russell diagram. Finally, the capability of the photometric system to deal with the main {\\Gaia} science case is outlined. ",
        "introduction": "photometric system} \\label{sec:BBPMBP}  In this section we describe in detail the baseline photometric system for {\\Gaia}, called C1, which consists of the C1B and C1M broad and medium passbands. The role of each of the passbands is described in relation to spectral features and astrophysical diagnostics.  \\subsection{The C1B broad passbands} \\label{sec:section_BBP}  The C1B component of the {\\Gaia} PS has five broad passbands covering the wavelength range of the unfiltered light from the blue to the far-red (i.e.\\ 400--1000~nm). The basic response curve of the filters versus wavelength is a symmetric quasi-trapezoidal shape.  The filters were chosen to satisfy both the astrophysical needs and the specific requirements for chromaticity calibration of the astrometric instrument (see Section~\\ref{sec:design_PS}). The C1B set of passbands evolved from the convergence of the \\citet{vytasbbp}, \\citet{lin2003a} and \\citet{jordibbp} proposals. The specifications of the filters are given in Table~\\ref{tab:filters5}. Figure~\\ref{fig:SED_BBP} shows the spectral response of the passbands.  The estimated end-of-mission precisions, computed as described in Section~\\ref{sec:aperture}, are shown in Fig.~\\ref{fig:errors_BBP}.  The Balmer discontinuity and the H$\\beta$ line limit the blue and red edges of the C1B431 filter, respectively. The reddest filter C1B916 is designed to measure the light between the Paschen jump and the red limit of the sensitivity of CCDs in the Astro focal plane. The filter C1B655 is centred on the H$\\alpha$ line and its width has been optimized together with C1M656 in C1M (see Section~\\ref{sec:section_MBP}). The blue and red limits of the two remaining filters (C1B556 and C1B768) are consequently set in order to provide full coverage of the wavelength range in $G$ (i.e.\\ avoiding gaps between passbands). Four passbands are enough for chromaticity calibration, but five passbands are preferred to four, when the classification and astrophysical parametrization of stars is considered. Since only four strips of CCDs are available in BBP, the two broadest passbands (C1B556 and C1B768) are implemented together in one strip. Hence, the number of observations with these two passbands will be half as much as with the other C1B passbands.  The current design of the payload foresees no ultraviolet sensitivity for the Astro/BBP instrument (see Section~\\ref{sec:instrument}).  The near ultraviolet is the most important for stellar classification. The Balmer jump is the feature in the spectra of B-A-F type stars most sensitive to the temperature and gravity. It also contains information on the metallicity of F-G-K type stars. The absence of an ultraviolet passband in BBP is compensated with the inclusion of a broad UV passband (C1M326) in MBP (see Section~\\ref{sec:section_MBP}).  The classification and parametrization of objects in {\\Gaia} is done using the BBP and MBP measurements together and therefore the lack of a UV passband in BBP should not be a drawback. However, since the Spectro instrument has a lower angular resolution than Astro, the combination of BBP and MBP data is not always possible.  BBP will be the only tool for classification of stars in the crowded fields with stellar densities larger than $200\\,000$--$400\\,000$ stars/sq~deg (see Section~\\ref{sec:crowded}). Such stellar densities are found in some areas of the bulge and of the disk \\citep{robin,drimmel2005} and most of these areas have low interstellar extinction (such as Baade's window). In dense areas the trigonometric parallax and the Paschen jump will provide luminosity parametrization.   \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig5.eps,width=1.0\\linewidth} \\end{center} \\caption{Response curves of the filters in C1B folded with the optics transmission and the QE of the CCDs. The spectral energy distributions of solar metallicity A0~V, G2~V and K7~V type stars in units of W m$^{-2}$ Hz$^{-1}$, taken from the BaSeL-2.2 library, are overplotted with (black) dotted lines (the vertical scale is arbitrary). The dashed (blue and violet) lines correspond to G2~V and K7~V type stars with $\\feh=-1$ (when compared with the dotted lines they show the effect of a change of \\feh). The dashed (red) line shows the change of the spectrum of an A0 type star due to a change of luminosity. The solid (orange) line is the N-type carbon star AW Cyg extracted from the \\citet{gunn} library of stellar spectra.} \\label{fig:SED_BBP} \\end{figure}  \\begin{table} \\caption{{\\small Specifications of the filters of C1B implemented in Astro.}} \\label{tab:filters5} \\begin{center} \\begin{tabular}{lccccc} \\hline Band       &     C1B431&C1B556&C1B655&C1B768&C1B916 \\\\ \\hline $\\lambda_\\mathrm{blue}$ (nm)  & 380  &  492  &  620  &  690  & 866 \\\\ $\\lambda_\\mathrm{red}$ (nm)   & 482  &  620  &  690  &  846  & 966 \\\\ $\\lambda_\\mathrm{c}$ (nm)     & 431  &  556  &  655  &  768  & 916 \\\\ $\\Delta \\lambda$ (nm) & 102  &  128  &   70  &  156  & 100 \\\\ $\\delta \\lambda$ (nm) &10,40 & 10,10 & 10,10 & 10,40 &40,10 \\\\ $\\epsilon$ (nm)      & 2,2  &  2,2  &  2,2  &  2,2  & 2,2 \\\\ $T_\\mathrm{max}$ (\\%)        &  90  &   90  &   90  &   90  &  90 \\\\ ${n_\\mathrm{strips}}$    &   1  &  0.5  &    1  &  0.5  &   1 \\\\ \\hline \\multicolumn{6}{l}{\\footnotesize{$\\lambda_\\mathrm{blue}$, $\\lambda_{\\rm red}$: wavelengths at half-maximum transmission}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$\\lambda_\\mathrm{c}$: central wavelength$= 0.5(\\lambda_{\\rm blue} + \\lambda_{\\rm red})$}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$\\Delta \\lambda$: FWHM}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$\\delta \\lambda$: edge width (blue, red) between 10 and 90\\% of $T_{\\rm max}$}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$\\epsilon$: manufacturing tolerance intervals centred on $\\lambda_\\mathrm{blue}$ and $\\lambda_\\mathrm{red}$}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$T_\\mathrm{max}$: maximum transmission of filter}}\\\\ \\multicolumn{6}{l}{\\footnotesize{$n_\\mathrm{strips}$: Number of CCD strips carrying the filter: C1B556 and C1B768}}\\\\ \\multicolumn{6}{l}{\\footnotesize{\\hspace{0.4cm}share one CCD strip}}\\\\ \\end{tabular} \\end{center} \\end{table}    \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig6.eps,width=1.0\\linewidth} \\end{center} \\caption{Estimation of the end-of-mission precisions for the five passbands in C1B as a function {of~~}$V$, computed according to equation (\\ref{eq:eqerror}) and taking $\\sigma_{\\rm cal}=0$~mag. In this paper, the contribution of calibration errors to the end-of-mission precision is estimated to be 3~millimag in every passband (horizontal line). The errors for C1B556 are all nearly the same for a given $V$ magnitude because the mean wavelengths of these passbands are very similar.} \\label{fig:errors_BBP} \\end{figure}   \\subsubsection {Astrophysical diagnostics}  The response of the C1B431 filter at the shortest wavelength is asymmetrical, with a red edge that is less steep. This is done to compensate the shift of the maximum of the response function redward due to the slope of the quantum efficiency curve and the reflectance curve for six silver surfaces. As a result, the response function of the C1B431 passband becomes similar to that of the $B$ passband of the {\\it UBV} system. The mean wavelengths of both passbands are also similar:  445~nm for C1B431 and 442~nm for $B$.  The mean wavelength and half width of the C1B556 response function is very similar to that of the $V$ passband.  As a result the colour index C1B431--C1B556 will be easily transformable to Johnson's $B$--$V$ and vice versa.  Analogously, the C1B768 passband can be easily related with the Cousins $I$, the SDSS $i'$ and the HST~814 passbands and the colour index C1B556--C1B768 is transformable to $V$--$I_C$, $r'$--$i'$ and HST~555--814. This will facilitate the comparison of the numerous ground-based investigations in the {\\it BV} system and the large number of observations being done in the far-red passbands with the {\\Gaia} results. As can be seen from Fig.~\\ref{fig:SED_BBP}, the differences of the `blue minus green', `blue minus red' or `blue minus far-red' magnitudes may serve as a measure of metallic-line blanketing, although with much less sensitivity than a colour index containing the ultraviolet. In the C1B431--C1B556 vs. C1B556--C1B768 diagram the deviations of F-G metal-deficient dwarfs and G-K metal-deficient giants from the corresponding sequences of solar metallicity are up to 0.07 and 0.20 mag, respectively.  The combination of the fluxes measured in the C1B655 passband and the narrow passband C1M656 in MBP form an H$\\alpha$ index primarily measuring the strength of the H$\\alpha$ line. The H$\\alpha$ index shares the same properties as the $\\beta$ index in the Str\\\"omgren-Crawford photometric system. It is an indicator of luminosity for stars earlier than A0 and of temperature for stars later than A3, almost independent of interstellar extinction and chemical composition. The same reddening-free index may be used for the identification of emission-line stars.  The two remaining red and far-red passbands, C1B768 and C1B916, give the height of the Paschen jump which is a function of temperature and gravity. Although the maximum height of the Paschen jump is 0.3 mag only, i.e., about four times smaller than the Balmer jump, it still provides the needed information if its height, C1B768--C1B916, is measured with high accuracy (not lower than 1\\%) and this is reachable up to about $V\\sim 17-18$ (see Fig.~\\ref{fig:errors_BBP}). The colour indices C1B556--C1B768 and C1B556--C1B916 for unreddened late-type stars may be used as indicators for the temperature. These colour indices (and C1B768--C1B916) also allow for the separation of cool oxygen-rich (M) and carbon-rich (N) stars.  The C1M326--C1B431 vs.\\ C1B431--C1B556 diagram has the same properties as the $U-B$ vs.\\ $B-V$ diagram of Johnson's system or similar diagrams for other systems \\citep{vytas1992}. Supergiants are well separated from the main-sequence stars. Metal deficient F-G dwarfs and G-K giants exhibit ultraviolet excesses up to $0.4$ mag. Blue horizontal branch stars show ultraviolet deficiencies up to $0.3$ mag, while white dwarfs are situated around the interstellar reddening line of O-type stars.  \\subsubsection{Chromaticity evaluation} \\label{sec:chroma}  Although no refracting optics is used for the astrometric field, the precise centre of a stellar image is still wavelength-dependent because of diffraction and its interplay with the optical aberrations of the instrument. Differential shifts by up to $\\sim$10\\% of the width of the diffraction image (i.e., several milliarcsec) may be caused by odd aberrations such as coma, even though the resolution remains essentially diffraction-limited. As a result, the measured centres of stellar images will depend on their spectral energy distributions, and a careful calibration of the effect, known as \\emph{chromaticity}, is mandatory in order to attain the astrometric accuracy goals. The gross spectral energy distribution of each observed target is therefore needed in the wavelength range of the astrometric CCDs; moreover, these data are needed with the same spatial resolution as in the astrometric field. As stated before, the BBP set of passbands was designed with this requirement in mind, as well as on astrophysical grounds.  \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig7.eps,width=1.0\\linewidth} \\end{center} \\caption{An example of the colour-dependent shifts (chromaticity) of stellar images that may be obtained in the astrometric field as the result of a coma-like optical aberration (see text for details). The shift relative to the geometric image centre is plotted versus a $V-I_c$-like colour index for a range of synthetic stellar spectra without extinction (filled circles) and for $A_V=2$~mag (open circles).} \\label{chrom1} \\end{figure}  \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig8.eps,width=1.0\\linewidth} \\end{center} \\caption{The chromatic shifts of a quasar image before (dash-dotted curve) and after (solid) correction for the effect calibrated by means of stellar spectra. A standard quasar spectrum was assumed to be observed at different red-shifts, and subject to the same aberrations as in Fig.~\\ref{chrom1}. (The dash-dotted curve has been displaced 4000~microarcsec downwards in the diagram to offset the overall shift caused by the coma).} \\label{chrom2} \\end{figure}  \\citet{lin_fil} showed that for the chromaticity calibration, near-rectangular filters are acceptable and that the choice of the separation wavelengths is more important than the edge widths. The author concluded that the use of four broad passbands covering the wavelength range of the astrometric $G$ passband should be enough to match the chromaticity constraints (r.m.s.\\ contribution to the parallaxes $<1$~microarcsec).  Following the design of the passbands in C1B, a more detailed evaluation of residual chromaticity effects was performed. A worst-case scenario, representing an extreme amount of coma, was considered where the wavefront aberration consisted of a third-degree Legendre polynomial in the normalized along-scan pupil coordinate, with r.m.s.\\ wavefront error (WFE) 45~nm. Polychromatic images in $G$ were generated for a library of synthetic stellar spectra \\citep{munari05}, including some reddened by interstellar extinction.  Image centres were computed through a modified form of Tukey's biweight formula \\citep{nrecipes}, with properties similar to the maximum-likelihood estimator to be used with the real {\\Gaia} data.  Figure~\\ref{chrom1} shows the stellar image shifts for a range of synthetic stellar spectra plotted versus one of the BBP colour indices. For the particular WFE assumed in this example, there is a general shift of the image centroid by $\\sim\\!4$~milliarcsec caused by the coma. However, it is only the variation of this shift with the spectral composition that is of concern here. The r.m.s.\\ variation of the shift versus colour index is 415~microarcsec. Using linear regression against the synthetic stellar C1B counts $\\phi_j$ ($j=1\\dots 5$, normalized to $\\sum_j\\phi_j=1$), the shifts could be reproduced with an r.m.s.\\ residual of 7~microarcsec. Further averaging between the $\\sim$800 astrometric CCD observations that are combined in a single parallax will reduce the astrometric effect of the chromatic residuals by a factor $0.02$--$0.2$ depending on the degree of correlation among the individual observations, thus leading to a chromatic contribution to the parallax errors of $0.14$--$1.4$~microarcsec. Since these numbers are based on a worst-case assumption for the WFE and a somewhat simplistic calibration model, it is reasonable to expect that a residual contribution of 1~microarcsec can be achieved, and in particular that the chosen C1B passbands provide sufficient information on the spectral energy distribution within the $G$ passband for this purpose. The residual effect is thus small compared with the statistical errors from photon noise and other sources even for bright stars.  The chromaticity correction based on an empirical calibration against stellar BBP fluxes will be less accurate for objects with strongly deviating SEDs. Prime examples of this are the quasars, which may exhibit strong emission lines at almost any wavelength depending on red-shift. Quasars are astrometrically important for establishing a non-rotating extragalactic reference system for proper motions; thus chromaticity correction must work also for these objects. This was investigated by applying the correction derived as described above from stellar spectra to synthetic quasar images, calculated from a mean quasar spectrum observed at red-shifts in the range $z=0$ to 6. Figure \\ref{chrom2} shows the uncorrected image shift as function of $z$ (dash-dotted curve) together with the residual shift after correction (solid curve). The r.m.s.\\ image shift decreases from 170~microarcsec before to 29~microarcsec after correction. The residual curve shows artifacts that are clearly attributable to the limited sampling of the spectral range. For example, the negative slopes for $z=2.3$--2.8 and $z=3.1$--3.9 correspond to the red-shifted Lyman-$\\alpha$ emission line moving through the C1B431 and C1B556 passbands, respectively.  With similar assumptions as for the stars concerning the statistical averaging of the effect when propagating to the astrometric parameters, the residual effect for the quasars will be a few microarcsec. This is acceptable since these are mostly faint objects with much larger photon-statistical errors.  Note that the data in Figs.~\\ref{chrom1}--\\ref{chrom2} only represent an example of the image shifts that may occur. The actual behaviour depends strongly on the shape and size of WFE, which vary considerably across the field of view, and on the detailed centroiding algorithm.  \\subsection{The C1M medium passbands} \\label{sec:section_MBP}  \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig9a.eps,width=1.0\\linewidth} \\epsfig{file=MF1354_fig9b.eps,width=1.0\\linewidth} \\end{center} \\caption  {Same as Fig.~\\ref{fig:SED_BBP} for C1M. Top and bottom figures show the `blue' and `red' passbands, respectively.} \\label{fig:SED_MBP} \\end{figure}  \\begin{table*} \\caption{{\\small Specifications of the filters in C1M implemented in Spectro.} } \\label{tab:filters6} \\begin{center} \\begin{tabular}{lccccccc} \\hline Band            &C1M326&C1M379 &C1M395& C1M410& C1M467& C1M506& C1M515 \\\\ \\hline $\\lambda_\\mathrm{blue}$ (nm) & 285 &  367  &  390 &   400 &   458 &   488 &   506 \\\\ $\\lambda_\\mathrm{red}$ (nm)  & 367 &  391  &  400 &   420 &   478 &   524 &   524 \\\\ $\\lambda_\\mathrm{o}$ (nm)    & 326 &  379  &  395 &   410 &   468 &   506 &   515 \\\\ $\\Delta \\lambda$ (nm)&  82 &   24  &   10 &    20 &    20 &    36 &    18 \\\\ $\\delta \\lambda$ (nm)&   5 &    5  &    5 &     5 &     5 &     5 &     5 \\\\ $\\epsilon$ (nm)     & 2,2 &  2,2  &  2,1 &   1,2 &   2,2 &   2,2 &   2,2 \\\\ $T_\\mathrm{max}$ (\\%)        &  90 &   90  &   90 &    90 &    90 &    90 &    90 \\\\ Type of CCD      &blue & blue  & blue &  blue &  blue &  blue &  blue \\\\ $n_\\mathrm{strips}$ &   2 &    2  &    1 &     1 &     1 &     1 &     1 \\\\ \\hline Band             &C1M549& C1M656 &C1M716& C1M747& C1M825& C1M861& C1M965 \\\\ \\hline $\\lambda_\\mathrm{blue}$ (nm) &   538&   652.8&   703&   731 &   808 &   845 &   930 \\\\ $\\lambda_\\mathrm{red}$ (nm)  &   560&   659.8&   729&   763 &   842 &   877 &  1000 \\\\ $\\lambda_\\mathrm{c}$ (nm)    &   549&   656.3&   717&   747 &   825 &   861 &   965 \\\\ $\\Delta \\lambda$ (nm)&    22&     7  &    26&    32 &    34 &    32 &    70 \\\\ $\\delta \\lambda$ (nm)&     5&     2  &     5&     5 &     5 &     5 &     5 \\\\ $\\epsilon$  (nm)   &   2,2&   1,1  &   2,2&   2,2 &   2,2 &   2,2 &   2,2 \\\\ $T_\\mathrm{max}$ (\\%)        &    90&    90  &    90&    90 &    90 &    90 &    90 \\\\ Type of CCD     &  blue&   red  &   red&   red &   red &   red &   red \\\\ $n_\\mathrm{strips}$ &     1&     1  &     1&     1 &     1 &     1 &     1 \\\\ \\hline \\multicolumn{8}{l}{\\footnotesize{$\\lambda_\\mathrm{blue}$, $\\lambda_\\mathrm{red}$: wavelengths at half-maximum transmission}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$\\lambda_\\mathrm{c}$: central wavelength$= 0.5(\\lambda_\\mathrm{blue} + \\lambda_\\mathrm{red})$}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$\\Delta \\lambda$: FWHM}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$\\delta \\lambda$: edge width (blue, red) between 10 and 90\\% of $T_\\mathrm{max}$}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$\\epsilon$: manufacturing tolerance intervals centred on $\\lambda_\\mathrm{blue}$ and $\\lambda_\\mathrm{red}$}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$T_\\mathrm{max}$: maximum transmission of filter}}\\\\ \\multicolumn{8}{l}{\\footnotesize{$n_\\mathrm{strips}$: Number of CCD strips carrying the filter}}\\ \\end{tabular} \\end{center} \\end{table*}  \\begin{figure} \\begin{center} \\leavevmode \\epsfig{file=MF1354_fig10.eps,width=1.0\\linewidth} \\end{center} \\caption{Estimation of the end-of-mission precisions as a function of $V$ for several photometric indices constructed from the passbands in C1M, and computed according to equation (\\ref{eq:eqerror}) with $\\sigma_{\\rm cal}=0$~mag. In this paper, the contribution of calibration error to the end-of-mission precision is estimated to be of 3~millimag in every passband.  This is indicated by the horizontal line (which here includes a factor of $\\sqrt2$).} \\label{fig:errors_MBP} \\end{figure}  The C1M component of the {\\Gaia} PS consists of 14 passbands and evolved from the convergence of the proposals by \\citet{grenon}, \\citet{knude}, \\citet{vladas1X}, \\citet{vytas2004} and \\citet{jordimbp}. The guidelines by \\citet{taut2002} for $\\alpha$-elements abundance determination were taken into account.  The basic response curve of the filters versus wavelength is a symmetric quasi-trapezoidal shape. Their parameters are listed in Table~\\ref{tab:filters6} and the response of the corresponding passbands is shown in Fig.~\\ref{fig:SED_MBP}. Six strips with red-enhanced CCDs are available for MBP and six red passbands have been designed, implemented as one filter for each strip, with the only constraint that MBP has to measure the flux entering the RVS instrument (see Section~\\ref{sec:principles_PS}). For the blue passbands eight filters are implemented on the ten strips with blue-enhanced CCDs. Two strips have been allocated to each of the two UV passbands to increase the signal-to-noise ratio (SNR) of the measurements.  The primary purpose of the medium passbands is the classification and astrophysical parametrization of the observed objects (stars, QSOs, galaxies, solar-system bodies, etc.). In the case of the stars the goal is to determine effective temperature \\teff, gravity $\\log g$ or luminosity $M_V$, chemical composition [M/H], \\afeh\\ and C/O abundances, peculiarity type, the presence of emission, etc., in the presence of varying and unknown interstellar extinction. Taxonomy classification for solar-system objects and photometric red-shift determination for QSOs are also aimed for.  \\subsubsection {Astrophysical diagnostics}  The filter at the shortest wavelength is C1M326 with wavelengths at half-maximum transmission of 285 and 367~nm (thus it is of broad passband type although implemented in Spectro). Below 280~nm, strong absorption lines, metallicity dependent, are present already in A type stars. The interstellar extinction increases rapidly below 280~nm and reaches a maximum at 218~nm. In space the UV passband can be extended down to 280~nm which improves the determination of {\\feh} for F-G-K stars because of the presence of many atomic lines, ionized or of high excitation. For late G dwarfs, the line blocking in the extended UV passband is about 2.7 times larger than in the violet 376--430~nm domain. The red side of C1M326 is set at the Balmer jump. The colour indices C1M326--C1B431 or C1M326--C1M410 give the height of the Balmer jump, which is a function of {\\teff} and {\\logg} in B-A-F stars. In F-G-K stars these colour indices measure metallic line-blocking in the ultraviolet which can be calibrated in terms of {\\feh}.  The C1M379 passband is placed on the wavelength range where the lines corresponding to the higher energy levels of the Balmer series crowd together in early-type stars. The integrated absorption in these lines is very sensitive to {\\logg} (or $M_V$). For late-type stars the position of this passband coincides with the maximum blocking of the spectrum by metallic lines. Hence the colour index C1M379--C1M467 is a sensitive indicator of metallicity. Analogs in other photometric systems are the $P$ magnitude in the Vilnius system and $L$ in the Walraven system.  The C1M395 passband is introduced mainly to measure the Ca\\,II~H line. The index C1M395--C1M410 shows a strong correlation with W(CaT*), the equivalent width of the Calcium triplet measured by the RVS instrument, corrected for the influence of the Paschen lines. The C1M395--C1M410 vs.\\ W(CaT*) may be used as a {\\logg} estimator (\\citealt{kaltcheva}, Knude \\& Carrasco, priv. comm.).  The C1M395--C1M410 vs.\\ W(CaT*) plane is particularly useful since the effect of reddening on the colour index is only minor due to the small separation of the two passbands. Additional uses of the C1M395 passband are in assisting the {\\afeh} determination and in the identification of very metal-poor stars.  The violet C1M410 passband measures the spectrum intensity redward of the Balmer jump. In combination with C1M326, it gives the height of the jump. For K--M stars it is the shortest passband which when combined with longer passbands can provide temperatures and luminosities of solar metallicity stars in the presence of interstellar reddening (i.e.\\ when the stars are too faint in the ultraviolet).  Its analogs are: $v$ in the Str\\\"omgren system, $B_1$ in the Geneva system and $X$ in the Vilnius system.  The blue C1M467 and green C1M549 passbands measure domains where the absorption by atomic and molecular lines is minimal.  The flux in these domains correspond to a pseudo-continuum. The colour indices C1M467--C1M549, C1M467--C1M747 and C1M549--C1M747 may be used as indicators of the temperature for stars of all spectral types.  The analogs of C1M467 are: $b$ in the Str\\\"omgren system, $B_2$ in the Geneva system, and $Y$ in the Vilnius system.  The analogs of C1M549 are:  $y$ in the Str\\\"omgren system and $V$ in the Vilnius system.  The green C1M515 passband is placed on a broad spectral depression seen in the spectra of G and K type stars and formed by crowding of numerous metallic lines. Among them, the strongest features are the Mg\\,I triplet and the MgH band.  The depth of this depression, the intensity of which reaches a maximum around K7~V, is very sensitive to gravity, being deeper in dwarfs than in giants. The same passband is also useful for the identification of Ap stars of the Sr-Cr-Eu type.  The same passband ($Z$) is used in the Vilnius photometric system.  The C1M506 is much broader and includes the C1M515 passband region. The combination of both provides an index which is almost reddening free and its combination with the contiguous pseudo-continuum passbands (C1M467 and C1M549) provides an index sensitive to Mg abundances and gravity. If the luminosity is known from parallax, Mg abundances can be determined. The Ca\\,II and Mg\\,I spectral features show inverse behaviour when {\\feh} and {\\afeh} change \\citep{taut2002}, and hence, indices using C1M395 and C1M515 allow the disentangling of Fe and $\\alpha$-process element abundances.  The narrow passband C1M656 is placed on the H$\\alpha$ line.  As mentioned in Section~\\ref{sec:section_BBP}, the H$\\alpha$$=$C1B655--C1M656 index is a measure of the intensity of the H$\\alpha$ line, yielding luminosities for stars earlier than A0 and temperatures for stars later than A3. The index is most useful for identification of emission-line stars (Be, Oe, Of, T Tau, Herbig Ae/Be, etc.).  C1M716 coincides with one of the deepest TiO absorption bands with a head at 713 nm \\citep{wahlgren}, while C1M747 measures a portion of the spectrum where the absorption by TiO bands is minimum. So, the index C1M716--C1M747 is a strong indicator of the presence and intensity of TiO, which depends on temperature and TiO abundance for late K and M type stars. For earlier type stars, both passbands provide measurements of the pseudo-continuum. Earlier PS proposals for MBP considered the inclusion of a filter centred on the TiO absorption band at 781~nm. FoM computations and analysis of the \\skpost values showed that the AP determination improves by more than 10\\%, for [Ti/H] and {\\teff}, if the passband is centred on 716~nm instead of on 781~nm. A similar but narrower passband has been used in the Wing eight-colour far-red system.  The passband C1M825 is designed to measure either the continuum blueward of the Paschen jump (hence its limitation at 842~nm for the red side mid-transmission wavelength) or the strong CN band for R- and N-type stars.  For M stars, C1M825 measures a spectral domain with weak absorption by TiO. The distinction between M and C stars is realized with all red passbands. At a given temperature, the fluxes are similar in the C1M747 and C1M861 for O-rich stars (the M sequence) and for C-rich stars (the C sequence), but very different in the C1M825 and C1M965 passbands, namely because of strong CN bands developing redward of 787 nm. The separation between M and C stars is possible even if they are heavily reddened.  Similarly, C1M965 measures the continuum redward of Paschen jump (and in combination with C1M825 yields the height of the jump) or strong absorption bands for R- and N-type stars (see Fig.~\\ref{fig:SED_MBP} bottom). Having a passband at these very red wavelengths at the egde of the CCD QE curve improves the interstellar extinction determination, which was proven through the FoM computations.  The C1M861 passband, in between C1M825 and C1M965, is constrained by the wavelength range of the RVS instrument (i.e.\\ 848-874~nm) and hence includes the Ca IR triplet. The measurement of the flux of the star in this passband will help the RVS data reduction. The index C1M861--C1M965 measures the gravity-sensitive absorption of the high member lines of the Paschen series.  Finally, the indices C1M825--C1M861 and C1M861-C1M965 are a sensitive criterion for the separation of M-, R- and N-type stars (see Fig.~\\ref{fig:SED_MBP} bottom). ",
        "conclusions": ""
    },
    "0512/astro-ph0512512_arXiv.txt": {
        "abstract": "{ Stable lumps of quark matter may be present in cosmic rays at a flux level, which can be detected by high precision cosmic ray experiments sensitive to anomalous \\textquotedblleft nuclei\\textquotedblright\\ with high mass-to-charge ratio. The properties of these lumps, called strangelets, are described, and so is the production and propagation of strangelets in cosmic rays. Two experiments underway which are sensitive to a strangelet flux in the predicted range are briefly described. Finally it is summarized how strangelets circumvent the acceleration problem encountered by conventional candidates for ultra-high energy cosmic rays and move the Greisen-Zatsepin-Kuzmin cutoff to energies well above the observed maximum energies. }  \\FullConference{29th Johns Hopkins Workshop on current problems in particle theory: Strong Matter in the Heavens\\\\ August 1--3, 2005\\\\ Budapest, Hungary}  \\begin{document} ",
        "introduction": "Iron and nickel nuclei are normally assumed to be the most stable form of hadronic matter at zero external pressure. In principle, this should be testable from the basic theory of strong interactions, Quantum Chromo Dynamics (QCD), but in practice this is impossible in the foreseeable future. QCD is not suited for finite density calculations and calculations involving many degrees of freedom, so much of our theoretical knowledge about dense matter is based on phenomenological model calculations that try to incorporate and parametrize some of the main features of the strong interactions, such as confinement and asymptotic freedom. In some of these studies, notably studies based on the MIT bag model, it has been shown, that there is a significant range of parameters for which a three-flavor quark phase with roughly equal numbers of up, down and strange quarks (called strange quark matter in bulk, and strangelets in small lumps) has lower energy than a nucleus with the same baryon number. Thus, strangelets rather than nuclei may be the ground state of hadronic matter \\cite{Bodmer:1971we, Chin:1979yb, Witten:1984rs, Farhi:1984qu, Madsen:1998uh, Weber:2004kj}.  A range of questions immediately occur if strangelets are more stable than nuclei:  \\begin{enumerate} \\item Why are we here---why don't our nuclei decay?  \\item Can strangelets be created in the laboratory?  \\item Can strange quark matter be found in space? \\end{enumerate}  The answers to all three questions rely on the properties of strangelets summarized in the following. The answers to the questions (some of which will be further explored below) are briefly:  \\begin{enumerate} \\item Very low baryon number strangelets are likely to be unstable, even if strange matter in bulk is absolutely stable, and for intermediate baryon numbers the transformation of a nucleus requires an improbable high-order weak interaction to simultaneously transform of order $A$ up and down quarks into strange quarks, where $A$ is the baryon number. This explains the stability of our nuclei. Strangelets have a positive electric charge, and therefore repel nuclei from the surface, thus minimizing the risk of growth via absorption of nuclei \\cite{Witten:1984rs, Farhi:1984qu, Dar:1999ac, Busza:1999pr, Madsen:2000kb, Blaizot:2003bb}.  \\item In principle strangelets might be formed in ultrarelativistic heavy-ion collisions by coalescence or through a distillation mechanism leading to strangeness enhancement. However, the available baryon number is low, which makes it difficult to cross the low-$A$ stability cutoff, and furthermore the environment is hot, so the process has rightly been compared to the creation of ice cubes in a furnace. Experiments have been performed, but with negative results \\cite{Sandweiss:2004bu, Klingenberg:2001qs, :2005cu}.  \\item The most promising place to search for strange quark matter is in the cosmos. Originally strange quark matter \\textquotedblleft nuggets\\textquotedblright\\ surviving from the cosmological quark-hadron phase transition were suggested as a possible and even natural candidate for the cosmological dark matter \\cite{Witten:1984rs}, but later studies showed, that this was probably unlikely due to the high temperature environment that would lead to evaporation of the nuggets \\cite{Alcock:1985vc, Madsen:1986jg, Alcock:1988br, Olesen:1991zt, Olesen:1993iz, Bhattacharjee:1993ah}. Recent lattice results indicating, that the quark-hadron transition at low chemical potential (like the early Universe) is not a first order phase transition seems to rule out the possibility of cosmological strange quark matter. But strange quark matter, if stable, is almost unavoidable in dense stellar objects. Pulsars and compact x-ray sources will contain strange stars containing bulk strange quark matter \\cite{Witten:1984rs,Haensel:1986qb, Alcock:1986hz}, and a number of observational signatures have been suggested to distinguish these from compact stars made of ordinary hadronic matter. This is an interesting story in its own right (see \\cite{Madsen:1998uh, Weber:2004kj} for reviews), but the focus of the present presentation is the possibility of directly observing debris from collisions of strange stars in the form of strangelets in cosmic rays. \\end{enumerate}  In the following I will summarize the main properties of strangelets, describe the propagation of strangelets in cosmic rays, estimate the flux to be expected in cosmic ray detectors, and discuss the possible detection in a couple of future experiments. Finally I will argue that strangelets may also be relevant for the puzzle of ultra-high energy cosmic rays. ",
        "conclusions": "The total strangelet flux reaching the Moon or a detector in Earth orbit is in a regime that could be within experimental reach, and therefore provide a crucial test of the hypothesis of absolutely stable strange quark matter. Experiments are underway which are sensitive to the high mass-to-charge signature expected for such events within the flux-range predicted. Strangelets may also provide an interesting explanation of ultra-high energy cosmic ray events."
    },
    "0512/astro-ph0512348_arXiv.txt": {
        "abstract": "We present metallicities, [Fe/H], and elemental abundance ratios, [X/Fe], for a sample of 24 Cepheids in the outer Galactic disk based on high-resolution echelle spectra. The sample have Galactocentric distances covering 12 $\\le$ \\rgc~(kpc) $\\le$ 17.2 making them the most distant Galactic Cepheids upon which detailed abundance analyses have been performed. We find sub-solar ratios of [Fe/H] and overabundances of [$\\alpha$/Fe], [La/Fe], and [Eu/Fe] in the program stars. All abundance ratios exhibit a dispersion that exceeds the measurement uncertainties. As seen in our previous studies of old open clusters and field giants, enhanced ratios of [$\\alpha$/Fe] and [Eu/Fe] reveal that recent star formation has taken place in the outer disk with Type II supernovae preferentially contributing ejecta to the ISM and with Type Ia supernovae playing only a minor role. The enhancements for La suggest that AGB stars have contributed to the chemical evolution of the outer Galactic disk.  Some of the young Cepheids are more metal-poor than the older open clusters and field stars at comparable Galactocentric distances. This demonstrates that the outer disk is not the end result of the isolated evolution of an ensemble of gas and stars. We showed previously that the older open clusters and field stars reached a basement metallicity at about 10-11 kpc. The younger Cepheids reach the same metallicity but at larger Galactocentric distances, roughly 14 kpc. This suggests that the Galactic disk has been growing with time as predicted from numerical simulations.  The outer disk Cepheids appear to exhibit a bimodal distribution for [Fe/H] and [$\\alpha$/Fe]. Most of the Cepheids continue the trends with Galactocentric distance exhibited by Andrievsky's larger Cepheid sample and we refer to these stars as the ``Galactic Cepheids''. A minority of the Cepheids show considerably lower [Fe/H] and higher [$\\alpha$/Fe] and we refer to these stars as the ``Merger Cepheids''. One signature of a merger event would be composition differences between the ``Galactic'' and ``Merger'' Cepheids. The Cepheids satisfy this requirement and we speculate that the distinct compositions suggest that the ``Merger Cepheids'' may have formed under the influence of significant merger or accretion events. The short lifetimes of the Cepheids reveal that the merger event may be on-going with the Monoceros ring and Canis Major galaxy being possible merger candidates. ",
        "introduction": "The Galactic disk accounts for the vast majority of stars in the Milky Way. Abundance analyses of nearby field stars \\citep{bdp93,bdp03}, open clusters \\citep{friel95}, planetary nebulae \\citep{henry04}, and H II regions \\citep{shaver83} have provided insight into the mean metallicity and metallicity gradient of younger and older stars and stellar remnants. Radial abundance gradients, as measured in disk stars, provide crucial constraints for models of the formation and evolution of our Galaxy \\citep{hou00,chiappini01}. However, to explore the origin and continuing evolution of the Galactic disk, we need much more information than mean metallicities alone can provide. We need detailed elemental abundance ratios, which contain vital information on the relative contributions of Type II supernovae, Type Ia supernovae, and asymptotic giant branch (AGB) stars.  In order to address this situation, we commenced an observing program to measure metallicities and detailed chemical compositions for stars in the outer Galactic disk. An analysis of the radial velocities and chemical abundance patterns of four old open clusters with Galactocentric distances between 12 and 23 kpc was presented in Paper I \\citep{open}. In Paper II \\citep{warp}, we continued our attack upon the outer disk by initiating a successful selection process for identifying distant field giants then conducting an abundance analysis of three such stars. From these analyses, two principal findings may be noted: (1) at large Galactocentric distances, \\rgc~$> 10$ kpc, the expected metallicity gradient vanishes and the stars exhibit a constant value of [Fe/H] $\\approx$ $-$0.5; (2) field and cluster stars in the outer Galactic disk show enhancements for the ``$\\alpha$'' elements, [$\\alpha$/Fe] $\\approx$ 0.2. Our interpretation was that these abundance patterns reflected the episodic growth of the disk via accretion or merger events. These events triggered rapid star formation with Type II supernovae preferentially contributing to the chemical enrichment.  The open clusters are old, with ages between 2 - 6 Gyr. The ages of the field giants are unknown, but presumably these stars are quite old too. The absence of the radial abundance gradient inferred from our sample of distant field and clusters stars may therefore not reflect the current situation in the outer disk. The time variation of the Galactic radial abundance gradient offers a more comprehensive test of Galactic evolution models than a single epoch (or time integrated) abundance gradient. The possibility of an on-going merger event in the outer disk (e.g. \\citealt{newberg02,ibata03,yanny03,martin04}) reinforces the need to analyze a sample of young stars at large distances and to measure detailed abundance ratios [X/Fe] in addition to the metallicity, [Fe/H].  Cepheids are high-mass stars whose short lifetimes ensure that their atmospheres reflect the present-day composition of the ISM. Abundance analyses of Cepheids appear feasible \\citep{fry97,andrievsky02a,andrievsky02b,andrievsky02c,andrievsky04,luck03} and the Cepheid period-luminosity relation allows for accurate distance determinations. Due to their luminosity, high-resolution spectroscopic observations of Cepheids located at large distances can be conducted with modest-sized telescopes. In this paper, the third of this series, we present metallicities, [Fe/H], and elemental abundance ratios, [X/Fe], for a sample of two dozen Cepheids with Galactocentric distances 12 $\\le$ \\rgc~$\\le$ 17.2 kpc. These Cepheids allow us to study the Galactic radial abundance gradient as a function of time. Detailed abundance ratios [X/Fe] offer an insight into the events currently taking place in the outer disk and such abundance ratios may reveal the signatures of recent merger events. ",
        "conclusions": "In the previous two papers in this series, we presented the chemical compositions for old open cluster giants and field giants in the outer Galactic disk. In this paper we conduct an abundance analysis of 24 young Cepheids located at large Galactocentric distances. The program Cepheids therefore allow us to study the time evolution of the Galactic radial abundance distribution as well as the current radial abundance distribution.  The short lifetimes of Cepheids ensures that their compositions reflect the recent state of the ISM. In general, the abundances measured in our program Cepheids continue the trends with Galactocentric distance seen in the large sample of Cepheids analyzed by Andrievsky. We also find enhancements for [La/Fe] and [Eu/Fe] in the outer disk. For all elements, we find a scatter that exceeds the measurement uncertainties. The enhancements in [Eu/Fe] and [$\\alpha$/Fe] suggest that Type II supernovae have played a greater role in the chemical evolution than have Type Ia supernovae in the outer disk, compared to the inner disk. The short lifetimes of the Cepheids demonstrate that recent star formation has taken place in the outer disk. The high ratios of [La/Fe] suggest that AGB stars have also played a role in the evolution of the outer disk. We find that the ratio [La/Eu] is centered at the solar value with no star showing a scaled-solar pure $s$-process or $r$-process value.  The sample of Cepheids in the outer disk, although numbering only 24 stars, has provided some tantalizing clues to the evolution of our Galaxy. As expected, the outer disk Cepheids are on average more metal-poor than Cepheids with smaller Galactocentric distances. However, the outer disk Cepheids show higher abundances of the $\\alpha$ elements, despite having very similar ages to Cepheids in the inner disk. The simplest conclusion to be drawn is that recent star formation in the outer disk has provided a recent enhanced production of Type II supernovae relative to Type Ia supernovae. Either the inner disk is undergoing a slower pace of chemical evolution, or there is enhanced star formation underway in the outer disk. The question would then be what is causing this phenomenon in such a low density environment?  There are two hints in the data that point toward accretion as being the underlying cause of the enhanced [$\\alpha$/Fe] ratios. First, the most distant Cepheids appear to have bimodal distributions of [Fe/H] and [$\\alpha$/Fe]. The more metal-rich Cepheids of the outer disk appear to share a similar [$\\alpha$/Fe] vs.\\ [Fe/H] trend as the inner disk Cepheids. But the outlying more metal-poor outer disk Cepheids show even higher [$\\alpha$/Fe] values, suggesting a separate history. Among these stars, it is intriguing that the most recently formed such star (with the longest period) has a significantly lower [$\\alpha$/Fe] values than the four shorter period (older) Cepheids, hinting that the sources of nucleosynthesis may have been changed on a timescale of only a few tens of millions of years. Alternately, if accretion is responsible for the difference, it may be that the four older stars formed from gas that was richer in material from a merging galaxy while the more recently formed star emerged from gas more throughly mixed with Galactic gas.  A comparison between the young Cepheids and older field stars and open cluster in the outer disk also provides clues about the evolution of the outer Galactic disk. Disregarding the apparent bimodality of the outer disk Cepheids, one must confront the observation that the younger Cepheids are more metal-poor than the older clusters and field stars. Accepting the bimodality, and concentrating on only that Cepheids that appear to continue the trend of [$\\alpha$/Fe] vs.\\ [Fe/H] seen in inner disk Cepheids, one finds a behavior similar to the older clusters and field stars. [Fe/H] appears to reach a basement value of about $-0.5$ while [$\\alpha$/Fe] reaches a ceiling of +0.15. However, the Cepheids reach such values only for $R_{\\rm GC} > 14$ kpc, while the older clusters and stars do so at 10 to 11 kpc. This suggests that the outer {\\em stellar} disk has grown in radius by several kpc in the past several billion years. This suggests a past history of accretion, in addition to that apparently recently underway.  Accretion events have been a common theme in other recent studies of the Galaxy's outer disk. The Sloan Digital Sky Survey has detected the ``Monoceros Ring\" \\citep{newberg02}. \\citet{yanny03} have found clear dynamical evidence for a structure at $\\ell$ = 198, $b = -27$ that may be part of a more extensive merger remnant. The 2MASS database and its ability to distinguish M dwarfs from M giants has led to the identification of the candidate ``Canis Major galaxy\", whose center lies near $\\ell\\ \\approx\\ 244$, $b \\approx\\ -8$ \\citep{martin04,martin05,bellazzini04,delgado05}. Radial velocities of the photometrically identified streams has confirmed the existence of unique streams \\citep{martin05,penarrubia05,conn05}. Are any of these related to the unusual star formation history we are suggesting for the outer disk? Certainly all the evidence points to on-going accretion episodes, but how may we distinguish individual events?  The referee has drawn our attention to the fact that the metal-poor but $\\alpha$-rich Cepheids are not distributed as widely as the other Cepheids in Table~1. These seven stars, CI~Per, EW~Aur, FO~Cas, GP~Per, IO~Cas, NY~Cas, and OT~Per, all lie at low Galactic latitudes (and six between $b = -2$ and $-4$) and between $\\ell =$ 119 and 166 degrees. Could these stars' locations be a clue about a possible merger origin? If there is a relation to an already-suggested merger event, its location is more consistent with Canis Major than with the Monoceros Ring. But even so, we are reluctant to ascribe much significance to the spatial locations of these ``Merger Cepheids\". They are, after all, very spread out in longitude. Further, it is not clear how uniform are the searches for distant Cepheids. More extensive work in even one hemisphere than another could create such an apparent grouping (and all these Cepheids are northern ones). A more thorough search for Cepheids and, especially, the more abundant but comparably young OB stars may be needed to resolve this question.  We argue that detailed chemical composition studies such as we have undertaken here and in Papers I and II are critical to understanding the complex history of our Galaxy. Distinct chemical evolution patterns are as important, if not more so, in probing the relationship between candidate streams and merging galaxies to one another and to the stars in the outer Galactic disk. Unfortunately, little work has been done as yet. So far, only 3 stars in Canis Major have been analyzed \\citep{sbordone05}, and the data employed had S/N=40 per pixel at 5800\\AA~which is less than ideal (in our opinion). But the results are intriguing. The abundance ratios found by \\citet{sbordone05} do not appear to match those seen in our samples of outer disk stars. Specifically, the Canis Major candidates have low [$\\alpha$/Fe] while our outer disk stars all show enhancements in [$\\alpha$/Fe]. Low abundances of [$\\alpha$/Fe] are a signature of the current dwarf spheroidals orbiting our Galaxy \\citep{venn04}. Enhancements in [$\\alpha$/Fe] in the ``Merger Cepheids'' may therefore suggest that these Cepheids formed as a result of star formation triggered by the merger rather than forming within the dwarf galaxy. Unfortunately, detailed stellar abundance ratios are unable to offer a clearer picture of the exact mechanism. While we speculate that the star formation was triggered by merging gas, we cannot say whether that gas was pristine or pre-enriched. Detailed abundance ratios in a large sample of candidate members of the Monoceros ring and/or Canis Major Galaxy are required. The measured abundance ratios would confirm if the outer disk has been growing via the merger of dwarf galaxies as well as providing the chemical history of these small galaxies. Such a result would have profound implications not only for our understanding of the evolution of our Galaxy but also for $\\Lambda$CDM cosmology."
    },
    "0512/astro-ph0512381_arXiv.txt": {
        "abstract": "We  investigate statistically  the dynamical  consequences  of cosmological fluxes of matter  and related moments on progenitors  of today's dark matter haloes. These  haloes are described  as open collisionless systems  which do not  undergo strong interactions  anymore. Their  dynamics is  described via canonical perturbation theory which accounts for two types of perturbations: the  tidal field  corresponding  to  fly-bys and  accretion  of dark  matter through the halo's outer boundary.  The non-linear evolution of both the  entering flux and the particles of the halo  is  followed  perturbatively.   The  dynamical  equations  are  solved linearly, order by order, projecting on a biorthogonal basis to consistently satisfy  the  field equation. Since our perturbative solution of the Boltzmann Poisson is explicit, we obtain, as a result, expressions for the N-point correlation  function of the halo's response to the  perturbative environment.   It allows  statistical predictions  for the ensemble distribution of  the inner dynamical features of  haloes. We demonstrate the feasibility of the implementation via a simple example in the appendix.  We argue that the  fluid description  accounts for the  dynamical drag and  the tidal stripping  of  incoming structures.   We  discuss  the  realm of  non-linear problems which  could be addressed statistically  by such a  theory, such as differential  dynamical friction, tidal  stripping and  the self  gravity of objects within the virial sphere.   The secular evolution of open galactic haloes is investigated: we derive the kinetic  equation which governs  the quasi-linear  evolution of  dark matter profile induced by infall  and its corresponding gravitational correlations. This yields a Fokker  Planck-like equation for the angle-averaged underlying distribution function. This equation  has an explicit source term accounting for  the net  infall  through the  virial  sphere. Under  the assumption  of ergodicity  we then  relate the  corresponding source,  drift  and diffusion coefficients for the ensemble-average  distribution to the underlying cosmic two-point statistics of the infall and discuss possible applications.  The internal dynamics of  sub-structures within galactic haloes (distortion, clumps   as  traced  by   Xray  emissivity,   weak  lensing,   dark  matter annihilation, tidal streams  ..), and the implication for  the disk (spiral structure, warp \\etc) are then  briefly discussed.  We show how this theory could be  used to (i)  observationally constrain the statistical  nature of the  infall (ii)  predict  the observed  distribution  and correlations  of substructures in upcoming surveys, (iii)  predict the past evolution of the observed  distribution of  clumps,  and finally  (iv)  weight the  relative importance of the intrinsic (via the unperturbed distribution function) and external (tidal and/ or infall) influence of the environment in determining the  fate  of galaxies.   We  stress that  our  theory  describes the  {\\sl perturbed } distribution function  (mean profile removed) directly in phase space. ",
        "introduction": "It now  appears clearly that  the dynamical (azimuthal  instabilities, warps, accretion),  physical (heating, cooling)  {and} chemical  (metal-poor  cold gas fluxes) evolution of galaxies are   processes which are {partly driven} by the boundary conditions imposed by their cosmological environment. It is therefore of prime importance {to formulate the effects of such an interaction in a unified framework}.   Modern digital all-sky surveys, such as the  SDSS, 2MASS or the 2dF provide for the first time the opportunity to build  statistically relevant  constraints on the dynamical states  of galaxies {which can} be  used as observational input. Other projects, like Gaia or Planck, will provide small-scale information on our Galaxy  and its  environment and will  soon allow detailed confrontation  of the predictions of {models} with the  observations. We ought to be  able to draw conclusions on the internal dynamics of the halo and its inner components and constrain their statistical properties.   Unfortunately, it is difficult to study the response of haloes to moderate amplitude perturbations. Current N-body techniques suffer from resolution limitations (due  to particle number and drift in orbit integration, see e.g. \\citet{Power}, \\citet{Binney04}, for a discussion of such effects)   that {hide to some extent} linear collective effects which  dominate the response of the halo (\\cite{weinberg2}, \\cite{murali})\\footnote{it has been argued that shadowing (\\cite{Tremaine2}) will in practice allow for another orbit to correct for the drift, but this is of no help to resonant processes because it requires that the {\\sl same} orbit does not diffuse for a few libration periods.}. Simulations on galactic scales are also often carried without any attempt to {represent} the cosmological variety arising from the possible boundary conditions (the so-called cosmic variance problem). This {is because the} dynamical range required to describe both the environment and the inner structure is considerable, and can only be achieved for a limited number of simulations (e.g. \\citet{Knebe}, \\citet{Gill}, \\citet{Diemand}). {By}  contrast, the method presented below {circumvents} this difficulty while relying on an {\\sl explicit} treatment of the inner dynamics of the halo, in the perturbative regime. Specifically, our purpose  is to develop a tool to study  the dynamics of an open stellar system and apply it to  the dynamic of a halo which is embedded into  its cosmological  environment. One can think of this project as an attempt to produce a semi-analytic explicit re-simulation tool, in the spirit of what is done in N-body simulations with zoomed-in initial conditions.  {The} concept of an initial power spectrum describing the statistical properties of the gravitational perturbations has {proved very useful in cosmological studies} (e.g. \\citet{Peebles}, \\citet{BCGS}). The underlying paradigm{,} that gravity {drives} cosmic evolution, is likely to {be} a good description at the megaparsec scale. We  show  below  that a  similar  approach  {to}  galactic haloes  is  still {acceptable,} and marginally within  the reach of our modeling capabilities. {The} description  of the  boundary is significantly  more complex,  but the inner dynamics of  hot components is better behaved.   Here, we {describe} a stable system  which undergoes small  interactions, rather than  an unstable system in comoving coordinates undergoing catastrophic multi-scale collapse.     The purpose of this investigation is  to derive analytically the dynamical response of a galactic halo,  induced  by  its (relatively weak)  interaction   with  its  near environment.   Interaction  should be  understood  in  a  general sense  and involve  {tidal} potential  interactions (like  that  corresponding  to a  satellite orbiting around the galaxy), or an infall where an external quantity (virialized or not) is advected  into the galactic halo.  With  a suitable formalism, we derive the  propagation of an  external perturbation from the near  galactic environment  down to the  scale of the galactic disk through  the dark matter halo. {We essentially} solve the coupled collisionless Boltzmann-Poisson equations as a Dirichlet initial value problem to determine the response of the halo to infall and tidal field. {The} basis over which the response is projected {can} be  customized to, say,  the universal  profile of  dark matter  haloes, {which makes it possible to}  consistently and efficiently {solve} the coupled dynamical and  field equations {,} so long as the entering fluxes of dark matter amounts to  a small perturbation in mass compared to the underlying equilibrium.        {In a} pair of companion papers, \\citet{aubert1,aubert2} described the statistical properties of the infalling distribution of dark matter at the virial radius, $R_{200}$ as a function of cosmic time between redshift $z=1$ and today. These papers focused on a description of the one- and two-point statistics of the infall towards well formed $L_\\star$ dark matter haloes. All measurements were carried for 15~000  haloes undergoing minor mergers. The  two-point  correlations were measured both angularly  and temporally for the flux  densities, and over the whole 5D phase space for the expansion coefficients of the source.  Together with the measurements presented there, we show in this paper that the formalism described below {will} allow astronomers to address globally and coherently dynamical issues on galactic scales. Most importantly it will allow them to tackle problems in a {\\sl statistically representative } manner. This investigation has a broad field of  possible applications.   Galaxies are subject to boundary conditions  that reflect {motions} on  larger scales  and their dynamics may constrain the cosmology  through the rate of merging events for example,  or the  mass distribution  of satellites. Halo transmission and amplification  also fosters communication between spatially separated regions, (see \\eg  \\cite{murali}) and  continuously {excites the}  disk structure. {For example}, spirals  can  be  induced  by  encounters with  satellites  and/or  by  mass injection (e.g. \\citet{TT}, \\citet{Howard}), while warps results from torque interactions with the surrounding matter (\\citet{Lopez}, \\citet{JiangBinney}). Therefore  the proportion of spirals and  warps contains information on the  structure's formation and  environment. The statistical link between the inner properties of galactic haloes, and their cosmic boundary can be reversed  to attempt and constrain the nature of the infall while investigating the one and two point statistics of the induced perturbations. This is best done by transposing down to galactic scales the classical  cosmic probes for the large-scale structures (lensing, SZ, \\etc) which have been used successfully to characterize the power-spectrum of fluctuations on larger scales.    The outline of this paper is the following: we describe in \\Sec{spherical} the linear response of a spherical halo which undergoes cosmological infall of dark matter,  and compute the induced correlations in the inner halo; \\Sec{nonlinear} presents the second-order perturbative response of the galactic halo to the infalling flux; (Appendix \\ref{s:apendperturb} gives the higher order corrections to the dynamics and addresses the issue of dynamical friction). \\Sec{quasi1} derives the Fokker-Planck equation that the cosmic mean halo profile {obeys} in such an open environment. \\Sec{applications} describes briefly possible astrophysical applications. In particular, it is discussed how {the}  statistical analysis of mean and {variance} properties of  galactic haloes and galaxies can be compared to the quantitative prediction of  the  concordant $\\Lambda$CDM cosmogony on those scales. We also show how to revert in time observed tidal features within our Galaxy, or in external galaxies. The last section draws conclusions and discusses prospects for future work. ",
        "conclusions": "In  the last  few  years,  with the  observational  convergence towards  the concordant    cosmological   model,  a significant  fraction  of  the interest  has shifted towards smaller  scales.  Indeed it now becomes  possible to project down to  these scales some  of the predictions  of the model.  This  in turn offers the prospect of transposing there what has certainly been a key asset of  modern cosmology,  both observationally  and  theoretically: statistics. This is a requirement both from the point of view of the (often understated) variety of objects  falling onto an $L_{\\star}$ galaxy,  but also because of the  sheer  size of  the  configuration  space for  infall.   It  is also  a requirement from  the point  of view of  the non-linear dynamics  within the dark  matter halo  in order  to account  for the  relative time  ordering of accretion events.  % was pointless to describe % continuous infall, haloes are typically not in fully phase mixed equilibria, and  the resulting  fluctuation spectrum  may  seed or  excite the  observed properties of galaxies.  %   In this paper, we aimed at constructing a self-consistent description of dynamical issues  for dark matter  haloes embedded in a moderately active cosmic environment. It relied entirely on the assumption that the statistics of the infall is well-characterized, as described in  \\citet{aubert1,aubert2}, and that the mass of the infalling material (or to a lesser extend that of the fly-by) should be small compared to the mass of the halo. It also assumed that the halo was spherical and static or evolving adiabatically (\\Sec{quasi1}). The emphasis was  on  the theoretical framework, rather than the details of the actual implementation. In other words, we aimed at  describing a self-consistent setting which allows us to propagate the cosmological environment into the core of galactic haloes.   In \\Sec{spherical}, we derived the dynamical equations governing the linear evolution of the induced perturbation by direct infall or tidal excitation of a spherically symmetric (integrable)  stationary dark matter halo. The simplified geometry of the initial state allowed us to focus on the specificities of an open system. Specifically, we  revisited the influence of the external perturbations on the spherical halo, and extended the results of the literature by considering an advection term in the Boltzmann equation. This approach was compared to the classical Green solution in Appendix~\\ref{s:propag}. Note that both the intrinsic properties of the halo, via the distribution function, $F$ (\\Eq{defkk}), and the environment,  via $(s^e,\\psi^e)$ (\\Eq{defhh}),  of  the infall and the tidal distortion were accounted for. Clearly the subclass of problems corresponding to tidal perturbations only will turn out to be easier to implement at first. Appendix~\\ref{s:implementation} presents the details of the angle action variables on the sphere together with an explicit expression for the kernel, $\\bK$,  and carried out a test case implementation of the statistical propagation of an ensemble  of  radial excitations with a  powerspectrum scaling like $\\nu^{-2}$.   In \\Sec{nonlinear}, we derived the non-linear response of the galactic halo to second-order (\\Eq{defdiag2}) in the perturbation (and to order $\\rn$ in Appendix~\\ref{s:apendperturb} together with the corresponding N-point correlation function) to account for tidal stripping and dynamical friction. The dynamics was ``solved'' iteratively, in the spirit of the successful  approach initiated in cosmology by \\citet{Fry} and considerably extended by \\citet{francis}. In particular we presented and illustrated a set of diagrams (\\Fig{diag1},\\Fip{diag2}),  each  corresponding to the contribution of the perturbation expansion. Though the actual implementation of the non-linear theory is going to be CPU intensive, we argue that it will improve our understanding of the competing dynamical processes within a galactic halo. In particular, we discussed how this {explicit theory of non-linear dynamics} provides the setting in which substructure evolution (and destruction) will have to be carried, in order to account for \\eg tidal stripping.  In \\Sec{quasi1} we presented the Fokker-Planck equation governing the quasi-linear evolution of the mean profile of the ensemble {averaged} halo embedded in its cosmic environment. Specifically we showed how the infall, drift and diffusion coefficients (\\Eqs{D2def}{D0def}) are related to the two-point correlation of the tidal field, and incoming fluxes. Appendix~\\ref{s:secularAppendix} {gives} a derivation of this equation from first {principles}, while in the main text, we focused on the bibliographic context and possible applications. The key physical ingredient behind this secular evolution theory was  the stochastic  fluctuation caused by the  incoming cosmic substructures. The key technical assumption was that the two time scales corresponding to the relaxation processes and the dynamical evolution decouple. Hence we could assume a hierarchy in time so  that the distribution function is constant in time when computing the polarization.   Finally, in \\Sec{applications} we considered in turn a few classical probes of the large-scale structures which had been used in the past to constrain the main cosmological  parameters { and } the initial power spectrum, which we transposed to the galactocentric context. Note that these are built upon observables, hence they may be used  to constrain  the boundary power spectrum of the $a_{\\bn}$. Since  \\Eq{defLensing}, \\Ep{defSZ}, \\Ep{defFDM}, and \\Ep{defXB}  involve different combinations of $\\langle a_{\\bn} a_{\\bn'} \\rangle$, they will constrain them at different scales with different biases, which should ultimately allow us to better characterize the power spectrum. This situation is the direct analog of the cosmic situation, where the different tracers (weak lensing, Ly-$\\alpha$ forest, CMB etc..), constrain different scales of the cosmological power spectrum (with different biases). Note also that  our knowledge of the statistical properties of the boundary (via the $b_{\\bn}$ and $c_{\\bn}$ coefficients) together with some assumptions on the equilibrium $F_{0}$ allows us to generate given realizations of the $a_{\\bn}$ as shown in \\citet{aubert1} and therefore virtual observables for any of these data sets, for the purpose of \\eg  validating inverse methods. We investigated the consequences of the infall down to galactic scales and showed how it could  be used to account for the observed distribution of disk properties (spiral winding, warps \\etc). We demonstrated how the analytical model (both linear and non-linear) are quite useful when attempting to``invert'' the observations for the past accretion history a given galaxy. This stems mainly from the fact that perturbation theory provides an explicit scheme for the response of the system, in contrast to the algorithmic procedure corresponding to N-body simulations.  Again let us emphasize that \\Eq{SOL} and its non-linear generalization \\Ep{soln} and \\Ep{SOLQuasi} yield in principle the detailed knowledge of the full {\\sl perturbed} distribution (inside $R_{200}$) at later times. (This is to be contrasted to the situation in N-body simulations where the response of the system is partially hidden by the mean profile of the halo, which requires first identifying substructures (\\citet{aubert0})). Hence we should be in a position to weigh the relative importance of the environment (via $s^{e}$ and $\\psi^{e}$) against the inner properties of the galaxy: the unperturbed distribution function of the halo, $F(\\bI)$, (its level of anisotropy, the presence of a central cusp \\etc) the disk, (its mass, its profile, its distribution function, $F(\\M{J})$ \\etc).  The work presented here derives from the fact that it was  realized that the biorthogonal projection pioneered  by \\citet{Kalnajs2} could be applied order by order to the perturbative expansion of the dynamical equations. Yet this in turn required the knowledge of the relative phases involved in the perturbation, which involves characterizing the properties of the perturber. The characterization only made sense statistically in order to retain the generality of the approach of  \\citet{Kalnajs2}. Hence the emphasis on statistics.    \\subsection{discussion \\& Prospects} \\label{s:prospects}  Our purpose  in this paper was to address in a statistically representative manner dynamical issues on galactic scales. We also advocated using perturbation theory in angle-actions in order to explicitly propagate this cosmic boundary inwards in phase space. As was demonstrated in the paper (and shown quantitatively in \\Sec{statprogex}), this task remains in many respects quite challenging.  One of  the limitations of the above method is the reliance on numerous expansions combined to  the special care required in their implementation. One could argue that this level of sophistication might not be justified in the light of the weakness of some of the assumptions. Indeed, we are limited to systems with spherical geometry whereas galaxies most likely come in a variety of shapes. This assumption could be lifted provided we compute the modified actions of the flattened spheroidal equilibrium using perturbation theory for the equilibrium in the spirit of \\citet{binney}, but implies a higher level of complexity; (it would also require statistically specifying the orientation of the halo relative to the infall, as discussed in part in \\citet{aubert0}). We assumed here that the perturbation was relatively light, which excludes a fraction of cosmic event which might dominate the distribution of some of the observables.     \\Sec{Universal}, \\Sec{applications}  and Appendix~\\ref{s:AppendAppli} presented a few possible applications for the framework described here and in \\citet{aubert1}. These galactic probes would need to be further investigated, in particular in terms of observational and instrumental constraints. The biasing specific to each tracer should be accounted for. The second-order perturbation theory needs to be implemented  in practice together with the diffusion coefficients of \\Sec{quasi1}, following \\Sec{implementation} and extending \\Sec{statprogex}. Similarly, the identification and evolution of substructures within the halo mentioned in  \\Sec{satcount} deserves more work. {In \\citet{aubert0}, we showed that the accretion onto L$^{\\star}$ haloes was anisotropic; the dynamical implication of this anisotropy will require some specific work in the future. }  We will need to demonstrate against N-body simulations the relevance of perturbation theory for dynamical friction; in particular, we should explore the regime in which the second-order truncation is appropriate, and at what cost ? Note that truncated perturbation theory implies that modes will ring forever. At some stage, one will therefore have to address the problem of energy dissipation.  Implementing a realistic treatment of the infalling gas will certainly be amongst the more serious challenges ahead of us. This is a requirement both from the point of view of the dynamics but also from the point of view of converting the above predictions into baryon-dependent observables.  The description of the gas will require a proper treatment of the various cooling processes, which can be quite important on galactic scales. In particular, the thickening of galactic disks is most likely the result of a fine-tuning between destructive processes such as the tidal disruption of compact substructures on the one hand, and the  adiabatic coplanar infall of cold gas within the disk. In fact, the nonlinear theory presented in \\Sec{nonlinear} and \\Sec{quasi1} could  be extended to the geometry of disks to account for the adiabatic polarization towards the plane of the disk.    Note  that we  assumed here  that  transients corresponding  to the  initial conditions where damped out so that  the response of the system was directly proportional to the excitation.  The  underlying picture is that of a calmer past, which in fact is very much in contradiction with both our measurements and common knowledge on the more violent past accretion history of galaxies. Indeed, infalling subclumps will contribute via the external tidal potential at  some earlier  time, and  the larger  the lookback  time,  the relatively stronger the importance  of the perturbation (since the  intensity of infall is  in fact  an increasing  function of  lookback time).   We  are therefore facing   a  partially   divergent  boundary   condition.   Because   of  the characteristics of hierarchical clustering,  the actual bootstrapping of the analytical framework  is therefore challenging.  This could be a  problem in particular for  non linear  dynamics, where the  coupling of  transients may turn out to be as important as the driven response.  The importance of these shortcomings will need to be addressed in the future.     Finally  let   us  note  that  the  theory   described  in  \\Sec{spherical}, \\Sec{nonlinear}  and  \\Sec{quasi1} describe  perturbative  solutions to  the collisionless Boltzmann  Poisson equation in angle-action  variables, and as such are  not specific to the  description of dark matter  haloes.  It could straightforwardly be  transposed to other situations  or geometries provided the system  remains integrable.  As  mentioned in \\Sec{source},  the stellar dynamics around a massive black hole  would seem to be an obvious context in which  this  theory  could be  applied.   For  instance,  we might  want  to investigate the capture of streams of  stars by an infalling black hole.  In a slightly  different context, note in  passing that the  above theory could also  be applied  to celestial  mechanics, since  an  angle-action expansion corresponds to an all eccentricity scheme.   Let us  close this paper  by a summary  of the pros  and cons of  the theory presented here. \\vskip 0.2cm \\vskip 0.2cm \\parn {\\it Possible assets:} \\vskip 0.2cm \\begin{itemize} \\item fixed boundary: localized statistics; \\item fluid description : no  {\\it a priori} assumption on the possibly time- dependent nature of the objects; \\item non-linear explicit treatment of the dynamics: proper account of the self-gravity of incoming objects and statistical  accounting of causality; \\item dynamically-consistent statistically-representative treatment of the cosmic environment; \\item customized description of resonant processes within the halo via angle action variables of universal profile; \\item ability to construct one- and two-point statistics for a wide range of galactic observables. \\item theoretical framework for dynamical inversion and  secular evolution \\end{itemize}  \\vskip 0.2cm \\parn {\\it Possible drawbacks:} \\vskip 0.2cm \\begin{itemize} \\item weak perturbation w.r.t. spherical stationary equilibrium: not representative of \\eg equal mass mergers; \\item complex time dependent 5D boundary condition; \\item {\\it ad hoc} position of the boundary; \\item no obvious truncation of  two-entry perturbation theory; \\item no account of baryonic processes; \\item inconsistency in relative strength of merging events versus time; \\item non-Gaussian environment probably untractable; \\item finite temporal horizon given finite $\\ell_{\\rm max}$; \\item no statistical accounting of linear instabilities. \\end{itemize}  \\vskip 1cm \\parn{\\handfont Acknowledgments}  {\\sl  We   are  grateful  to   E.  Bertschinger,  J.~Binney,   S.   Colombi, J.~Devriendt,    J.~Heyvaerts,   A.~Kalnajs,    J.~Magorrian,   D.~Pogosyan, S.~Prunet,  A.~Siebert, S.~Tremaine, \\&  E.~Thi\\'ebaut for  useful comments and  helpful  suggestions.  We are especially grateful to J. Heyvaerts for carefull reading of the manuscript and for introducing us to the  quasi-linear formalism. We thank the anonymous referee for constructive remarks. Support   from  the  France-Australia  PICS  is gratefully  acknowledged.  DA thanks  the Institute  of Astronomy  and the MPA for their hospitality and funding from a Marie Curie studentship.  }"
    },
    "0512/astro-ph0512125_arXiv.txt": {
        "abstract": "We provide a comprehensive study of the cosmic-ray muon flux and induced activity as a function of overburden along with a convenient parameterization of the salient fluxes and differential distributions for a suite of underground laboratories ranging in depth from $\\sim$1 to 8 km.w.e.. Particular attention is given to the muon-induced fast neutron activity for the underground sites and we develop a Depth-Sensitivity-Relation to characterize the effect of such background in experiments searching for WIMP dark matter and neutrinoless double beta decay. ",
        "introduction": "Underground laboratories provide the overburden necessary for experiments sensitive to cosmic-ray muons and their progenies. Muons traversing a detector and its surrounding material that miss an external veto serve as a background themselves and secondary backgrounds are induced in the production of fast neutrons and cosmogenic radioactivity. In this study we have focused on the muon-induced fast neutron background as a function of depth and the implications for rare event searches for neutrinoless double beta decay and WIMP dark matter. One of our main goals is to develop a Depth-Sensitivity-Relation (DSR) in terms of the total muon and muon-induced neutron flux and to put this into the context of existing underground laboratories covering a wide range of overburden.  In Section II we review the experimental data available for differential muon fluxes and provide a definition of depth in terms of the total muon flux that removes some confusion regarding the equivalent depth of an underground site situated under a mountain versus one with flat overburden. The muon fluxes and differential distributions are parameterized and used as input in Section III to generate, via FLUKA simulations~\\cite{fluka}, the production rate for fast neutrons. The total neutron flux and salient distributions are compared with the available experimental data and we provide some convenient parameterizations that can be used as input for detector-specific simulations at a given underground site. We quantify the agreement between FLUKA simulation and experimental data and provide an explanation for the discrepancy between neutron flux and energy spectra as measured in the LVD detector. Muon-induced cosmogenic radioactivity is discussed in terms of depth and the average muon energy in Section IV. In Section V we apply our results to a generic study of germanium-based experiments in search of neutrinoless double beta decay and WIMP dark matter and demonstrate the utility of the DSR in projecting the sensitivity and depth requirements of such experiments. We conclude with a summary of the results and an outline of new studies underway. ",
        "conclusions": "We have provided a comprehensive study of the cosmic-ray muon flux and salient distributions as a function of depth and specific to a set of existing underground laboratories around the globe. We have applied these distributions to simulate the induced background at various underground sites and, where possible, made direct comparison to the available experimental data in order to assess the accuracy of our predictions. A Depth-Sensitivity-Relation has been developed and applied to examples of germanium-based detectors used in the search for cosmological dark matter and neutrinoless double beta decay.  The cosmic-ray muon flux is well described by a simple exponential law over a broad range in depth extending from about 1 to 8 km.w.e. We have defined depth in terms of the total muon flux obtained at an equivalent vertical depth to a site with flat overburden. This removes some of the confusion regarding the average depth often quoted for laboratories sited beneath mountains where the measured total muon flux is $\\sim$ 15 to 20\\% greater than what would be predicted based upon the average depth alone.  Good agreement can be found between the output of FLUKA simulations and the available experimental data on muon-induced fast neutrons provided one accepts our argument to correct the LVD data on both flux and energy distribution due to quenching effects. In that case we find that our simulations reproduce the data well, albeit with an overall normalization for the total neutron flux that appears to be underestimated by $\\sim$ 35\\%. This normalization appears to be greatly improved when one corrects the output of the FLUKA simulation to agree with experimental data on neutron multiplicity. Clearly, more data on the fast neutron energy spectrum and multiplicity induced by muons would be valuable to further bench-mark and tune the FLUKA simulations.  Our example DSR for dark matter searches is developed based on a model for the CDMS detector and demonstrates that depths in excess of $\\sim$ 5 km.w.e. will be required in order to circumvent background from the elastic scattering of fast neutrons contaminating the low-energy region of interest for recoiling WIMPs. A similar conclusion can be made for neutrinoless double beta decay, modeled after the Majorana detector, where background following nuclear excitation due to the inelastic scattering of fast neutrons is the main culprit. Shallower depths make such experiments feasible provided the fast neutron flux can be adequately shielded and/or actively vetoed. The muon and muon-induced activity increases by approximately one order of magnitude for every decrease in depth of 1.5 km.w.e..  The program developed here has been applied to these specific types of experiments and detector geometries, however, the distributions presented in parameterized form can now be used as input to new simulations and background studies in other detectors of interest. The program could also be easily extended to underground sites under development that have not been considered in this work. More recently, we have begun an experimental program to verify some of our specific predictions by irradiating a Ge-detector with fast neutrons. Preliminary results indicate that the data agree well with our specific predictions for the Majorana detector. The details of that work is beyond the scope of this paper and will be communicated separately."
    },
    "0512/astro-ph0512639_arXiv.txt": {
        "abstract": "A two-dimensional mean field dynamo model is solved where magnetic helicity conservation is fully included. The model has a negative radial velocity gradient giving rise to equatorward migration of magnetic activity patterns. In addition the model develops longitudinal variability with activity patches travelling in longitude. These patches may be associated with active longitudes. ",
        "introduction": "Active regions are complexes of magnetic activity out of which sunspots, flares, coronal mass ejections, and several other phenomena emerge with some preference over other regions. These regions tend to be bipolar, i.e.\\ they come in pairs of opposite polarity and are roughly aligned with the east--west direction.  There is some controversy as to whether or not active regions appear preferentially along so-called active longitudes and what their long term stability properties are (e.g., Bai 2003, Berdyugina \\& Usoskin 2003, Pelt et al.\\ 2005). Some degree of recurrence of sunspots has frequently been reported over the years (Vitinskij 1969, Bogart 1982, Bai 1987, 1988), but only now these ideas are becoming more quantitative.  As the recent analysis of Usoskin et al.\\ (2005) has shown, active longitudes have characteristic angular velocities that depend on the phase of the cycle which, in turn, determines the typical latitude of their occurrence. The analysis of solar magnetograms by Benevolenskaya et al.\\ (1999) showed already that at the beginning of each cycle, when most of the activity occurs at about $\\pm30^\\circ$ latitude, the rotation rate of the active longitudes is $\\Omega/2\\pi\\approx446\\nHz$, while at the end of each cycle, when the typical latitude is only $\\pm4^\\circ$ latitude, the rotation rate of the active longitudes is $\\Omega/2\\pi\\approx462\\nHz$. The recent work of Usoskin et al.\\ (2005) demonstrates that the active regions can also be detected in sunspot data. Unlike the analysis of Benevolenskaya et al.\\ (1999), who isolated two different active longitudes at the beginning and the end of the cycle, Usoskin et al.\\ (2005) determined a continuous latitudinal dependence of these active longitudes on the phase of the cycle. The success of their analysis lies in the way they calculated time-dependent reference values. In particular, they find that only 10\\% of the spots participate in this nonaxisymmetric effect. Thus, the effect is real, but weak, although with a well determined strength.  Given that active longitudes suggest the presence of well preserved activity patches in an otherwise turbulent medium, one has to look for a quantity  that has the capability to be long-lived. An obvious candidate for anything long lived in hydromagnetic turbulence at large magnetic Reynolds numbers is the magnetic helicity. In the absence of magnetic helicity fluxes, magnetic helicity is nearly perfectly conserved in resistive magnetohydrodynamics at large magnetic Reynolds numbers. Even in the presence of magnetic helicity fluxes, magnetic helicity is only transferred from one place to another, but it is not lost. It is therefore plausible that there might be a connection between the nearly perfect magnetic helicity conservation and the long-lived features on the sun such as active longitudes. In the simplest case, magnetic helicity could be considered frozen into the plasma so that local patches of enhanced magnetic helicity would just propagate with the ambient velocity of the gas. This could in principle be a very simplistic picture of active longitudes that might explain their long life times, but not how they came into existence.  An alternative that is traditionally discussed in connection with active longitudes is the idea that there are a number of different axisymmetric and nonaxisymmetric dynamo modes present in the sun that are mixed in the right proportions such that their superposition corresponds to the observed field configuration (R\\\"adler et al.\\ 1990, Moss 1999, 2004, Moss \\& Brooke 2000, Bigazzi \\& Ruzmaikin 2004, Berdyugina et al.\\ 2006). An obvious problem is that these modes are solutions of the linearized problem and that their superposition does not constitute a solution to the nonlinear problem. At first instance only one of the modes gets selected, so the final solution is sill mainly either axisymmetric or nonaxisymmetric, but not easily anything in between, as originally anticipated (R\\\"adler et al.\\ 1990, Moss et al.\\ 1995).  Magnetic helicity conservation only applies to the total field, i.e.\\ the sum of small scale and large scale fields. The large scale field produced by a mean field (large scale) dynamo does by itself not conserve magnetic helicity, because the $\\alpha$ effect leads to a transfer of magnetic helicity from smaller to larger scales (Pouquet et al.\\ 1976, Ji 1999). It is therefore necessary to include also the contribution from the small scale magnetic helicity, which enters into the mean field description through a magnetic contribution to the $\\alpha$ effect. This approach is now fairly well developed and is important in reproducing the slow saturation (Field \\& Blackman 2002, Blackman \\& Brandenburg 2002, Subramanian 2002) found in simulations (Brandenburg 2001). For a review of these recent developments see Brandenburg \\& Subramanian (2005a).  We begin by explaining this approach in more detail. In order to focus on the essentials, we consider only a minimalistic model. The basic dynamo wave can already be described in a one-dimensional model with only latitudinal extent. In the present context we still need the longitudinal extent, but we ignore the variation of the magnetic field with depth. Furthermore, in order to study the basic effect introduced by the longitudinal extent and by magnetic helicity conservation, it suffices to restrict oneself to a cartesian model. Since there is no radial extent, there are also no radial boundaries, and hence no magnetic helicity flux out of them. This might be an important limitation to keep in mind. ",
        "conclusions": "The present investigations can only be regarded as preliminary, because we have not been able to explore the big parameter space given by the large number of unknowns. Most worrisome is the restriction to only small values of the magnetic Reynolds number. At the moment we are experiencing problems when we try larger values, suggesting that the problem may not we well-posed for larger values of $R_{\\rm m}$ and may require modifications. However, this is still surprising, because similar problems have not been encountered in simpler problems, where however no mean advection or shear was taken into account (e.g.\\ Brandenburg \\& Subramanian 2005b).  One surprising aspect emerging from the present simulations is the fact that magnetic patterns do not move with the local gas velocity. Instead, the field propagation is governed by the nonaxisymmetric dynamics of the dynamo modes. Thus, these simulations would not support the notion of field line anchoring. This pictures is occasionally used in connection with sunspot proper motions. Long before the internal angular velocity was determined via helioseismology, it was known that sunspots rotate faster than the surface plasma (Howard et al.\\ 1984). Moreover, young sunspots rotate faster than old sunspots (Tuominen \\& Virtanen 1988, Pulkkinen \\& Tuominen 1998). A common interpretation is that young sunspots are still anchored at a greater depth than older ones, and that therefore the internal angular velocity must increase with depth (see also Brandenburg 2005). This provided also the basis for the classical mean field dynamo theory  of the solar cycle according to which the radial angular velocity gradient has to be negative (Parker 1987).  Given the lack of agreement between the speed of magnetic patches and the local flow speed, one it led to believe that magnetic structures can therefore not be used as tracers of the local flow speed. Alternatively, it is also possible that the tracer properties of the advected field are fully displayed only in three dimensions, or at larger magnetic Reynolds numbers. However, the present results should be interpreted with care, because the present calculations have only been possible in a rather limited parameter range and for rather small magnetic Reynolds numbers. It would be important to test the notion of field line anchoring in direct simulations of the original non-averaged equations, i.e.\\ in the presence of developed turbulence."
    },
    "0512/astro-ph0512313_arXiv.txt": {
        "abstract": "We examine the variability of the high-ionizaton C~\\textsc{iv}$\\lambda{1549}$ line in a sample of 105 quasars observed at multiple epochs by the Sloan Digital Sky Survey. We find a strong correlation between the change in the C~\\textsc{iv} line flux and the change in the line width, but no correlations between the change in flux and changes in line center and skewness.  The relation between line flux change and line width change is consistent with a model in which a broad line base varies with greater amplitude than the line core. The objects studied here are more luminous and at higher redshift than those normally studied for variability, ranging in redshift from 1.65 to 4.00 and in absolute $r$-band magnitude from roughly $-$24 to $-$28.  Using moment analysis line-fitting techniques, we measure line fluxes, centers, widths and skewnesses for the C~\\textsc{iv} line at two epochs for each object.  The well-known Baldwin Effect is seen for these objects, with a slope $\\beta = -0.22$.  The sample has a median intrinsic Baldwin Effect slope of $\\beta_{int} = -0.85$; the C~\\textsc{iv} lines in these high-luminosity quasars appear to be less responsive to continuum variations than those in lower luminosity AGN.  Additionally, we find no evidence for variability of the well known blueshift of the C~\\textsc{iv} line with respect to the low-ionization Mg~\\textsc{ii}$\\lambda$2798 line in the highest flux objects, indicating that this blueshift might be useful as a measure of orientation. ",
        "introduction": "}  Quasar emission lines represent light reprocessed by high-velocity ionized gas surrounding a central continuum source.  As the central source continuum varies, the emission lines vary in response.  Quasar continuum variability is well studied.  Long-known anticorrelations between variability amplitude and luminosity \\citep[e.g.,][]{uomoto76,cristiani97} and between variability amplitude and wavelength \\citep[e.g.,][]{giveon99,trevese01,wilhite05}, as well as the correlation between variability and time lag \\citep[e.g.,][]{hook94,hawkins02,devries03} were recently parameterized by \\citet{vandenberk04}.  Variability of quasar emission lines is also well studied, but for a much smaller number of objects.  Much of this work has come as a part of reverberation mapping efforts to determine black hole masses and study the structure of the broad-line region \\citep[e.g.,][]{peterson93,wandel99}, and has generally focused on line response to continuum variability.  These emission line variability studies have been critical in characterize the structure of the broad-line region.  By measuring the response delays of various emission lines, with respect to the continuum, reverberation mapping has shown that the broad-line emitting region species are stratified by ionization potential \\citep[e.g.][]{peterson93}, and that the size of the BLR is dependent upon continuum luminosity \\citep[e.g.][]{wandel99,kaspi05}.  Until recently, C~\\textsc{iv}$\\lambda{1549}$ had been monitored in only a few low-redshift, low-luminosity objects, such as NGC5548, monitored for years with both IUE and HST \\citep{clavel91,korista95}, and NGC 5141, observed for short-term variability with IUE \\citep{crenshaw96}.  \\citet{kaspi03} have begun a campaign to use C~\\textsc{iv} variability in reverberation mapping of high-redshift, high-luminosity quasars, but the results are not yet conclusive.  Profile variability of AGN emission lines has been studied \\citep[e.g.,][]{wanders96,sergeev01}, but most of this has been done for nearby, low-luminosity Seyferts and has been limited to the rest-frame optical.  The reasons for this are well motivated---low-luminosity objects are known to be more variable and optical spectroscopy is more common---but this means there has been relatively little study of line variability in higher luminosity objects, or of rest-frame ultraviolet lines like C~\\textsc{iv}.  Although not well studied in the time domain, the C~\\textsc{iv} line has demonstrated several intriguing properties at a single epoch which suggest study of C~\\textsc{iv} variability could prove useful in understanding the structure of the broad emission line region and quasars as a whole.  \\subsection{The C~\\sc{iv} Line Profile\\label{introprofile}}  \\citet{wills92} and \\citet{francis92} found that C~\\textsc{iv} line width is anticorrelated with the equivalent width of the line, a result even more clearly demonstrated by \\citet{wills93}.  The highest flux C~\\textsc{iv} lines tend to be the most narrow.  \\citet{wills93} suggested that this might be due to different relative importances of an intermediate width line region (ILR) and the very broad line region (VLBR) from quasar to quasar.  In this scenario, the narrow ($\\sim2000$ km/s) line core is produced in the ILR, which lies near the outer edge of the broad-line region (BLR).  The broad ($\\sim7000$ km/s) line base is a product of the VBLR, which comprises the inner portion of the BLR.  According to the ILR model, possible line core fluxes extend over a larger range in values than the line base fluxes.  Thus, the width of an individual line is strongly dependent upon the strength of the ILR line core.  Strong C~\\textsc{iv} lines have dominant cores and are therefore narrow.  Similarly, weak C~\\textsc{iv} lines have less dominant cores and, thus, relatively more important line bases, and are preferentially broader as a result.  \\citet{murray97} found that the C~\\textsc{iv} profile could be reproduced with a continuous disk-wind model and did not distinguish between intermediate-width and very broad line regions.  \\citet{wills93} also found that the C~\\textsc{iv} line is typically asymmetric, in the sense that C~\\textsc{iv} lines are generally skewed to shorter wavelengths.  \\citet{richards02a} found that this asymmetry tends to increase with increasing C~\\textsc{iv} blueshift; the most blueshifted lines also tend to be the most skewed toward the blue end of the spectrum.  \\subsection{The Baldwin Effect in the C~\\sc{iv} Line\\label{introbaldwin}}  \\citet{baldwin77}, and others later \\citep{kinney90,baskin04} demonstrated that the equivalent width of the C~\\textsc{iv} line is anticorrelated with the luminosity of the nearby continuum for quasars observed at a single epoch:  \\begin{equation} W_{C~\\textsc{iv}} \\propto L_{\\lambda}^{\\beta} \\end{equation}  The initial fit to the slope was $\\beta=-0.64$ \\citep{baldwin78}.  \\citet{kinney90} found a lower value of $\\beta = -0.17 \\pm 0.04$ and later studies \\citep[c.f.,][]{dietrich02} have found similar results.  The Baldwin Effect may also be recast in terms of line luminosity, giving the similar  \\begin{equation} L_{C~\\textsc{iv}} \\propto L_{\\lambda}^b, \\end{equation} where $b=\\beta + 1$. In this form, the relation is easier to understand.  From quasar to quasar, as the continuum luminosity increases, the C~\\textsc{iv} line luminosity increases, but at a slower rate.  By using equivalent width as a proxy for luminosity,   it had originally been hoped that the Baldwin Effect could be used as a cosmological probe.  Unfortunately, the roughly half-magnitude scatter about the original relation is too large to allow for precision cosmology.  \\citet{richards02a} found that the Baldwin Effect in C~\\textsc{iv} appears to be related to the blueshift of the line with respect to lower ionization lines, such as Mg~\\textsc{ii} (see \\S\\,\\ref{introlineshifts}). They separated almost 800 quasars into four equally populated bins, splitting by the size of the blueshift; from the quasars in each bin, they created a composite spectrum.  They found a clear anticorrelation between the strength of the C~\\textsc{iv} line and the size of the blueshift.  The bin with the largest C~\\textsc{iv}-Mg~\\textsc{ii} blueshift also had the lowest equivalent width composite C~\\textsc{iv} line, and vice versa.   It has been suggested that the C~\\textsc{iv} Baldwin Effect could be largely reproduced through a softening of the continuum slope with increasing luminosity and luminosity-dependent quasar metallicity \\citep{korista98}.  \\citet{wang98} found that the ultraviolet to X-ray spectral index is correlated with quasar luminosity: more luminous quasars have softer ionizing continua slopes.  They also found that the UV to X-ray index is strongly correlated with C~\\textsc{iv} equivalent width.  The combination of these effects leads directly to the Baldwin Effect: quasars with high luminosity display low C~\\textsc{iv} equivalent width.  Though these relations are consistent with the Baldwin Effect, the physical driver itself is not yet understood.    Recently, \\citet{baskin04} found that the correlation between C~\\textsc{iv} equivalent width was much stronger with $L^{1/2}(H_{\\beta} FWHM)^{-2}$, a proxy for $L/L_{EDD}$ (since the black hole mass scales as $L^{1/2}(H_{\\beta} FWHM)^{2}$), than it was with the simple continuum luminosity. They have suggested that the Baldwin Effect may in fact be a secondary effect spawned by a more fundamental relation between C~\\textsc{iv} equivalent width and the relative accretion rate, $L/L_{EDD}$. However, the potential physical mechanism driving the relation is unknown.   The roughly half-magnitude scatter in the Baldwin Effect was shown by \\citet{kinney90} to be at least partially due to continuum and C~\\textsc{iv} line variability. As a quasar's continuum luminosity increases or decreases, the C~\\textsc{iv} line luminosity (which consists largely of reprocessed continuum photons) increases or decreases in turn, with a small delay owing to the light travel time.  An intrinsic relationship between continuum and line luminosities may be written in forms identical to those for the global Baldwin relation ($W_{C~\\textsc{iv}} \\propto L_{\\lambda}^{\\beta_{int}}$ and $L_{C~\\textsc{iv}} \\propto L_{\\lambda}^{b_{int}}$).  \\citet{kinney90} found that the so-called ''intrinsic Baldwin Effect'' (IBE) slope ranged from $\\beta$ = -0.4 to -0.9 for 6 Seyfert galaxies and 3C 273 with an average of $\\beta_{int} \\approx -0.65$ ($b_{int} \\approx 0.35$).  The intrinsic Baldwin Effect is, for historical reasons, usually cast in terms of equivalent width, but is more straightforward when expressed in terms of luminosity.  As was the case with the global Baldwin Effect, the slope of the IBE is between 0 and 1.  This indicates that, for an individual quasar, the BLR reprocessing of the incident continuum light is not perfectly efficient.  As the continuum luminosity of an individual quasar fluctuates, so too does the CIV line luminosity, but to a lesser degree.  If an object had an intrinsic Baldwin Effect slope of 0.35, a doubling in a quasar's continuum luminosity would only lead to a roughly 25\\% increase in CIV line luminosity (after allowing for the light-travel time delay).  The intrinsic Baldwin Effect (IBE) slope itself has been found to vary.  Over 13 years of monitoring, the IBE slope of the H$_{\\beta}$ line in the Seyfert I galaxy NGC 5548 ranged from $b=0.4$ to 1.0 on time scales of roughly one year \\citep{goad04}.  The slope was strongly anti-correlated with continuum flux, indicating a lower line responsivity at higher continuum flux levels, which is consistent with photoionization models \\citep{korista04}.  \\citet{pogge92} determined that the Baldwin Effect scatter may be further reduced by accounting for the light travel time, $\\tau$, between the continuum source and the broad emission line region: $L_{C~\\textsc{iv}}(t) \\propto L(t-\\tau)^{\\beta}$.  \\subsection{C~\\sc{iv} Line Shifts\\label{introlineshifts}}  \\citet{gaskell82} first demonstrated that high-ionization quasar broad emission lines (such as C~\\textsc{iv}) are typically blueshifted by hundreds of kilometers per second with respect to the low-ionization lines, thought to represent the true systemic redshift of the quasar.  This was verified in a number of later studies \\citep{wilkes84,espey89,corbin90,tytler92,mcintosh99,sulentic00}.  Recently, \\citet{richards02a} measured the blueshift of the C~\\textsc{iv} line with respect to Mg~\\textsc{ii} for $\\sim800$ quasars in the SDSS Early Data Release Quasar Catalog \\citep{schneider02}.  A possible correlation between C~\\textsc{iv} blueshift and radio-determined orientation measures, as well as a similarity between the spectra of broad absorption line quasars and quasars with large C~\\textsc{iv} blueshifts, prompted \\citet{richards02a} to suggest the possibility that C~\\textsc{iv} blueshift could be used as a measure of quasar orientation, either internal (related to the disk wind opening angle) or external (related to the line of sight to the observer). They proposed that the blueshift might be a result of the obscuration or suppression of the C~\\textsc{iv} flux on the red side of the line.  If the blueshift of the C~\\textsc{iv} line, relative to low-ionization lines like Mg~\\textsc{ii}, is related to the observer's viewing angle, it could represent the first technique to measure orientation for radio-quiet quasars.  \\subsection{The Present Work\\label{presentwork}}  This is the second paper reporting results of a quasar spectral variability program using data from the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}.  The first paper \\citep[][hereafter Paper I]{wilhite05} examined the detailed wavelength dependence of quasar variability.  This paper focuses on the high-ionization C~\\textsc{iv}$\\lambda1549$ line.  We briefly summarize the SDSS data acquisition, our previous spectrophotometric re-calibration work, and the creation of the variable quasar sample in \\S\\,\\ref{dataset}.  In \\S\\,\\ref{linefitting}, we describe the line-fitting algorithm used here.  The variability of the C~\\textsc{iv} line flux and profile is studied in \\S\\,\\ref{civvariability}.  Interesting individual objects are identified in \\S\\,\\ref{individual}.  The results are discussed in \\S\\,\\ref{discussion} and we conclude in \\S\\,\\ref{conclusions}.  Throughout the paper we assume a flat, cosmological-constant-dominated cosmology with parameter values $\\Omega_\\Lambda = 0.7, \\Omega_{M} = 0.3,$ and $H_{0}=70$km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "}  Using a sample of 105 quasars observed multiple times by the Sloan Digital Sky Survey, we have studied the variability of the C~\\textsc{iv} line.  Spectra were fit using moment analysis techniques and four main conclusions are drawn:  1)  We find a strong correlation between the change in C~\\textsc{iv} line flux and the change in line width.  As an individual quasar's C~\\textsc{iv} line flux increases, so does the C~\\textsc{iv} line width.  This is consistent with any picture of the BLR in which the broad line base is produced nearer to the central engine and the portion of the BLR nearer to the central engine exhibits more coherent line flux variability.  2)  We demonstrate that there is no apparent variability in the the blueshift of the C~\\textsc{iv} line with respect to the Mg~\\textsc{ii} line for the highest flux C~\\textsc{iv} lines, a possibly positive sign for the use of line shifts as an orientation measure.  3)  With our measurements of continuum and line fluxes, we are able to reproduce the Baldwin Effect, deriving a slope of $b=0.78$.  We also calculate a median slope for the Intrinsic Baldwin Effect of $b_{int}=0.15$, shallower than the $b_{int} \\approx 0.35$ determined by \\citet{kinney90} for lower luminosity AGN.  4)  Using the continuum flux at the position of the line center, we reproduce well-known dependences of continuum variability amplitude on quasar luminosity and rest-frame time lag.  However, these same dependences are not evident for the amplitude of the C~\\textsc{iv} line variability.  This may be due to the \"smearing out\" of continuum variability by the extended BLR.  Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is {\\tt http://www.sdss.org/}.  The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0512/astro-ph0512549_arXiv.txt": {
        "abstract": "We review the X-ray spectra of the cores of clusters of galaxies.  Recent high resolution X-ray spectroscopic observations have demonstrated a severe deficit of emission at the lowest X-ray temperatures as compared to that expected from simple radiative cooling models.  The same observations have provided compelling evidence that the gas in the cores is cooling below half the maximum temperature.  We review these results, discuss physical models of cooling clusters, and describe the X-ray instrumentation and analysis techniques used to make these observations.  We discuss several viable mechanisms designed to cancel or distort the expected process of X-ray cluster cooling. ",
        "introduction": "Observations show that the X-ray emission from many clusters of galaxies is sharply peaked around the central brightest galaxy. The inferred radiative cooling time of the gas in that peak, where the temperature drops to the center, is much shorter than the age of the cluster, suggesting the existence of a cooling flow there \\cite{fabian94}. X-ray spectroscopy over the past 5 yr shows that the temperature drop toward the center is limited to about a factor of three. Just when the gas should be cooling most rapidly it appears not to be cooling at all. This is sometimes known as the cooling flow problem. Careful observations show that gently distributed heat is required over a radius of up to 100~kpc to balance radiative cooling in these regions.  The issues of cooling and heating of hot gas have broad relevance to the gaseous part of galaxy formation and evolution. Brightest cluster galaxies (BCG) are the most massive galaxies known. Calculations of the clustering behaviour of cold dark matter predict a power-law mass distribution for large galaxies whereas the stellar mass observed has an exponential distribution \\cite{benson03}. The truncation of the stellar mass distribution in massive galaxies is likely due to the process which stops cooling flows. Simple cooling flows are an ingredient of semi-analytical models for galaxy formation. The cooling of hot gas to form stars is essential for the growth of massive galaxies and cannot be studied directly for isolated systems due to Galactic absorption. The cores of galaxy clusters offer examples which can be directly observed. However they do not appear to operate in any simple manner.  The problem appears to be widespread, from the most massive clusters to the centers of individual elliptical galaxies. Heating and cooling problems of hot gases are common in astronomy, with examples ranging from the interstellar medium of our own Galaxy to the Solar Corona.  The diffuse hot ionized plasma in clusters is magnetized which means that MHD processes may be important \\cite{schekochihin04}.  Here we briefly review the main X-ray properties and emission processes of the intracluster medium (ICM) before showing the X-ray spectra of cool cores. We then discuss the main solutions which have been proposed for the cooling flow problem.  \\newpage ",
        "conclusions": "The simplest explanation for the common appearance of cold core, X-ray peaked clusters is that, when averaged over tens of Myr, the radiative cooling is balanced in part by distributed heating. Thermal conduction as a means of distributing heat from outer gas is ruled out for low and intermediate temperature clusters. It may however have a role in spreading the energy in the central parts.  A plausible mechanism is the dissipation of energy propagating through the ICM from a central radio source. Such a process stems massive cooling onto the BCG which would otherwise gain a total stellar mass $\\gg10^{12}~{\\rm M_{\\odot}}$. The process is therefore a vital ingredient in stopping the growth of the most massive galaxies \\cite{fabian02b,benson03,binney04}. NOte that most semi-analytic models for galaxy formation (e.g. Kauffmann et al 1999) already needed to suppress cooling in massive haloes in order to match observation.  Difficulties and doubts remain with regard to the issues of the energy dissipation and distribution processes which are tied in with the transport processes in the gas. Similarly, it is not clear how the feedback manages to produce such similar cooling time profiles in systems where temperatures and thus masses differ by over an order of magnitude?  There still remains the possibility that some process not yet foreseen, or at least not well studied will eventually prove more important than the effect of the central radio source, or will at least be important in mediating its effect. The central radio source is so common and so energetic however that it must at least be part of the solution. Similarly, the motions of galaxies or interacting dark matter, if it exists, could be important in heating cluster cores along with the AGN. Given the wide range of objects in which a balance is required, we suspect that a single mechanism is dominant, rather than several.  The need for a heating-cooling balance is in a time-averaged sense, over intervals of about $10^8~{\\rm yr}$. In most cases the heating has not been so energetic as to drive gas out of the inner regions nor so weak as to allow much cooling at very high rates. Examples of objects at the extremes are Hydra~A and Cygnus~A where heating is high (but is dumped mostly at large radii outside the cool region) and A2597 and RXC1504.1-0248 where cooling appears to be high (in the latter object over 70 per cent of the total X-ray luminosity emerges from the cool core; \\citett[boehringer05]).  In most objects residual cooling at a rate of about 10 per cent of the simple unheated cooling rates appears to occur. It could be larger if non-radiative cooling, due to mixing say, is occurring. Stars form from the cooled gas giving the excess blue light seen in the BCGs. Mass loss from such stars can make the cooled gas dusty and radio bubbles drag some of it back out to large radii.  The time evolution of the heating / cooling balance is little understood. We suspect that the common temperature drop associated with the short central cooling times {\\em is} due to radiative cooling and that heating only came into balance when the overall temperature structure was in place. Perhaps the central galaxy and its central BH grew until balance was achieved, and growth required cooling. Comparison of samples of clusters at $\\bar z=0.22$ with $\\bar z=0.056$ \\cite{bauer05} shows little surprisingly change in the distribution of cooling times at 50 kpc. The results imply that any balance was established well beyond $z\\sim 0.3$.  Few massive cool core clusters are known at much higher redshifts than that of RXJ1347 at $z=0.44$. This may however be a selection effect. They will be absent from X-ray cluster samples if the central BH is a bright X-ray source such as a quasar. If a central quasar outshines the host cluster in X-rays then the object will generally be classified as a quasar. H\\,1821+643 is a good example of a bright quasar in an X-ray peaked cluster at intermediate redshift ($z=0.297$, \\citett[fang02]).  Searches for clusters around powerful radio-loud quasars and galaxies have found some examples at $z=0.5-1.1$ \\cite{worrall01,crawford03,siemiginowska05} but no complete searches have been done.  \\clearpage"
    },
    "0512/astro-ph0512063_arXiv.txt": {
        "abstract": "Near-infrared imaging of the emission from molecular hydrogen is a powerful method for discovering outflows in star-forming regions. We present new near-infrared images, long slit and integral field spectroscopy of the ultra-compact \\textsc{H~ii} region G25.65+1.05. These new observations reveal shocked ${\\mathrm H_2}$ emission associated with a bipolar outflow from a young high mass star at the centre of the source. The physical parameters of the outflow are discussed and compared with outflows from lower mass stars. ",
        "introduction": "A number of recent surveys have invigorated the study of high mass star formation by providing data on coherent samples of candidate high mass young stellar objects (HMYSOs) (e.g. \\citealt{palla91}, \\citealt{molinari96}, \\citealt{molinari98}, \\citealt{ridge01}, \\citealt{sridharan02}). Until these surveys became available the relative scarcity of HMYSOs, the large distances to the closest examples (a few kpc) and the high extinction to these objects had limited their study to a few, well-known examples. Thus, there are many questions still remaining as to the nature of high-mass star formation, and the similarities and differences in high and low mass stellar evolution.  Outflows are associated with the formation of stars of all masses and are the subject of a number of the HMYSO surveys (\\citealt{shep2}, \\citealt{shep1}, \\citealt{zhang01}). While the role of outflows in the formation of low-mass stars is comparatively well studied (see for example the review by \\citealt{richer00}), far less is known about the properties and role of outflows associated with the formation of high-mass stars (M greater than $\\sim$~8~M$_{\\odot}$). A radio survey of molecular line emission from high-mass star forming regions showed that high-velocity molecular gas (CO) is associated with around 90\\% of these regions (\\citealt{shep1}). Two sources (G25.65+1.05 and G240.31+0.07) were mapped at higher spatial resolution and found to have bipolar outflows. In a follow-up survey, \\citet{shep2} mapped a further ten of the best candidates for having bipolar molecular outflows and confirmed the presence of such an outflow in 5 of those sources. \\citealt{zhang01} observed 69 candidate high mass protostars in the CO J=2-1 transition and argue that as many as 90\\% of the sources may have outflows. \\citet{beuther02a}, mapping at a higher spatial resolution, found evidence of bipolar outflows in 21 of their 26 sources suggesting that bipolar outflows may be associated with most young high mass stars. The first report of molecular hydrogen line emission associated with high mass star formation was from \\citet{lee01}, following up a search for the near-infrared counterparts of ultra-compact H~{\\sc ii} regions by \\citet{walsh99}. \\citet{lee01} found evidence for outflows in IRAS~15278-5620 and IRAS~16076-5134.  These surveys confirm the importance of outflows for HMYSO evolution. Further work is required to determine whether the physics of the HMYSO outflows is the same as for low mass YSOs. The observed outflows appear to have much lower collimation factor -- between 1 and 1.8~-- than those seen from low-mass stars which often have a collimation factor of around 10 (see, for example, the interferometric observations of \\citet{richer00}). This would be hard to explain if the outflows are formed by the same jet entrainment model as that believed to describe the outflows from low-mass stars.  However, \\citet{beuther02a} argued that the observed degree of collimation could be significantly reduced by the low spatial resolution of the maps, which could be consistent with well collimated high-mass flows. Interferometric observations have shown collimation factors as high as 10 \\citep{beuther02b} in flows from high-mass stars and have revealed that some of these apparently uncollimated flows can be resolved into several well collimated flows from separate young stars. Detection of multiple flows in a region in which high-mass stars are forming would not be surprising. High-mass stars are known to form in dense clusters \\citep{garay99}, so high spatial resolution observations are essential for identifying the source of an individual outflow. The presence of a collimated outflow would imply the presence of a stable accretion disk and hence strengthen the view that high mass stars are formed by steady accretion, in a similar way to low mass stars, rather than by merging of intermediate-mass protostars in the centre of dense clusters.  We have embarked upon a near-infrared survey of HMYSOs, which will be presented in full in Varricatt et al. (2005, in preparation). The survey is designed to reveal outflows by mapping at sub-arcsecond spatial resolution in the v=1-0 S(1) line of molecular hydrogen. NIR observations of $\\mathrm H_2$ have frequently been used to study outflows from low mass YSOs. In this paper we present near-IR imaging, integral-field and long-slit spectroscopic observations of new outflows in the region of the ultra-compact \\textsc{H~ii} region G25.65$+$1.05 .  G25.65$+$1.05 (also IRAS 18316$-$0602 or RAFGL7009S) is an irregular, compact radio source, first identified at 3.6cm by \\citet{kurtz94}. It is located at a distance of 3.2kpc \\citep{molinari96}. The radio peak is coincident with an unresolved infrared source, identified as a young B1V star with a large K-band excess \\citep{zavagno02} and is also closely associated with methanol \\citep{molinari96} and ammonia maser emission (\\citealt{walsh03}, \\citealt{szymczak00}). Submillimetre continuum observations at 350$\\mu$m \\citep{hunter00}, 450$\\mu$m and 850$\\mu$m \\citep{walsh03} are all peaked at the position of the radio and maser sources. Observation of the CS (2-1) line by \\citet{bronfman96} shows an excellent match of the observed radial velocity from the masers (40.8-42.4~kms$^{-1}$, \\citealt{walsh03}) and the line emission (41.4~kms$^{-1}$) indicating a strong link between the dense gas, the maser sources and the massive star. \\citealt{zavagno02} propose that one explanation for the observed $K$-band excess from the central source may be the presence of a disk. ISO spectroscopy of this source shows a rich spectrum of ice features including absorption features attributed to H$_2$O, CH$_3$OH, CO$_2$,$^{13}$CO$_2$,CO, OCS, HCOOH, HCOO$^-$, CH$_3$HCO, CH$_4$, NH$_3$ and Silicate (see \\citealt{gibb04} and references therein). Laboratory spectra fitted to the ISO observations suggests that the ice features arise in dense material with temperatures in the range 10K-100K. ",
        "conclusions": "\\subsection{Morphology}  The K--band image of G25.65+1.05 (Figure~1a) is dominated by the nebulous continuum emission surrounding the central, unresolved, source at $\\alpha=18^{\\rm h}34^{\\rm m}20.9^{\\rm s}, \\delta=-6^\\circ02'42.3''$ seen by Zavagno et al. (2002). The morphology of the $\\mathrm{H_2}$ emission in Figure~4 is strongly suggestive of a moderately well-collimated outflow, or outflows, centred on G25.65+1.05. Sources A, B, C and D are connected by diffuse emission and appear to form one lobe of an outflow with a flow axis through G25.65+1.05 to source E. The outflow has a position angle on the sky of 130$^\\circ$ East of North. This is consistent with the direction of the highly energetic bipolar outflow found to be centred on or close to the \\textsc{H~ii} region by \\citet{shep2}. Sources ABCD are spatially co-incident with the blueshifted lobe seen in the CO outflow and E with the redshifted lobe. The relative faintness of E compared with ABCD may then be interpreted as being due to E being embedded further in the molecular cloud. The projected length of the outflow, taken to be the distance from C to E, is $\\sim$70arcsec. In the absence of any information, we assume an angle relative to the plane of the sky of 45$^\\circ$. This gives an estimate of the total length of 1.4pc and a collimation factor $\\sim$3. Interpreted as a single outflow, we find a lower degree of collimation than is seen in outflows from low-mass YSOs, but higher than the 1-1.8 often assumed for high-mass YSOs.  \\citet{shep2} found two velocity components associated with G25.65+1.05 at 39km${\\rm s}^{-1}$ and at 49km${\\rm s}^{-1}$. One possible interpretation of the $\\mathrm{H_2}$ emission in this region would be to see sources A and B as bow-shocks in a highly collimated flow from west to east, as suggested by their morphology. This would imply the presence of another, more deeply embedded source to the south and west of G25.65+1.05. The MSX point source catalogue \\citep{Egan98} shows a single source associated with G25.65+1.05 ($\\alpha=18^{\\rm h}34^{\\rm m}21.2^{\\rm s}, \\delta=-6^\\circ02'36.96''$ with a positional accuracy of $4''-5''$) so does not support an argument for a second source offset south-west of G25.65+1.05. Furthermore, a recent submillimetre survey by \\citet{hill05} finds a single core of mass 1.$\\times10^3 \\rm M_{\\odot}$ associated with G25.65+1.05, located at the position of the methanol maser \\citep{molinari96}, consistent with an earlier result by \\citet{faundez04}.  In the following Section, the near-IR spectra of sources ABCD are analysed and shown to support our model of the emission arising due to a single outflow.  \\subsection{$\\mathrm H_2$ Excitation in the G25.65+1.05 Outflow}  The near-infrared emission from $\\mathrm{H_2}$ can be produced either by thermal excitation in shock-fronts or fluorescent excitation by non-ionizing ultraviolet photons ($91.2 < \\lambda < 110.8$~nm) from hot young stars. These cases may be distinguished using the near-IR spectrum, though the detection of lines from a broad range of energy levels is required to obtain a secure result. For shock-excited $\\mathrm{H_2}$, the lower energy levels (v=1) are typically populated as for a gas in local thermal equilibrium (LTE) with a characteristic excitation temperature of a few thousand Kelvin, though in reality the temperature of the cooling gas will vary over a range from a few hundred to a few thousand Kelvin. The emission lines would also be broadened to tens of kms$^{-1}$. This is less than our instrumental resolution and undetectable in the spectroscopy of G25.65+1.05.  For radiatively excited gas, the populations follow a non-LTE distribution, characterised by high excitation temperatures ($\\sim10,000$K). Lines from high vibrational levels (v$>>1$) may be detected and the ratio of the intensity of the $1-0~\\mathrm{S}(1)$ line to that of the $2-1~\\mathrm{S}(1)$  $\\mathrm{H_2}$ line is $\\sim 2$ if the density is low \\citep{sternberg89}.  For gas above the critical density (${\\rm n}_{\\rm H_2}\\sim 10^5~\\mathrm{cm^{-3}}$), collisional de-excitation becomes important and the level populations tends towards an LTE population.  The spectra of G25.65+1.05 provide high signal/noise observations of lines from the v=1 and v=2 vibrational levels. Deeper observations would be required to obtain line strengths for the higher vibrational levels and thereby to confirm the excitation. However, there are indicators that the $\\mathrm{H_2}$ is shock excited which, taken in aggregate, lead us to conclude that this is indeed the case.  In the following sections, we obtain the excitation temperatures for the sources in G25.65+1.05. These are in the range 1000-3000K indicative of thermal excitation. When molecular hydrogen is radiatively excited, 90$\\%$ of the $\\mathrm{H_2}$ is dissociated by the ionizing photons ($\\lambda < 91.2$~nm) that are also present. This may result in emission in the 2.166~\\micron\\ Brackett-$\\gamma$ recombination line. The spectra of sources ABCD contain only emission lines of $\\mathrm{H_2}$. We note that the absence of Brackett-$\\gamma$ emission is not a strong constraint as some fluorescently excited photodissociation regions emit only $\\mathrm{H_2}$ lines. Finally we cite the bow-shock morphology seen clearly in Figure~4, perhaps the strongest indicator that the $\\mathrm H_2$ is excited by a shock front driven by an outflow rather than heated from within the knots by embedded point sources.  \\subsubsection{The Excitation Temperature of Sources A and B}  The measured intensity, $I$, of a given $\\mathrm{H_2}$ line can be used to calculate the column density of the upper excitation level of the transition: \\begin{equation} N_j = \\frac{4 \\pi \\lambda_j I}{A_j hc} \\label{col_dens_eq} \\end{equation} where $A_j$ is the Einstein $A$-coefficient of the transition. The relative column densities of any two excitation levels can be expressed in terms of an excitation temperature $T_{\\mathrm{ex}}$: \\begin{equation} \\frac{N_i}{N_j} = \\frac{g_i}{g_j} \\exp \\left(\\frac{-(E_i - E_j)}{kT_{\\mathrm{ex}}} \\right) \\label{thermal_eq} \\end{equation} where $g_j$ is the degeneracy and $E_j$ is the energy of the level. The values of $\\lambda_j$, $E_j$, $A_j$ and $g_j$ for the lines detected in our spectra are shown in Table~1.  Before we could use our measured intensities to derive the excitation temperature of the gas it was necessary to measure and compensate for extinction.  In the absence of more information it was assumed that the extinction was constant across the IFU field of view. An extinction law of the form $\\tau(\\lambda) = A_k\\left(\\lambda/2.2 \\micron\\right)^{-1.75}$ was used, giving a corrected intensity of $I_\\mathrm{corr} = I/e^{-\\tau(\\lambda)}$. Plotting $\\log(N_j/g_j)$ against $E_j$ should give a straight line for low vibration levels and for a single characteristic excitation temperature in each pixel. If the value of $A_k$ used to correct the line intensities from which the column density is calculated is wrong then the scatter of the points will increase. Our measurements are particularly sensitive to this because we have measured one \\textit{H}-band line ($1-0~\\mathrm{S}(7)$ at 1.748~\\micron) which comes from an upper level close in energy to the upper level of two \\textit{K}-band lines ($2-1~\\mathrm{S}(1)$ at 2.248~\\micron\\ and $2-1~\\mathrm{S}(3)$ at 2.074~\\micron).  The value of $A_k$ which minimised the scatter of the points was measured for each spatial pixel in the two brightest columns in the rebinned data-cube. It was found that the intensities of the $1-0~\\mathrm{S}(3)$ line and all the Q-branch lines other than the $1-0~\\mathrm{Q}(4)$ line were not consistent with any non-negative values of $A_k$.  These lines were therefore assumed to be partly absorbed by the very narrow atmospheric lines which dominate the edges of the \\textit{K}-band window which, being spectrally unresolved, are not removed by dividing by the standard star. Once these lines were excluded the mean value of $A_k$ over all the spatial pixels used was found to be $A_k = 0.7 \\pm 0.1$, where the error on the derived value was estimated from the scatter of the values measured from one spatial pixel to another. The variation in the measurements appeared random, with no evidence for any systematic variation in extinction across the source. There was also no evidence of curvature in the $N_j/g_j$ versus $E_j$ plots, one of which is shown in Figure~6, though the absence of any points with upper energy levels between 8000 and 12000~K would make it hard to detect small deviations from a straight line. A curved line in this plot would provide evidence for non-LTE processes, such as would be seen for radiatively excited gas.  \\begin{figure} \\centering \\includegraphics[width=8cm]{images/corrected.eps}  \\caption{Plotting $N_j/g_j$ against $E_{\\mathrm{upper}}$ on a logarithmic scale gives a straight line after correcting for extinction using the extinction law described in the text. The solid line is given by Equation~(3) using the maximum likelihood values of $T$ and $\\alpha$. These measurements are taken from the 0.48arcsec$\\times$0.48 arcsec region of the source marked 'd' Figure~4.} \\label{col_dens_plot} \\end{figure}  The line fluxes were measured at each spatial pixel in the rebinned data-cube and corrected for extinction using the extinction law described above. The value of $N_j/g_j$ was then derived from each corrected line flux using Equation~(1). The excitation temperature was measured by fitting Equation~(2) to all of the measurements made at a single spatial pixel using a maximum likelihood method. When a single intensity is used there is a degeneracy between $T$ and the constant of proportionality, causing the most likely value of the constant of proportionality to vary exponentially with $T$. For this reason Equation~(2) was reformulated as \\begin{equation} \\frac{N_j}{g_j} = \\exp \\left(\\frac{-E_j}{kT} + \\alpha \\right) \\label{thermal_eq2} \\end{equation} where $\\alpha$ is the logarithm of the constant of proportionality.  The likelihood of a range of values of $T$ and $\\alpha$ was calculated for each measured line flux assuming a gaussian error distribution on the measurements of the line intensity and using a uniform prior probability density function for both parameters: \\begin{equation} L_j(T, \\alpha) \\propto  \\frac{\\exp[{-(\\delta/2\\sigma)^2}]}{\\sigma \\sqrt{2\\pi}} \\label{likelihood}. \\end{equation}  In Equation~(4) $\\delta$ is the difference between the value of $N_j/g_j$ derived from the measured intensity using Equation~(1) and that calculated using Equation~(3) and $\\sigma$ is the $1\\sigma$ error on the measured value of $N_j/g_j$.  When this is evaluated for a single spectral line there is a degeneracy between the two parameters, as shown in Figure~7a. The likelihood of the parameters using all the line fluxes is calculated by taking the product of all the individual likelihoods: \\begin{equation} L(T, \\alpha) = \\prod_j{L_j(T, \\alpha)} \\end{equation} This reduces the degeneracy (Figure~7b).  The likelihood of each value of $T$, taking into account the uncertainty in $\\alpha$, can be calculated by marginalising the likelihood, or summing over all values of $\\alpha$ to produce a one dimensional likelihood curve (Figure~8). The 68\\% confidence levels (the narrowest possible range of values of $T$ to include 68\\% of the total area under the likelihood curve) were measured as an equivalent of $1\\sigma$ error bars on a gaussian distribution.  \\begin{figure} \\centering \\includegraphics[width=7cm]{images/likelihood1.eps} \\includegraphics[width=7cm]{images/likelihood2.eps} \\caption{(a) The likelihood of the parameters $T$ and $\\alpha$ were evaluated for a single spatial pixel using the intensity of a single spectral line. (b) When other spectral lines are introduced the degeneracy is partly broken.  On both plots the contours are at the 68\\%, 95\\% and 99\\% confidence levels, equivalent to $1\\sigma$, $2\\sigma$ and $3\\sigma$ error bars on a gaussian distribution.} \\label{likelihood_1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=7cm]{images/likelihood3.eps}  \\caption{The marginal distribution of $T$ can be obtained by integrating the distribution shown in Figure~7b over $\\alpha$. The likelihood shown here is normalised to have a maximum value of 1. The maximum likelihood value of $T$ (2178~K in this case) and confidence intervals can be measured from this. The 68\\% confidence interval of 1973 -- 2371~K is shown by the dotted lines.} \\label{likelihood_2} \\end{figure}   These measurements were carried out along three 0.48~arcsec wide strips marked a, b and c on the image shown in Figure~4. The results are shown in Figure~9.  The excitation temperature clearly increases from around 1800~K to $(2840 \\pm 230)$~K (averaging over the most easterly pixel of all three slices) at what may be the shock front. In source A there also appears to be a decrease in excitation temperature of $\\sim 200$~K to the right (east) in the image shown in Figure~4.  \\begin{figure} \\centering \\includegraphics[width=7cm]{images/slices.eps}\\\\ \\caption{The temperature was measured along three 0.48~arcsec wide strips running west-east on the left of Figure~4. The points and error bars on these plots show the maximum likelihood value and 68\\% confidence interval on the value of $T$ marginalised over $\\alpha$, as described in the text.} \\label{temperatures} \\end{figure}  In Figure~4 it can be seen that both sources are brighter on the left (south) side than on the right. The temperatures measured in source A are approximately 200~K higher on the south edge of the IFU field than 1.5~arcsec to the north. We interpret this asymmetry in the north-south direction as supporting our interpretation of the $\\mathrm{H_2}$ emission being associated with a single outflow. If sources A and B are were produced by a flow from west to east a symmetrical bow-shock would be expected. However, we cannot exclude the possibility that the asymmetry is produced by variations in density in the medium with which the outflow is interacting.  \\subsubsection{The Excitation Temperature of Sources C and D}  The spectra of sources C and D are shown in Figure~3 and the relative line intensities in Table~2. Although the weather was non-photometric, both sources were observed simultaneously and the line ratios may be considered to be accurate, at the limit of the gaussian errors. The maximum likelihood analysis was repeated for C and D, giving an excitation temperature of 2730~K with 68\\% confidence levels of (2727$\\pm$227)K for source C and (2256$\\pm$250)K for source D, in each case averaged over the observed extent of the source.  \\begin{table*} \\begin{minipage}{60mm} \\begin{center} \\caption{Line ratios from the spectra of sources C and D.} \\begin{tabular}{| l | c | c |} \\hline Line &  \\multicolumn{2}{c}{Ratio to 1-0 S(1)} \\\\ &  Source C  & Source D \\\\ \\hline 1-0 S(1) & 1 & 1 \\\\ 2-1 S(1) & 0.15 $\\pm$ 0.02  & 0.10 $\\pm$ 0.04 \\\\ 1-0 S(0) & 0.20 $\\pm$ 0.05  & 0.36 $\\pm$ 0.07 \\\\ 2-1 S(3) & 0.22 $\\pm$ 0.04 & 0.15 $\\pm$ 0.05 \\\\ 1-0 S(2) & 0.28 $\\pm$ 0.08  & 0.52 $\\pm$ 0.08 \\\\  \\hline \\end{tabular} \\end{center} \\medskip  The ratio of line fluxes measured in the CGS4 spectra of sources C and D relative to the 1-0 S(1) line flux. The spectra were observed in cloudy conditions, so the absolute fluxes are unknown.  \\label{cgs4tab} \\end{minipage} \\end{table*}   \\subsubsection{Shock excitation}  The properties of four classes of shocks were summarised by \\citet{davies00}: \\begin{enumerate} \\item \\textit{fast J shock} (100--300~$\\mathrm{km~s^{-1}}$): hydrogen molecules are dissociated and reform producing a spectrum similar to that of UV fluorescence. The flux of Brackett-$\\gamma$ and $H$-band Fe~\\textsc{ii} lines are comparable to the $1-0~\\mathrm{S}(1)$ flux.  \\item \\textit{slow J shock}: strong $\\mathrm{H_2}$ lines are produced and molecules are not dissociated. However, in normal ionization fractions, magnetic field strengths and gas densities, such shocks are expected to evolve into C-type shocks.  \\item \\textit{fast C shock}: shock velocities of $\\sim 40~\\mathrm{km~s^{-1}}$ heat the gas to $\\sim 2000$~K or more (dependent on the velocity) producing strong $\\mathrm{H_2}$ lines.  \\item \\textit{slow C shock}: peak temperature 300~K or less, resulting in very weak $\\mathrm{H_2}$  emission.  \\end{enumerate}  The absence of Brackett-$\\gamma$ emission from atomic hydrogen or Fe~\\textsc{ii} lines and the excitation temperatures measured in all the sources ABCD leads us to suggest that the fast C shock model may be most appropriate. A maximum post-shock excitation temperature of $(2840 \\pm 230)$~K would be produced by a shock velocity of $(37\\pm3)~\\mathrm{km\\:s^{-1}}$ in the models of \\citet{kaufman96}. This would be consistent with the $29~\\mathrm{km\\:s^{-1}}$ along the line of sight measured by \\citet{shep2} when the uncertainty on the angle of incidence is taken into account. The temperatures measured in source A (offsets along the slice of between 4.5 and 6.5~arcsec) decrease with distance from the shock front.  The C-shock models of \\citet{flower96} predict a steady decrease in temperature over around 0.03~pc (2~arcsec at a distance of 3~kpc) behind the shock front, beyond which point the flux falls rapidly. This is consistent with what we see, though the proximity of source B and the uncertain geometry of the outflow inhibits more detailed comparison with the models. A more detailed analysis, including lines from higher vibrational levels examined as function of distance across the shock-front or high spectral resolution to explore the lines shapes would be required to allow us to draw stronger conclusions. Integral field spectroscopy is a technique that would lend itself readily to these observations."
    },
    "0512/astro-ph0512255_arXiv.txt": {
        "abstract": "We consider the structure of self-gravitating marginally stable accretion disks in galactic centers in which a small fraction of the disk mass has been converted into proto-stars. We find that proto-stars accrete gaseous disk matter at prodigious rates. Mainly due to the stellar accretion luminosity, the disk heats up and geometrically thickens, shutting off further disk fragmentation. The existing proto-stars however continue to gain mass by gas accretion. As a results, the initial mass function for disk-born stars at distances $R \\sim 0.03-3$ parsec from the super-massive black hole should be top-heavy. The effect is most pronounced at around $R\\sim 0.1$ parsec. We suggest that this result explains observations of rings of young massive stars in our Galaxy and in M31, and predict that more of such rings will be discovered. ",
        "introduction": "\\label{sec:intro}  Accretion disks around super-massive black holes (SMBHs) have been predicted to be gravitationally unstable at large radii where they become too cool to resist self-gravity and can collapse to form stars or planets \\citep{Paczynski78,Kolykhalov80,Lin87,Collin99,Gammie01,Goodman03}. There is now observational evidence that the two rings of young massive stars of size $\\sim 0.1$ parsec in the centre of our Galaxy were formed in situ \\citep{NS05,Paumard05}, confirming the theoretical predictions. In our neighbouring Andromeda Galaxy (M31), \\cite{Bender05} recently discovered a population of hot blue stars in a disk or ring of similar size, i.e. with radius of $\\sim 0.15$ parsec. The significance of this discovery is that SMBH in M31 is determined to be as massive as $\\mbh \\approx 1.4 \\times 10^8 \\msun$, or about 40 times more massive than the SMBH in the Milky Way. This fact alone rules out (Eliot Quataert, private communication) the other plausible mechanism of forming stellar disks around SMBHs, e.g. the massive cluster migration scenario \\citep[e.g.,][]{Gerhard01}, because the shear presented by the M31 black hole is much stronger than it is at same distance from \\sgra, and its hard to see how a realistic star cluster would be able to survive that \\citep{Gurkan05}.  In this paper we shall attempt to understand what happens with the gaseous accretion disk around a SMBH when the disk crosses the boundary of the marginal stability to self-gravitation \\citep{Toomre64} and forms first stars. We find that in a range of distances from SMBH, interestingly centered at $R\\sim 0.1$ parsec, creation of first low-mass proto-stars should lead to very rapid accretion on these stars. The respective accretion luminosity greatly exceeds the disk radiative cooling, thus heating and puffing the disk up. The new thermal equilibrium reached is that of a disk stable to self-gravity where further disk {\\em fragmentation} is shut off. Star formation is however continued via accretion onto the existing proto-stars, which then grow to large masses. We therefore predict that stellar disks around SMBHs should generically posses top-heavy IMF, as seems to be observed in \\sgra\\ \\citep{NS05,Nayakshin05b}. In the discussion section we note three main differences between star formation process in a ``normal'' galactic environment and that in an accretion disk near a SMBH. ",
        "conclusions": "In this paper we have shown that the birth of even a small number (by mass fraction) of low-mass proto-stars inside a marginally stable accretion disk near a galactic center will unleash a very strong thermal feedback onto the gaseous disk. In a sub-parsec range of radii, the disk will be heated and thickened so that it becomes stable to further fragmentation. The feedback however is not strong enough to unbind the gas from the deep potential well of the SMBH. Therefore, while the feedback stops a further disk fragmentation, accretional growth of stars already present in the disk proceed. Quite generically, this scenario should lead to the average star created in the SMBH accretion disk being ``obese'' compared to its galactic cousins.   The author acknowledges very useful comments on the draft by Yuri Levin, and fruitful discussions with Andrew King, Jim Pringle and Jorge Cuadra."
    },
    "0512/astro-ph0512533_arXiv.txt": {
        "abstract": "{ A catalogue of astrometric (positions, proper motions) and photometric (B, V, R, r$'$, J) data of stars in fields with ICRF objects has been compiled  at the Observatory of the National Academy of Sciences of Ukraine and the Kyiv University Observatory. All fields are located in the declination zone from 0$^{\\circ}$ to +30$^{\\circ}$; the nominal field size is 46$'$(right ascension)x24$'$ (declination). The observational basis of this work is 1100 CCD scans down to V=17 mag which were obtained with the Kyiv meridian axial circle in 2001--2003. The catalogue is presented in two versions. The version KMAC1-T contains 159 fields (104\\,796 stars) and was obtained with reduction to the Tycho2 catalogue. For another 33 fields, due to a low sky density of Tycho2 stars, the reduction was found to be unreliable. Transformation to the ICRF system in the second version of the catalogue (KMAC1-CU) was performed using the UCAC2 and CMC13 catalogues as a reference; it contains 115\\,032 stars in  192 fields and is of slightly  better accuracy.  The external accuracy of one catalogue position is about 50--90~mas for V$<$15~mag stars. The average error of photometry is better than 0.1~mag for stars down to 16~mag. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512019_arXiv.txt": {
        "abstract": "{Recent \\emph{Chandra} observations of the Galactic center region (GCR) have uncovered a population of faint discrete X-ray sources.  A few theoretical works have been made to investigate the nature of these sources.} {We examine the contributions and luminosity functions of various kinds of candidate objects which are proposed either by previous authors or by ourselves.} {We conduct a population synthesis calculation based on Hurley et al.'s rapid binary evolution code.  Several candidate models, i.e. wind-accreting neutron stars, intermediate polars, low mass X-ray binaries, young pulsars and massive stars with strong winds, are incorporated into our calculation.  We also take the geometric effect of the accretion disk into account for Roche lobe overflow X-ray binaries.} {Our results show that neutron star low-mass X-ray binaries contribute significantly to the observed sources.  We also point out that wind-accreting neutron stars contribute negligibly to these sources due to propeller effect, and the intermediate polars play a minor role in accounting for the faint X-ray sources in both Wang et al. and Muno et al. survey. It should be mentioned that the majority of the sources in the survey field of Wang et al. are still beyond our expectation.} {} ",
        "introduction": "Instrumental development always leads people to study sources looking fainter. Recent {\\em Chandra} observations of the Galactic center region (GCR) have uncovered a population of faint discrete X-ray sources.  Wang, Gotthelf, \\& Lang (\\cite{Wang}) have reported hundreds of X-ray sources with luminosities $L_{\\rm X}\\sim 10^{33}- 10^{35}\\,\\ergs$ in the GCR which contains $\\lsim 1\\%$ of the total Galactic population (Pfahl et al. \\cite{Pfahl}, hereafter PRP02). Many of these sources are emitting in a hard energy ($>3$ keV) band since the X-ray absorption, varying across the field, is particularly high in the $1-3$ keV band. Muno et al. (\\cite{Muno03}) presented a catalog of 2357 point sources detected during their 590 ks {\\em Chandra} observations of the $17'\\times 17'$ field around Sgr A$^{\\ast}$. It was estimated that there are $\\sim 2000$ individual point-like sources with luminosities $L_{\\rm X}\\sim 10^{30}-10^{33}\\,\\ergs$ lying within a cylinder of radius 20 pc and  depth 440 pc centered on the Galactic center at a distance of 8.5\\,kpc (Muno et al. \\cite{Muno04a}). The spectra of these sources can be fitted by an absorbed power law with photon index $\\Gamma$. More than 1000 sources have relatively hard spectra with $\\Gamma<1$.  The nature of the faint X-ray sources in the GCR has been investigated by several authors. Pfahl et al. (\\cite{Pfahl}) attributed a significant fraction of the X-ray sources revealed by Wang et al. (\\cite{Wang}) to wind-accreting neutron stars (WNSs). By calculating the flux distribution of WNSs, they concluded that the detected number ranges from ten to several hundred. The similar idea for pre-low-mass X-ray binaries was suggested by Willems \\& Kolb (\\cite{Willems}, see however, Popov \\cite{Popov}). Muno et al. (\\cite{Muno04a}) proposed that magnetized cataclysmic variables (CVs), referred to as polars and intermediate polars (IPs), could account for $\\sim 1000-2000$ X-ray sources in their survey. The population synthesis work by Belczynski \\& Taam (\\cite{Belczynski}, hereafter BT04) suggested that transient neutron star low-mass X-ray binaries (NS LMXBs) may contribute primarily to the X-ray source population in Muno et al. survey.  In this paper, we employed the evolutionary population synthesis (EPS) method to calculate the expected numbers  and luminosity distributions of various types of candidates that shine X-rays in Wang et al. and Muno et al. surveys. We examined the spin evolution of WNSs proposed by PRP02, including the ejector, propeller and accretor stages (Davies \\& Pringle \\cite{Davies}, hereafter DP81; Lipunov \\cite{Lipunov}; Ikhsanov \\cite{Ikhsanov01}, \\cite{Ikhsanov02}). Besides traditional X-ray binaries and CVs, we also considered the contribution from rotation-powered pulsars and massive stars with strong winds. We describe the population synthesis method and the input physics for various types of X-ray sources in our model in \\S 2. The calculated results are presented in \\S 3. Our discussions and conclusions are in \\S 4. ",
        "conclusions": "By use of the EPS method, we have investigated the nature of the faint X-ray sources in the GCR. Our simulated objects include WNSs, magnetic CVs, LMXBs, rotation-powered pulsars and massive stars with strong winds. The main results are summarized in Tables 2 and 3, with $\\sim$10-30 and $\\sim$1000-2000 sources expected to be detected in Wang et al. and Muno et al. surveys, respectively. For Wang et al. survey, we find that a considerable fraction of the discrete sources may be rotation-powered pulsars and NS LMXB transients in quiescence, while WNSs proposed by PRP02 have negligible contribution due to the propeller effect (Note that we have not considered Be/X-ray binaries in which the NS evolution may be quite different from investigated here.). Recent near infrared imaging of the X-ray sources in the GCR shows that the colour distribution of the identified candidate counterparts of some of the X-ray sources is redder than expected for WNS systems, but consistent with later-type stars (Bandyopadhyay et al. \\cite{Bandyopadhyay}). For Muno et al. field, IPs present a minor contribution, and the majority of the X-ray sources seem to be NS LMXB transients with WD donors. The latter result is consistent with BT04.   Our calculations also suggest that some of the point sources detected by Muno et al. survey may be massive stars with strong winds. Radio emission from these winds should be detectable at centimeter wavelengths (Panagia \\& Felli \\cite{Panagia}; Wright \\& Barlow \\cite{Wright}).  We propose that a systematic radio observing campaign be undertaken to search for the stellar counterparts in the GCR surveyed by Muno et al. (some efforts have been made, e.g. by Lang et al. \\cite{Lang} most recently). Moreover, radio observations could reveal the existence of jets, which are common in X-ray binaries.  Sakano et al. (\\cite{Sakano}) recently reported the discovery in the GCR of two unusual X-ray transients XMM J174457$-$2850.3 and XMM J174544$-$2913.0 with flux variations in excess of a factor of 100 during roughly a year and peak X-ray luminosities of $\\sim5\\times10^{34}\\,\\ergs$ and with peculiar spectral features. Since no known classes of sources can well explain all their characteristics, these authors argued that these two sources may represent a new type of sources with different properties from those we have known. This also implies that what is happening in the GCR is much more complicated than we have already learnt."
    },
    "0512/astro-ph0512196_arXiv.txt": {
        "abstract": "The ``canonical behaviour\" of the early X-ray afterglows of long-duration Gamma-Ray Bursts (GRBs) ---observed by the X-Ray Telescope of the SWIFT satellite--- is precisely the one predicted by the Cannonball model of GRBs. ",
        "introduction": "Within a year of its launch on 20 November, 2004, SWIFT accomplished most of its goals (Gehrels et al.~2004): it saw and localized short-duration GRBs and discovered their X-ray afterglow (e.g.~Gehrels et al.~2005; Covino et al.~2005; Retter et al.~2005), heralding the discovery of their host galaxies in ground-based follow-up observations (e.g.~Bloom et al.~2005a, 2005b; Hjorth et al.~2005; Antonelli et al.~2005). SWIFT has detected GRBs at large redshifts (Jakobsson et al.~2005), with a record  $z=6.29$ (Haislip et al.~2005; Cusumano et al.~2005; Tagliaferri et al.~2005a). Its X-ray telescope (XRT) measured well the  X-ray afterglows (AGs) of long-duration GRBs in the first hours after burst (e.g.~Chincarini~2005). The latter observations have been claimed to provide two {\\it major surprises:}  \\noindent (a) The 0.2--10 keV light curves of the X-ray AGs exhibit a ``canonical behaviour\", to wit: (i) an initial very steep decay ($t^{-\\alpha}$ with $3<\\alpha<5$), followed by (ii) a shallow decay ($0.2<\\alpha<0.8$), which finally evolves into (iii) a steeper decay ($1<\\alpha<1.5$). These power-law segments are separated by the corresponding ``ankle'' and ``break'' at times 300 s $<t_{\\rm ankle}<$ 500 s and $10^3$ s $<t_{\\rm break}<10^5$ s (Chincarini et al.~2005; Nousek et al.~2005; Cusumano et al.~2005; Hill et al.~2005; Vaughan et al.~2005; Tagliaferri et al.~2005b; Barthelmy et al.~2005). This is illustrated in Figs.~\\ref{figs1+2}a,b for GRB 050315 and  GRB 050319.  \\noindent (b) While most of the early AG light curves decline smoothly, a substantial fraction has large X-ray flares on short time scales (see, e.g.~Burrows et al.~2005a, 2005b; Nousek et al.~2005).  It has also been claimed (e.g.~Chincarini et al.~2005; Nousek et al.~2005; Zhang~2005) that the ``canonical behaviour'' and the flares cannot easily be explained by current models. This may be true of the popular fireball (FB) models (for a general review, see e.g.~Zhang \\& M\\'esz\\'aros~2004). It is not true of the ``Cannonball\" (CB) model  (Dar \\& De R\\'ujula 2000, 2004 and references therein). In the CB model the observed canonical behaviour and the flares are {\\it predictions} (Dado, Dar \\& De R\\'ujula~2002, hereafter DDD2002). This is discussed in this letter for two recent representative SWIFT-XRT observations: GRB 050315 (Vaughan et al.~2005) and GRB 050319 (Cusumano et al.~2005). ",
        "conclusions": "The early-time X-ray afterglows of GRBs measured with SWIFT-XRT provide an excellent test for two contenders: the FB and CB models. We have shown that the CB model passes this test with flying colours: the observed ``canonical\" behaviour was correctly predicted  (DDD2002) long before the launch of SWIFT. To the best of our knowledge, no satisfactory FB model explanation of this behaviour has been found, though certain correlations have been proposed. Let the early X-ray flux be described as $F_\\nu(t)\\propto\\nu^{-\\beta_x}\\, t^{-\\alpha_1}$. The FB-model prediction is $\\alpha_1\\!-\\!\\beta_x \\!=\\! 2$, as recently discussed by Kumar et al.~(2006). The CB-model prediction is  $\\beta_x\\!\\simeq\\! 1.1$ at all times, so that $\\alpha_1\\!-\\!\\beta_x \\!\\approx\\! \\alpha_1\\!-1.1$. The models are compared with the data of O'Brein et al.~(2006) in Fig.~(\\ref{fig3new}). While the prediction of the CB model agrees with the data, the prediction of the FB model is in clear contradiction with the observations.  A celebrated prediction of  ``conical fireball\" models is a break in the AGs when the beaming angle of radiation from a decelerating cone increases beyond the opening angle of the jet, $\\gamma(t)^{-1}\\sim \\theta_j$, and the observer begins to see the full cone (e.g.~Rhoads~1999; Sari et al.~1999). If the observer happens to be nearly on the cone's axis, the break time (Sari et al.~1999) is: \\begin{equation} t_{\\rm break}\\sim 2.23\\, (1+z)\\left[{\\theta_j\\over 0.1}\\right]^{8/3} \\left[{n\\over 0.1\\, {\\rm cm}^{-3}}\\right]^{-1/3} \\left[{\\eta_\\gamma\\over 0.2}\\right]^{-1/3} \\left[{E_\\gamma^{\\rm iso}\\over 10^{53}\\, {\\rm ergs}}\\right]^{1/3}\\, {\\rm day} \\, , \\label{tbreakFB} \\end{equation} with $\\eta_\\gamma$ the conversion efficiency of the ejecta's energy into $\\gamma$ rays. Frail et al.~(2001) suggested that the $\\gamma$-ray energy of conical GRBs is a $\\theta_j$-independent standard candle. Consequently $E_\\gamma^{\\rm iso}\\approx E_\\gamma\\,\\theta_j^2/4$: the Frail Relation\\footnote{The fraction of the sky illuminated by a conical GRB is $f_b=(1-\\cos\\theta_j)/2\\approx \\theta_j^2/4$ and not $(1-\\cos\\theta_j)\\approx \\theta_j^2/2$, used by Frail et al.~(2001) and many other authors. Bipolar GRBs do light a fraction $\\theta_j^2/2.$}. From 16 GRBs of known $z$, Bloom et al.~(2003) found $E_\\gamma\\approx 1.3\\times 10^{51}$ ergs. Insert this and $\\theta_j^2=4\\,E_\\gamma/E_\\gamma^{\\rm iso}$ into Eq.~(\\ref{tbreakFB}), to obtain: \\begin{equation} t_{\\rm break}\\sim 4.33\\, (1+z) \\left[{n\\over 0.1\\, {\\rm cm}^{-3}}\\right]^{-1/3} \\left[{\\eta_\\gamma\\over 0.2}\\right]^{-1/3} \\left[{E_\\gamma^{\\rm iso}\\over 10^{53}\\, {\\rm ergs}}\\right]^{-1}\\, {\\rm day} \\, . \\label{tbreakFB2} \\end{equation} All published attempts to use Eq.~(\\ref{tbreakFB2}) to predict $t_{\\rm break}$ ---before it was measured--- from the observed $E_\\gamma^{\\rm iso}$, failed\\footnote{For instance, Rhoads et al.~2003 predicted $t_{\\rm break}>10.8$ days for GRB 030226, while Greiner et al.~2003, shortly after, observed $t_{\\rm break}\\sim 0.8$ day.}. In fact, in many GRBs the AG is not achromatic as expected in the jetted FB model, and the break in the X-ray AG is not matched by a similar break in the optical AG (e.g. Panaitescu et al. 2006). This further questions the FB-model interpretation of the AG break (Rhoads~1997, 1999; Sari, Piran and Halpern~1999) and the proclaimed success of the ``Frail Relation''. These suggest that the success of the Frail relation is an artefact resulting from an a-posteriori adjustment of free parameters. Moreover, $E_\\gamma^{\\rm iso}$ for all XRFs with known $z$ is much smaller than the ``standard-candle\" value of Frail et al. 2001, implying that XRFs and GRBs cannot be the same phenomenon viewed from different angles, contrary to indications (e.g.~Dado et al.~2004a). Many other more successful relations, such as the one between the equivalent isotropic energy and the ``peak'' $\\gamma$-ray energy of GRB pulses (the ``Amati Relation'') are predictions of the CB model (Dar \\& De R\\'ujula~2004; Dado \\& Dar 2005) but not of the FB model.  In the CB model, AG flares follow either from CB encounters with density inhomogeneities (DDD2002; Dado et al.~2004b) or from late accretion episodes on the compact central object. In the FB model, late-time flares result from late central activity (e.g.~Granot, Nakar and Piran~2003). Although such an activity can neither be predicted nor ruled out, it is not clear why the ensuing ejecta do not also produce $\\gamma$-ray pulses and why the duration and magnitude of the AG flares scale roughly with the time and magnitude of the declining AG."
    },
    "0512/astro-ph0512475_arXiv.txt": {
        "abstract": "We present the Shapley Optical Survey, a photometric study covering a $\\sim$ 2 deg$^2$ region of the Shapley Supercluster core at z$\\sim$0.05 in two bands (B and R). The galaxy sample is complete to B = 22.5 ($>$M$^*$+6, $\\mathrm{N_{gal}}$ = 16\\,588), and R = 22.0 ($>$M$^*$+7, $\\mathrm{N_{gal}}$ = 28\\,008). The galaxy luminosity function cannot be described by a single Schechter function due to dips apparent at B $\\sim$ 17.5 (M$_\\mathrm{B} \\sim$ - 19.3) and R $\\sim$ 17.0 (M$_\\mathrm{R} \\sim$ - 19.8) and the clear upturn in the counts for galaxies fainter than B and R $\\sim 18$ mag.  We find, instead, that the sum of a Gaussian and a Schechter function, for bright and faint galaxies respectively, is a suitable representation of the data.  We study the effects of the environment on the photometric properties of galaxies, deriving the galaxy luminosity functions in three regions selected according to the local galaxy density, and find a marked luminosity segregation, in the sense that the LF faint-end is different at more than 3$\\sigma$ confidence level in regions with different densities.  In addition, the luminosity functions of red and blue galaxy populations show very different behaviours: while red sequence counts are very similar to those obtained for the global galaxy population, the blue galaxy luminosity functions are well described by a single Schechter function and do not vary with the density.  Such large environmentally-dependent deviations from a single Schechter function are difficult to produce solely within galaxy merging or suffocation scenarios. Instead the data support the idea that mechanisms related to the cluster environment, such as galaxy harassment or ram-pressure stripping, shape the galaxy LFs by terminating star-formation and producing mass loss in galaxies at $\\sim{\\mathrm M}^*+2$, a magnitude range where blue late-type spirals used to dominate cluster populations, but are now absent. ",
        "introduction": "\\label{intro}  The properties and evolution of galaxies are strongly dependent on environment (e.g., Blanton et al. \\citeyear{bla05}; Rines et al. \\citeyear{rin05}; Smith et al.  \\citeyear{smi05}; Tanaka et al. \\citeyear{tan05}). In particular the cluster galaxy population has evolved rapidly over the last 4 Gyr (Butcher \\& Oemler \\citeyear{but84}). While distant clusters are dominated, particularly at faint magnitudes, by blue spiral galaxies, often with signs of disturbed morphologies and evidence of multiple recent star-formation events (Dressler et al. \\citeyear{dre94}), local clusters are completely dominated by passive early-type galaxies.  Recent observational studies on the luminosity, colours, morphology and spectral properties of galaxies have pointed out that the physical mechanisms which produce the transformation in galaxies affecting both the structure and the star formation are naturally driven by and related to the environment (e.g., Treu et al. \\citeyear{tre03}). These processes are linked in various ways to the local density and the properties of the intra-cluster medium (ICM).  In fact, galaxy-ICM interactions, such as ram-pressure stripping and suffocation, require a dense ICM and take place principally in the central cluster regions. The high density regions are also characterized by a steep cluster potential, and we can expect that galaxy-cluster gravitational interactions such as tidal stripping and tidal triggering are also dominant. On the contrary, in low-density environment galaxies have never been through the cluster centre and therefore have never experienced the effects of tidal stripping and tidal triggering of star formation, so the dominant mechanisms are galaxy-galaxy interactions, in terms of both low-speed interactions between galaxies of similar mass (mergers) and high-speed interactions between galaxies in the potential of the cluster (harassment).  The above mentioned mechanisms affect differently the observed morphology and/or the star formation properties of galaxies. In particular, galaxy harassment and ram pressure stripping cause a partial loss of gas mass and, depending on the fraction of gas removed and its rate, ram-pressure stripping can lead either to a rapid quenching of star formation or to a slow decrease in the star formation rate (Larson, Tinsley, \\& Caldwell \\citeyear{lar80}; Balogh et al. \\citeyear{bal00}; Diaferio et al. \\citeyear{dia01}; Drake et al. \\citeyear{dra00}).  Successive high-speed encounters between galaxies (galaxy harassment) lead to gas inflow and strong star formation activity (Fujita \\citeyear{fuj98}).  With the aim to investigate the effects of the environment on the galaxy population, we have undertaken an optical study of the Shapley Supercluster (SSC) core, one of the densest structures in the nearby Universe. The study of this region, selected because of its physical peculiarity in terms of density and complexity, but also for the availability of multiwavelength observations, will take advantage of deep optical photometry from the ESO Archive covering an area of 2 deg$^2$.  The Shapley Optical Survey (SOS) in B and R bands provides a galaxy sample complete and reliable up to 22.5 mag and 22.0 mag in B and R bands, respectively.  We plan to use the excellent SOS dataset to study, with respect to the supercluster environment, the distribution of galaxy populations both in luminosity and colour and the galaxy structural properties comparing the observations with theoretical predictions. In the present paper we will present the dataset, the catalogues and the galaxy luminosity functions in B and R bands as function of the SSC environment.  The galaxy luminosity function (LF), which describes the number of galaxies per unit volume as function of luminosity, is a powerful tool to constrain galaxy transformations, since it is directly related to the galaxy mass function. Moreover, the effect of environment on the observed galaxy LF could provide a powerful discriminator among the proposed mechanisms for the transformations of galaxies. The effects of galaxy merging and suffocation on the cluster galaxy population have been studied through combining high-resolution N-body simulations with semi-analytic models for galaxy evolution (e.g., Springel et al. \\citeyear{spi01}; Kang et al. \\citeyear{kan05}). These show that while galaxy merging is important for producing the most luminous cluster galaxies, the resultant LF can always be well described by a Schechter (Schechter \\citeyear{sch76}) function, although both M$^*$ and the faint-end slope can show mild trends with environment. Galaxy mergers are also inhibited once the relative encounter velocities become much greater than the internal velocity dispersion of galaxies, and so are rare in rich clusters (Ghigna et al. \\citeyear{ghi98}). In contrast, galaxy harassment and ram-pressure stripping may change the LF shape as galaxies lose mass in interactions with other galaxies, the cluster's tidal field, and the ICM. In particular \\citeauthor{moo98} (\\citeyear{moo98}) showed that harassment has virtually no effect on a system as dense as a giant elliptical galaxy or a spiral bulge and only purely disk galaxies can be turned into spheroidals, so these mechanisms produces a cutoff for Sd-Im galaxies. Since the luminosity function is strongly type specific, and those for Sc and Sd/Im galaxies can be described by narrow ($\\sigma\\sim1$ mag) Gaussian distributions centred at $\\sim$ M$^*+$1 and $\\sim$ M$^*+3$ (de Lapparent \\citeyear{del03bis}), the effects of galaxy harassment could be characterized by a dip in the LF at these magnitudes.  In order to further investigate and to assess the relative importance of the processes that may be responsible for the galaxy transformations, we have performed a photometric study of the SSC core, examining in particular the effect of the environment through the comparison of luminosity functions in regions with different local densities.  The SSC was observed by Raychaudhury (\\citeyear{ray89}) and the LF was firstly derived by Metcalfe, Godwin \\& Peach (\\citeyear{met94}; hereafter MGP94). By using photographic data, they investigate a region of 4.69 deg$^2$ around the cluster A\\,3558 considering a sample of 4599 galaxies complete and uncontamiminated by stars (to 2$\\%$ level) for ${\\it b}<19.5$. The derived LF for the central region of 1.35 deg$^2$ showed a broad peak in the number of galaxies at ${\\it b}=18$ which cannot be well fitted by a Schechter function. Moreover, MGP94 found a deficit of blue galaxies in the A\\,3558 core suggesting morphological segregation. However, their study is limited to bright magnitudes, preventing the determination of the faint-end slope while taking advantage of deeper photometry and larger sample of galaxies distributed in larger SSC area, we can provide clear evidence on the LF shape thus quantifying the environmental effect on the LF properties.  The layout of this work is the following. General information of the structure of the SSC core are summarized in Sect.~\\ref{sec:SSC}. We describe observations, data reduction and the photometric calibrations in Sect.~\\ref{sec:2}. The catalogues are presented in Sect.~\\ref{sec:4}. Sect.~\\ref{sec:5} is dedicated to the definition of the environment and in and Sect.~\\ref{sec:6} we show the LFs. Finally Sect.~\\ref{sec:7} contains the summary and the discussion of the results. In this work we assume H$_0$ = 70 \\ks\\,Mpc$^{-1}$, $\\Omega_m$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7. According to this cosmology, 1 arcmin corresponds to 0.060 Mpc at $z=0.048$. ",
        "conclusions": "\\label{sec:7}  We have presented a detailed analysis of the LFs for galaxies in the SSC core. All the luminosity functions were calculated through a weighted parametric fit of a single Schechter function and a composite function, given by the sum of a Gaussian for the bright-end and a Schechter for the faint-end of the LF. The main results of our analysis are the following.  \\begin{description}  \\item[-] The LFs in the whole SOS area have a bimodal behaviour both in B and R band. The weighted parametric fit of a S function is unable to describe the observed LF at faint magnitudes, in particular the dips apparent at B $\\sim$ 17.5 (M$_\\mathrm{B} \\sim$ - 19.3) and R $\\sim$ 17.0 (M$_\\mathrm{R} \\sim$ - 19.8) and the clear upturn for galaxies fainter than 18 mag. To successfully model these dip and changes in slope a composite G+S LF is required.  \\item[-] By deriving the LFs in regions with different local surface densities of R$<21.0$ galaxies we showed that, as observed in the LFs of the whole field, a dip is present at M$_{\\mathrm{R}} \\sim$ -19.8 for LFs in intermediate- and low-density regions, while for the high-density region, the data are well represented by the S function. Moreover the faint-end slope, $\\alpha$, shows a strong dependence on environment, becoming steeper at $>3\\sigma$ significance level from high- ($\\alpha_{\\mathrm{B}}$ = -1.46, $\\alpha_{\\mathrm{R}}$ = -1.30) to low-density environments ($\\alpha_{\\mathrm{B}}$ = -1.66, $\\alpha_{\\mathrm{R}}$ = -1.80) in both bands.  \\item[-] We derived the LFs separately for red and blue galaxy populations according to their B-R colours. The LFs of these two populations show a very different behaviour. In fact differently from the red sequence galaxy counts that are very similar to those obtained with a statistical background subtraction, the blue galaxy LFs are well described by a S function and do not vary with the density.  \\end{description}  These results confirm and extend those of MGP94 who found a peak in the number of galaxies at ${\\it b}=18$ and suggested that the Abell function is a better representation of the integral counts than the S function. However, their optical LF is limited at galaxies three magnitudes brighter than those analysed in the present work, preventing the determination of the steepening of the LF faint-end and a more clear definition of the LF shape.  On the other hand, MGP94 also noted that the CM red sequence galaxies show the broad peak at bright magnitudes in agreement with our findings.  The bimodality of the galaxy LF is commonly observed for rich clusters (e.g., Yagi et al. \\citeyear{yag02}; Mercurio et al. \\citeyear{mer03}), and using data from the RASS-SDSS galaxy cluster survey, Popesso et al. (\\citeyear{pop05}) find a similar variation of the LF with environment to that observed here, but using cluster-centric radius rather than local density (e.g., Haines et al. \\citeyear{hai04}) as a proxy for environment. This observed bimodality and its variation with environment can be best accomodated in a scenario where bright and faint galaxy populations have followed different evolution histories.  The SDSS and 2dFGRS surveys have indicated that the evolution of bright galaxies is strongly dependent on environment as measured by their local density, yet is independent of the richness of the structure to which the galaxy is bound, indicating that mechanisms such as merging or suffocation play a dominant role in transforming galaxies, rather than harassment or ram pressure stripping (G\\'omez et al. \\citeyear{gom03}; Tanaka et al.  \\citeyear{tan04}). However, it is difficult to reconcile the dramatic deviations from the S function observed for intermediate- and low-density regions with the transformation of field galaxies being due to just merging or suffocation, neither of which should alter the shape of LF, whilst they can be explained more easily by a scenario involving mass loss of low-luminosity galaxies.  One such mechanism is galaxy harassment (Moore et al. \\citeyear{moo96}, \\citeyear{moo98}), whereby repeated close ($<$50\\,kpc) high-velocity ($>$1000\\,km\\,s$^{-1}$) encounters with bright galaxies and the cluster's tidal field cause impulsive gravitational shocks that damage the fragile disks of late-type spirals. The cumulative effect of these shocks is the transformation of late-type spirals to spheroidal galaxies over a period of several Gyr.  An important aspect of galaxy harassment is that it has virtually no effect on systems as dense as giant elliptical or spiral bulges, and hence only pure disk systems (Sc or later) are affected. While these galaxies make up the vast majority of the faint (M$>$M$^*$+2) cluster galaxy population at $\\mathrm{z}\\gtrsim0.4$, they become rarer exponentially at brighter magnitudes. The encounters can drive the bulk of the dark matter and 20--75\\% of the stars over the tidal radius of the harassed galaxy, whereas in contrast the bulk of the gas collapses inward, and is consumed in a nuclear starburst. The combined results of these effects is a dimming of the harassed galaxy by $\\sim2$ magnitudes due to mass loss and passive aging of the remaining stars. These remnants are apparent in present day clusters as dwarf spheroids which often show blue cores suggesting nuclear star-formation, as well as remnant disk and bar components (Graham, Jerjen \\& Guzman \\citeyear{gra03}), and signs of rotational support (de Rijcke et al. \\citeyear{der01}).  In agreement with the recent work by Popesso et al. (\\citeyear{pop05}) we suggest that the observed dip at $M_{R}\\sim-19.8$ as well as the strong dependence on environment shown by the faint-end slope in the cluster galaxy luminosity can be explained naturally as the consequence of galaxy harassment.  Alternative mechanisms such as ram-pressure stripping by the ICM or tidal stripping can effect the galaxy population only in the cluster cores, which appears inconsistent with our observation that the dip is greatest in the low-density regions 1--2\\,Mpc from the nearest cluster. However, given the high infall velocities, any galaxy encountering the ICM is likely to be stripped rapidly of their gas, bringing star-formation to a swift halt. Given the high infall velocities, and the typical highly eccentric orbits of cluster galaxies, the low- and intermediate-density regions are likely to contain a significant fraction of galaxies that have already encountered the dense ICM.  In high-density regions, high-velocity dispersions inhibit merging processes (e.g., Mihos \\citeyear{mih04}), hence it is unlikely that dwarf galaxies merge to produce bigger galaxies at the cluster centres. The most likely explanation for the lack of dwarf galaxies near the centre is tidal or collisional disruption of the dwarf galaxies.  This interpretation is also confirmed when analysing separately red sequence galaxies. In fact the red galaxy counts exhibit a behaviour similar to those of the LFs obtained with a statistical background subtraction, confirming the excess of dwarf early type galaxies. Moreover, differently from red sequence galaxies, the blue galaxy LFs are well described by a S function with a slope $\\alpha \\sim$ -1.50 and do not vary with density. This slope is consistent with those recently derived by Blanton et al. (\\citeyear{bla05}) and Madgwick et al. (\\citeyear{mad02}) for field SDSS and 2dF galaxies respectively. This suggest that the observed blue galaxy population is characterized by infalling galaxies that have not yet interacted with the super cluster environment and transformed by the harassment mechanism.  In a forthcoming paper (Haines et al. \\citeyear{hai05}) we will investigate in detail the distribution of red and blue galaxies in the SSC environment."
    },
    "0512/astro-ph0512643_arXiv.txt": {
        "abstract": "We report the discovery of a prominent nonthermal X-ray feature located near the Galactic center that we identify as an energetic pulsar wind nebula. This feature, G359.95-0.04, lies 1 lyr north of \\sgra\\ (in projection), is comet-like in shape, and has a power law spectrum that steepens with increasing distance from the putative pulsar. The distinct spectral and spatial X-ray characteristics of the feature are similar to those belonging the rare class of ram-pressure confined pulsar wind nebulae. The luminosity of the nebula at the distance of \\sgra, consistent with the inferred X-ray absorptions, is $L_x \\sim 1 \\times 10^{34} {\\rm~ergs~s^{-1}}$ in the 2--10 keV energy band.  The cometary tail extends back to a region centered at the massive stellar complex IRS 13 and surrounded by enhanced diffuse X-ray emission, which may represent an associated supernova remnant. Furthermore, the inverse Compton scattering of the strong ambient radiation by the nebula consistently explains the observed TeV emission from the Galactic center. We also briefly discuss plausible connections of G359.95-0.04 to other high-energy sources in the region, such as the young stellar complex IRS 13 and SNR Sgr A East. ",
        "introduction": "The Galactic center (GC) provides a unique laboratory for a detailed study of the interplay between massive star formation and galactic nuclear environment.  Particularly interesting is the presence of a young massive stellar cluster around Sgr A$^*$ --- the super-massive black hole (SMBH) of the Galaxy (e.g., \\citet{gen03}). This cluster, with an integrated mass of $\\sim 10^4 M_\\odot$, has an age of $\\sim 6 \\times 10^6$ yrs and a flat initial mass function; therefore a considerable number of supernovae (SNe) should have occurred in this region. The presence of their stellar end-products (e.g., pulsars) and supernova remnants (SNRs) could be responsible for various high-energy phenomena observed in the GC region and could also have strong impacts on the environment, and hence on accretion onto the SMBH (e.g., \\citet{aha04, bel05}; for a review, see \\citet{meli01}).  However, no pulsar has yet been found within $\\sim 1^\\circ$ of the Galactic center.  Traditional radio searches at $\\lesssim 1$ GHz are insensitive to pulsars in this region because interstellar scattering causes severe pulse broadening \\citep{cor97}.  At high frequencies ($\\nu \\gtrsim 10$ GHz), where pulse broadening ($\\propto \\nu^{-4}$) is not as important, blind pulsar searches are difficult because radio telescope beams are small and because pulsar spectra are typically steep, with correspondingly less flux available. More productive are searches in the hard ($ \\gtrsim 2$ keV) X-ray band, where young pulsars, their wind nebulae (PWNe), and/or SNRs can be identified by X-ray imaging. Indeed, strong X-ray emission from SNR Sgr~A~East, near Sgr A$^*$, has long been known. In particular, \\citep{par05} have suggested that a point-like hard X-ray source CXOGC~J174545.5--285829, or ``cannonball'', is a candidate of a young pulsar, which may have been ejected from the SNR.  Recently, from the larger GC region \\citep{wan02a}, \\citet{wan02} and \\citet{lu03} have detected three X-ray threads associated with nonthermal radio filaments or radio ``wisps''. These highly polarized radio features, observed only in the GC, are due to synchrotron radiation from relativistic particles (electrons and/or positrons). The detection of the X-ray counterparts to these filaments provides strong constraints on the particle acceleration mechanism.  The X-ray emission, if also due to the synchrotron process, must arise from particles accelerated nearly {\\sl in situ}. A possible origin of these particles is PWNe, perhaps shaped by the strong magnetic field and/or ram-pressure of the GC \\citep{wan02}. Alternatively, some of the X-ray threads may represent shocks of young SNRs \\citep{ho85, sak03}.  Here we report on our study of a prominent nonthermal X-ray feature which probably represents the most convincing case for the presence of a young pulsar in the GC. This pulsar is separated from \\sgra\\ by only $\\sim$ 8\\farcs7 at its distance of $8$ kpc, corresponding to a projected separation of 0.32~pc. In the following, all error bars of our X-ray measurements are presented at the 90\\% confidence level. ",
        "conclusions": ""
    },
    "0512/astro-ph0512369_arXiv.txt": {
        "abstract": "We study the possibility of using the entire probability distribution function (PDF) of the aperture mass $\\Map$ and its related cumulative probability distribution function (CPDF) to obtain meaningful constraints on cosmological parameters. Deriving completely analytic expressions for the associated covariance matrices, we construct the Fisher matrix and use it to estimate the accuracy with which various cosmological parameters can be recovered from future surveys using such statistics. This formalism also includes the effect of various noises such as intrinsic ellipticity distribution of galaxies and finite survey size. The estimation errors are then compared with the ones derived from low order moments of the PDF (variance and skewness) to check how efficiently the high $\\Map$ tail can be used to constrain cosmological parameters such as $\\Om$, $\\sigma_8$ and dark energy equation of state $\\wde$. We find that for future surveys such as JDEM the full PDF does not bring significant tightening of constraints on cosmology beyond what is already achievable by the joint use of second and third order moments. ",
        "introduction": "Weak lensing surveys can be used as a very efficient probe to constrain the background cosmology as well as contents of the Universe (see Munshi \\& Valageas 2005 and references therein). Typically low order quantities such as two- and three-point functions or their reduced forms, variance and skewness, or the entire PDF are used to obtain cosmological constraints from observational data (Bernardeau et al. 2004). In weak lensing studies, observables such as $\\kappa$, $\\gamma$ or more commonly $\\Map$ and their low order moments are extensively studied (e.g. Jarvis et al. (2004) for recent results for $\\Map$). Use of $\\Map$ has the additional advantage of being able to separate gravity induced shear signals (``Electric'' modes) from the (``Magnetic'') modes due to various systematics such as point spread function. The one-point probability distribution function of $\\Map$ encodes information regarding non-Gaussianities at all orders (Bernardeau \\& Valageas 2002) and thus can be useful to pinpoint background cosmological parameters by breaking degeneracies which appear at the level of two-point correlation functions (Schneider et al. 2002). We develop an analytical formalism to study the covariance of binned PDF and employ Fisher formalism techniques to study covariance of estimation error of cosmological parameters from realistic weak lensing surveys. The purpose of this study is twofold. Firstly we check the estimation error associated with cosmological parameters while using binned PDF as the primary observable. Secondly we use these results to infer how much tightening of constraints, if any, can be achieved in general by using higher order informations regarding non-Gaussianities. This has been a topic of discussion in the recent years using the lowest order non-Gaussian statistics, e.g. bi-spectrum or related collapsed one-point skewness (Takada \\& Jain 2002, Kilbinger \\& Schneider 2004, Munshi \\& Valageas 2005). Thus our results based on PDF extend such studies to include all-orders of non-Gaussianity. This {\\it letter} is organised as follows: in section~2 we discuss the analytical results concerning covariance structure of the PDF and CPDF. Borrowing results from Valageas et al. (2004) we show how the Fisher matrix can be constructed from PDF and CPDF data. In section~3 the numerical results are provided. Section~4 is left for discussion of our results. ",
        "conclusions": "\\label{Discussion}    \\begin{figure} \\begin{center} \\epsfxsize=7. cm \\epsfysize=5. cm {\\epsfbox[129 60 432 242]{f_fisher.eps}} \\end{center} \\caption{The estimation error $\\Delta\\Theta_{\\alpha} $ for two cosmological parameter pairs are shown for the estimators $F_j$, $\\{\\Om,\\wde\\}$ in right panel and $\\{\\Om,\\sigma_8\\}$ in the left panel. A group of six neighbouring bins is used for each of these points to reduce the scatter in $\\Delta\\Theta_{\\alpha}$.} \\label{fig:fishpdf_err} \\end{figure}   We have checked the possibility of using the entire $\\Map$ PDF to constrain cosmological parameters as opposed to the use of its lower order moments which is more prevalent in the literature. We find that for a SNAP like survey the PDF does not yield significantly tighter constrainsts than those derived from the variance and skewness alone, despite the low-noise space based survey strategy. This also implies that higher order moments would not contribute either to further tightening of error ellipses. However, much larger surveys with high level of source galaxy distributions might still benefit from measurement of kurtosis, but higher number density requires increasing the depth of the survey which in turn makes the PDF itself more Gaussian. Our results only take into account the volume averages of higher order correlation functions and did not propose to quantify non-Gaussianity beyond collapsed one point objects. It is possible to extend our study by taking into account redshift binning or tomography which we will report elsewhere, as well as combining several angular scales. However, as seen in Munshi \\& Valageas (2005) for a space based survey such as SNAP with reasonably good sky coverage and high number density of source galaxies most of the useful information can in effect be obtained by studying one particular angular scale (e.g. $\\theta_s = 2'$ in case of SNAP, compare Fig.~\\ref{fig:fishmoments} with Figs.14, 15 in Munshi \\& Valageas 2005). Indeed, nearby scales are highly correlated and do not provide additional information whereas very large and very small angular scales are more affected by noise (due to the intrinsic ellipticity distribution of galaxies or the finite size of the survey).  The results provided here are for a monolithic survey strategy but real surveys will have more complicated topology. However our results can provide clues as to what extent the surveys can probe non-Gaussianity to extract meaningful constraints on cosmological parameters by using not only the first few lower order moments but the entire PDF.  We have not considered a Wiener filter based approach to reconstruct the PDF from noisy data as suggested by Zhang \\& Pen (2005) for convergence $\\kappa$ maps but such an approach can very easily be implemented by using the covariance matrices we have constructed. A complete Wiener filter based approach using compensated filter $\\Map$ as presented here will be investigated elsewhere. In a different context shape of the non-Gaussian PDF was used by  Huffenberger \\& Seljak (2005) to separate the kinetic-SZ contributions from primordial CMB. Similar approach can be useful in separating gravity generated $E$ mode using $M_{ap}$ PDF from various systematics.Techniques presented here can also be applied to the case of weak lensing of diffuse background such as CMB (Kesden et al. 2002) and high-redshift 21 centimeter radiation from neutral hydrogen during the era of reionization (e.g. Cooray 2004). Results of such analysis will be reported elsewhere."
    },
    "0512/hep-th0512274_arXiv.txt": {
        "abstract": "In theories with a hidden ghost sector that couples to visible matter through gravity only, empty space can decay into ghosts and ordinary matter by graviton exchange. Perturbatively, such processes can be very slow provided that the gravity sector violates Lorentz invariance above some cut-off scale. Here, we investigate non-perturbative decay processes involving ghosts, such as the spontaneous creation of self-gravitating lumps of ghost matter, as well as pairs of Bondi dipoles (\\ie lumps of ghost matter chasing after positive energy objects). We find the corresponding instantons and calculate their Euclidean action. In some cases, the instantons induce topology change or have negative Euclidean action. To shed some light on the meaning of such peculiarities, we also consider the nucleation of concentrical domain walls of ordinary and ghost matter, where the Euclidean calculation can be compared with the canonical (Lorentzian) description of tunneling. We conclude that non-perturbative ghost nucleation processes can be safely suppressed in phenomenological scenarios. ",
        "introduction": "Recently it has been suggested that the smallness of the observed cosmological constant can be attributed to an approximate ``energy symmetry\" \\cite{kasu}. The idea is that Nature is endowed with an exact copy of the matter sector, but with an overall minus sign in the action \\cite{li,kasu}, \\begin{equation} {\\cal L} = \\sqrt{-g}\\{ M_P^2 R + {\\cal L}_{matt}(\\psi,g) - {\\cal L}_{matt}(\\hat\\psi, g)+...\\}. \\label{esym} \\end{equation} Here, $g$ is the metric, $\\psi$ are ordinary matter fields (including those of the Standard Model), and $\\hat \\psi$ are the ghost fields. Energy parity is defined by $\\psi \\to \\hat \\psi, \\quad \\hat\\psi\\to \\psi, \\quad g\\to g.$ Ignoring gravity, the Hamiltonian $H$ transforms as \\beq H\\to -H\\,. \\label{htomh} \\eeq The vacuum state $|0\\rangle$ is defined as parity invariant, and from (\\ref{htomh}) the corresponding vacuum energy vanishes to all orders in perturbation theory. However, gravity breaks the energy symmetry and a cosmological constant is induced. It was argued in \\cite{kasu} that the magnitude of this vacuum energy can be comparable to the one suggested by observations provided that the gravitational cut-off scale $\\mu$ is sufficiently low (lower than the inverse of 30 microns, about an order of magnitude beyond the reaches of ongoing short distance probes of gravity).  ``Phantom\" matter has also been invoked in phenomenological studies of dark energy \\cite{de,ca,cline}, as a way of obtaining an effective equation of state parameter $w < - 1$. This violates all the standard energy conditions, but is not at all disfavoured by observations. A ghost sector is also present in the recently proposed B-inflation, based on effective theories with only second derivatives of a scalar field \\cite{an}.  In all of these cases, the Hamiltonian is unbounded below, and disaster would follow unless one postulates that the Lorentz symmetry is broken at a certain energy scale \\cite{cline}. The reason is simple. In the model (\\ref{esym}), empty space can decay into a pair of $\\psi$ ordinary particles and a pair of $\\hat \\psi$ ghost particles \\begin{equation}\\label{decay1} |0 \\rangle \\to \\psi\\psi\\hat\\psi\\hat\\psi \\label{1} \\end{equation} (if particles are charged, one in each pair should be understood as the antiparticle). Let the momenta of the ordinary particles be $p_1$ and $p_2$, and the momenta of the ghost particles $k_1$ and $k_2$. From translation invariance, the decay amplitude takes the form $\\langle p_1,p_2,k_1,k_2|0\\rangle={\\cal A}(p_1,p_2,k_1,k_2)\\ \\delta^{(4)}(p_1+p_2+k_1+k_2)$, which after integration over external momenta leads to the vacuum decay rate per unit volume \\begin{equation} \\Gamma = \\int d^4 P\\ \\gamma(P), \\label{3} \\end{equation} where $\\gamma(P)= \\int d {\\tilde p_1}\\ d{\\tilde p_2}\\ d {\\tilde k_1}\\ d{\\tilde k_2}\\ |{\\cal A}|^2 \\ \\delta^{(4)}(P+k_1+k_2)\\ \\delta^{(4)}(P-p_1-p_2)$. In a Lorentz invariant theory, $\\gamma(P)$ is just a function of $s=-P_\\mu P^\\mu$, and Defining $\\vec v = s^{-1/2} \\vec P$, we have \\beq \\Gamma = \\int ds\\ s\\ \\gamma(s) \\int {d^3 \\vec v\\over 2 \\sqrt{1+\\vec v^2}}. \\eeq Physically, the last integral corresponds to the fact that there is no preferred reference frame, and the total momentum $P^{\\mu}$ of the pair of particles (or the pair of ghosts) is equally likely to fall anywhere on the mass-shell of radius $s^{1/2}$. Particles only interact with ghosts gravitationally, and so the momentum $P^{\\mu}$ is transferred by gravitons. The decay rate is in principle infinite (due to the mass-shell integral) but it can be rendered finite if we postulate that Lorentz invariance is broken in the gravitational sector at some scale ${\\cal E}$ \\cite{cline}. The remaining integral over $s$ can be finite in a theory where gravity becomes soft at a certain cut-off scale $\\mu$, as it is in fact assumed. The process becomes completely negligible if ${\\cal E}$ is comparable to the cut-off scale $\\mu\\lsim(30\\, \\mu\\mathrm{m})^{-1}$ discussed above \\cite{kasu}. Similarly, empty space can decay to ghosts $\\hat\\psi$ and gravitons $h$,\\footnote{Ref.~\\cite{kasu} actually considered the decay $|0 \\rangle \\to h^*\\hat\\psi\\hat\\psi$, where $h^*$ is an ``excited\" (soft-scale) graviton, which is a more dominant process than \\reef{decay2}. We shall not consider the non-perturbative analogue of this process, since it cannot be described in terms of the low energy effective action (\\ref{esym}).} \\begin{equation}\\label{decay2} |0 \\rangle \\to hh\\hat\\psi\\hat\\psi. \\label{2} \\end{equation} In this case, the integrals over the momenta of the external gravitons must be cut-off at the Lorentz violating energy scale ${\\cal E}$. In this way, the vacuum can be made sufficiently stable to perturbative decay processes, in spite of the ghosts \\cite{kasu,ca}.   Although perturbative processes may be suppressed by the Lorentz-violating physics, it is conceivable that non-perturbative processes may quickly destabilize the present vacuum, through the production of lumps of non-relativistic ghost matter. The purpose of the present paper is to investigate the non-perturbative analogues of (\\ref{1}) and (\\ref{2}). Decays that proceed via non-perturbative tunneling are typically slower than their perturbative counterparts, but when ghosts are involved there are several reasons why this is not so obvious.  In accordance with the equivalence principle, a lump of ghost matter tends to fall towards the potential well created by a positive energy object. On the other hand, the repulsive gravitational field it produces tends to push the positive energy object away. It has been known for some time that this leads to a runaway behaviour, where the positive energy object is chased after by the ghost \\cite{bondi,iskh,bonnor,gibbons}, with a constant acceleration. Such configuration is known as a {\\em Bondi dipole} \\cite{bondi}. As we shall see, such self-accelerating solutions can be continued to the Euclidean section, leading to a semiclassical description of the spontaneous nucleation of pairs of Bondi dipoles. This would be the non-perturbative analogue of (\\ref{1}). A simple estimate (which we will confirm by rigorous calculation) gives the Euclidean action of this process as $\\sim m_+ d$, where $d$ is the size of the dipole and $m_+$ is the mass of its ordinary positive-mass component. Even if we impose $d\\gsim\\mu^{-1}$, we see that the action can dangerously approach a value of order one if the Compton wavelength of the particles produced is also close to the gravitational cutoff.  The analogue of (\\ref{2}) is the pair creation of self-gravitating lumps of ghost matter, which repel each other. The possibility of this process is suggested by the following weak field argument. The interaction energy of two ghost particles at rest, with identical mass $m<0$, is given by $E_{grav}=G m^2/ 2 r$. Here $G$ is Newton's constant and $2 r$ is the distance between the masses. For $r= -G m/4$ the positive gravitational energy is equal to minus the rest mass energy of the pair $E_{grav}= -2 m$, so this configuration can in principle pop out of the vacuum without violating energy conservation. Also, the initial acceleration of each particle is given by $a= 1/r$, suggesting that there is a Euclidean solution where the ghost matter runs around a circle of radius $1/a$. Note, however, that for $r \\sim -G m$ the gravitational field is of order one, and non-linear gravity must be taken into account. The corresponding instantons still exist, and have interesting peculiarities which make their interpretation non-trivial. First of all, they can produce a topology change, and second, the corresponding Euclidean action (defined as the bounce action minus the background action) can be {\\it negative}.  To shed some light into the meaning of such peculiarities, we shall first consider the simpler example of vacuum decay in a theory where the matter and the ghost sectors support domain wall solutions. In this case, the process of spontaneous nucleation of concentrical spherical domain walls of ordinary and ghost matter chasing after each other is the analogue of (\\ref{1}). The analogue of (\\ref{2}) is the spontaneous creation of spherical domain walls of ghost matter. If the cosmological constant is exactly vanishing, the instanton for the latter process changes the topology of space and a new boundary appears inside of the domain wall. This makes the calculation of the corresponding Euclidean action somewhat ill-defined. On the other hand, in the presence of a small cosmological constant (such as the one present in our universe), the topology does not change at all and the Euclidean action can be calculated unambiguously. Interestingly, it turns out to be negative. These features are quite similar to what happens in the case of pair creation of lumps of ghost matter, but the advantage here is that the geometry is much simpler, and the Euclidean calculation can be compared with a canonical (Lorentzian) description of tunneling.  \\bigskip  \\noindent The plan of the paper is the following. Since the subject of tunneling in theories with ghosts is fraught with many subtleties, we have developed it in quite some detail in Sections \\ref{sec:ghostwalls} through \\ref{sec:ghostpairs}. The specific application of our results to the energy-symmetric scenario of \\cite{kasu} is then discussed in the concluding Section \\ref{sec:disc}, to which phenomenologically-minded readers might want to jump directly if not interested in the theory of ghost tunneling.  The more technical sections \\ref{sec:ghostwalls}--\\ref{sec:ghostpairs} consist of a discussion of: the spontaneous nucleation of domain walls (Section \\ref{sec:ghostwalls}); the pair creation of Bondi dipoles, \\ie the non-perturbative analogue of \\reef{decay1} (Section \\ref{sec:bondidip}, where also we briefly review several technical aspects of the axisymmetric class of solutions for the benefit of readers unfamiliar with them); the spontaneous creation of self-gravitating lumps of ghost matter, \\ie the non-perturbative analogue of \\reef{decay2} (Section \\ref{sec:ghostpairs}). Some details of the canonical WKB construction of tunneling paths are deferred to an Appendix. ",
        "conclusions": "\\label{sec:disc}  We have studied theories with ghosts (``the Others\") that interact very weakly, through gravity, with the ordinary world. Empty space can decay into ghosts and ordinary particles, and we have focussed on the possibility that non-relativistic lumps of ghost matter be produced by tunneling processes, involving only curvature scales smaller than $\\mu$.  In ordinary field theory, nucleation rates are usually estimated by using instanton methods, as $\\Gamma\\sim e^{-I_E}$, where $I_E$ is the corresponding action. We have computed $I_E$ for a variety of such possible decays. In some cases (eqs.~\\reef{diwact}, \\reef{donkeyact}, \\reef{donkeydsact}), $I_E$ has turned out to be negative. Naively, this would appear to lead to an exponentially enhanced catastrophic decay rate. However, the Euclidean path integral (even in the absence of gravity) in theories with ghosts is ill-defined, and the standard Euclidean rules do not apply. In section \\ref{sec:ghostwalls} we have discussed nucleation processes using the canonical WKB approach, without reference to Euclidean methods. We have argued that for systems of a single degree of freedom, and for tightly bound systems of ghost and ordinary matter, the nucleation rates should be estimated as \\beq \\Gamma \\sim e^{-|I_E|}. \\label{tura2} \\eeq This formula may not be valid if, besides the gravitational interaction between ghosts and matter, there are other dominant forces contributing to the tunneling (such as the expansion of the universe). Here, we are mainly interested in the case when ghosts and matter nucleate from nearly flat space, and the tunneling is possible due to the gravitational energy-momentum transfer between matter and ghosts. In this situation, we expect (\\ref{tura2}) to be valid. In more general cases, we expect that Eq.~(\\ref{tura2}) represents only a conservative upper bound to the nucleation rate. Investigation of this more general case is left for further research.   A first class of decays involves the nucleation in flat space of lumps of ordinary matter together with their ghost counterparts. For the spontaneous nucleation of diwalls (an asymptotically flat configuration of concentrical walls, with an ordinary wall chasing after a ghost wall), the corresponding Euclidean action is given by Eq.~(\\ref{diwact}). If the tensions of the two walls are very different in magnitude, the action is of the form $|I_E| \\sim M_P^6/\\sigma_1^2$, where $\\sigma_1$ is the tension of the ordinary wall. If both tensions are very similar, then the action is parametrically given by \\beq\\label{diwrate} |I_E| \\sim {M_P^6 \\over \\sigma^2}{\\delta\\sigma\\over\\sigma}\\,. \\eeq The curvature radius of the walls is given by $R\\sim 1/ G\\sigma$, so the above formula takes the form $|I_E| \\sim M_P^2 R \\delta R$. The effective theory (\\ref{esym}) is supposed valid only at distance scales larger than the cutoff $\\mu^{-1}$, so within the regime of validity of this effective theory, \\beq\\label{extrasupp} |I_E|\\gsim \\left(\\frac{M_P}{\\mu}\\right)^2\\simeq 10^{60}\\, \\eeq so the process of diwall nucleation is extraordinarily suppressed, $\\Gamma \\sim e^{-10^{60}}$.  Such suppression does not appear to be at work in the nucleation of a pair of Bondi dipoles (the non-perturbative analogue of \\reef{decay1}), for which we have computed the action to be \\beq\\label{diprate} I_E\\simeq 2\\pi m_+ d\\,, \\eeq where $m_+$ and $d$ are the mass and size of the dipole. There is no obvious Planck-scale suppression here. However, since we require $d\\gtrsim \\mu^{-1}$, the process will be suppressed for all lumps whose mass $m$ is much larger than $\\mu$. In the context of the Standard Model with parity symmetry, the only elementary particles which might have a mass comparable to the gravitational cut-off scale are neutrinos. Could we then have non-perturbative nucleation of neutrinos and ghost-neutrinos? The existence of the corresponding instanton requires that the ordinary particle have larger mass than its ghost counterpart, eq.~\\reef{massdiff}. Now, the symmetry between the normal and ghost sectors is broken only by their different coupling to gravity. So any differences between the masses of a particle and its ghost must be due to self-energy corrections from the coupling to gravity. Since normal particles are gravitationally attractive, whereas ghosts are repulsive, the gravitationally-induced corrections to the mass will be \\beq (\\Delta m_+^2)_\\mathrm{grav}<0\\,,\\qquad (\\Delta m_-^2)_\\mathrm{grav}>0\\,. \\eeq These work in a direction opposite to \\reef{massdiff}, and therefore the dangerous decay appears to be not simply suppressed but actually forbidden!  Lastly, we have analyzed the non-perturbative analogue of \\reef{decay2}. This involves a dramatic topology change mediated by an instanton with negative Euclidean action \\beq\\label{ghpair} I_E\\simeq -2\\pi G m_-^2\\,. \\eeq As we have seen, the negative sign does not spell doom since we should use \\reef{tura2} for the decay rate. If we demand that the size of the nucleated configuration $\\sim A^{- 1}\\sim Gm_-$ be larger than the cutoff scale $\\mu^{-1}$, then this process is as strongly suppressed as \\reef{extrasupp}, and therefore imposes no phenomenological constraints.  The decays of flat space involving ghost matter that we have studied do not, by any means, exhaust all possibilities. We have considered ghost domain walls and ghost particles, but ghost cosmic strings can also unstabilize the vacuum. A ghost cosmic string coupled to gravity creates a conical excess. If the string is open with ghost monopoles at its endpoints (\\ie the string is topologically unstable) then the string will pull the monopoles together: such strings can not pop out of the vacuum spontaneously. However, string vortices, even topologically stable ones, can end on black holes \\cite{agk}. If a ghost string has black holes (which necessarily have positive mass) at its endpoints, then the string will push the black holes apart. One can then envisage a process in which one such ghost string with two black holes attached is nucleated out of the flat vacuum. Following \\cite{strbh}, the relevant instanton can be easily constructed, and the action is \\beqa\\label{ghstring} I_E\\simeq \\frac{\\pi m^2}{T}-\\pi Gm^2 =\\frac{\\pi m^2}{T}\\left(1-\\frac{T}{M_P^2}\\right)\\,, \\eeqa where $m$ is the mass of the black holes and $T$ the absolute value of the ghost string tension. The first contribution in \\reef{ghstring} is the nucleation rate of the string itself, and the second is an enhancement of the decay due to black hole entropy\\footnote{In order to construct the instanton the black holes must charged and extremal (or almost extremal) \\cite{strbh}, so their entropy is $S_{bh}\\simeq \\pi Gm^2$.}. The latter becomes subleading for string tensions $T\\ll M_P^2$, so the factor in brackets in \\reef{ghstring} can be taken to be of order one. In order that the size of the nucleated configuration be larger than the gravitational cutoff we must require $m/T\\gsim\\mu^{-1}$. Since the black holes must also be larger than Planck size, $m\\gsim M_P$, we find that \\beq I_E \\gsim \\frac{M_P}{\\mu}\\sim 10^{30}\\,. \\eeq Again, the nucleation rate is strongly suppressed.  The results presented in this paper are quite encouraging for the energy-parity-symmetric theory in (\\ref{esym}). In all the cases where we have found the possibility to conjure ghosts, we have seen they can be exorcized away for all of eternity (for practical purposes), in the sense that $|I_E|$ is always extremely large for the non-perturbative decay processes considered. Nevertheless, it should be kept in mind that the Euclidean path integral is ill-defined in theories with a ghost sector, and our justification of Eq.~(\\ref{tura2}) for the nucleation rates is rather heuristic. In this sense, our analysis should be regarded as preliminary. A more rigorous derivation of nucleation rates in theories with ghosts deserves further investigation. One might also worry that the suppression of the decay seems to be of a different sort in each of the cases we have studied: while \\reef{diwrate}, \\reef{ghpair}, \\reef{ghstring} are suppressed by the weakness of gravity, the Bondi pair prodution \\reef{diprate} is instead eliminated by dynamical considerations. There does not seem to be an underlying generic argument that guarantees that ghost decays are always phenomenologically harmless. Here we have considered what appear to be the simplest and most natural decay channels, but we cannot rule out the existence of a different and dangerous non-perturbative instability."
    },
    "0512/astro-ph0512657_arXiv.txt": {
        "abstract": "% The Energetic X-ray Imaging Survey Telescope (EXIST) is under study for the proposed Black Hole Finder Probe, one of the three {\\it Einstein Probe} missions in NASA's proposed Beyond Einstein Program. EXIST would have unique capabilities: it would survey the full sky at 5-600 keV each 95min orbit with 0.9-5 arcmin, 10$\\mu$sec - 45min and \\about0.5-5 keV resolution to locate sources to 10\\arcsec and enable black holes to be surveyed and studied on all scales. With 1y/5$\\sigma$ survey sensitivity \\Fx(40-80 keV) \\about5 \\X 10$^{-13}$ \\fcgs, or comparable to the  ROSAT soft X-ray (0.3-2.5 keV) sky survey, a large sample (\\gsim2-4 \\X 10$^4$) of obscured AGN will be identified and a complete sample of accreting stellar mass BHs in the Galaxy will be found. The all-sky/all-time coverage will allow rare events to be measured, such as possible stellar disruption flares from dormant AGN out to \\about100 Mpc. A large sample (\\about2-3/day) of GRBs will be located (\\lsim10\\arcsec) at sensitivities and bandwidths much greater than previously and likely yield the highest redshift events and constraints on Pop III BHs. An outline of the mission design from the ongoing concept study is presented. ",
        "introduction": "Of the wide-field surveys considered at this meeting, one is likely widest of all: temporally-resolved imaging in hard X-rays (\\about5-600 keV) of the full sky every 95min. This is possible with a satellite-borne coded aperture imaging telescope array with very wide ``fully-coded'' field of view that scans with a photon-counting detector array over the full sky each orbit. Such is the concept for the Energetic X-ray Imaging Survey Telescope (EXIST), originally recommended by the Decadal Survey as a possible ISS mission and then considered as a Free Flyer (Grindlay et al 2003), and now under study to be the Black Hole Finder Probe in NASA's {\\it Beyond Einstein} Program. The unique all-sky/all-time imaging capability of EXIST opens the temporal survey window and is highly complementary to LSST (or any Wide-Field OIR surveys) and  LISA. The hard X-ray band is particularly well matched for detection and study of accreting black holes, from stellar mass (\\about10\\Msun) in X-ray binaries to supermassive (\\about10$^{7-9}$ \\Msun in galactic nuclei, which characteristically emit their peak luminosity in a broad band around \\about100 keV. At energies above \\about7 keV, photoelectric absorption by column densities \\nh \\gsim10$^{23-24}$ \\cmsq in galactic nuclei or X-ray binaries in the Galaxy becomes negligible and black holes can be seen until \\nh \\gsim 10$^{25}$ \\cmsq and they are obscured by Compton scattering.  \\vspace*{-0.5cm} ",
        "conclusions": ""
    },
    "0512/astro-ph0512461_arXiv.txt": {
        "abstract": "We examine whether active galaxies obey the same relation between black hole mass and stellar velocity dispersion as inactive systems, using the largest published sample of velocity dispersions for active nuclei to date.  The combination of 56 original measurements with objects from the literature not only increases the sample from the 15 considered previously to 88 objects, but allows us to cover an unprecedented range in both stellar velocity dispersion (30--268 \\kms) and black hole mass ($10^5-10^{8.6}$~\\msun).  In the \\msigma\\ relation of active galaxies we find a lower zeropoint than the best-fit relation of Tremaine et al. (2002) for inactive galaxies, and an intrinsic scatter of 0.4 dex.  There is also evidence for a flatter slope at low black hole masses.  We discuss potential contributors to the observed offsets, including variations in the geometry of the broad-line region, evolution in the \\msigma\\ relation, and differential growth between black holes and galaxy bulges. ",
        "introduction": "Black holes (BHs) are a basic component of galaxies, and the existence of a tight correlation between the stellar velocity dispersion of bulges (\\sigmastar) and the BH mass (the \\msigma\\ relation; Gebhardt \\etal\\ 2000a; Ferrarese \\& Merritt 2000) suggests that the growth of the BH plays a fundamental role in the growth of the bulge, although exactly how remains unclear (e.g.,~Silk \\& Rees 1998; Kauffmann \\& Haehnelt 2000; Di~Matteo \\etal\\ 2005). The \\msigma\\ relation for inactive galaxies, in so far as it represents the final state of the BH-bulge system, represents a boundary condition for various evolutionary scenarios, and some clues are embedded in the scatter and possible skewness in the local \\msigma\\ relation (Robertson \\etal\\ 2005).  Unfortunately, the number of available points in the relation for inactive galaxies is limited, and statistics are poor.  Currently, our only recourse is to rely on BH masses from active galactic nuclei (AGNs).  To the extent that the BHs in AGNs continue to gain mass, the relation of BH masses in AGNs to their host bulges may carry additional information about the establishment of the relation for inactive sources.  Since the majority of BH mass was assembled at high redshift (e.g.,~Yu \\& Tremaine 2002), we might expect to find the strongest evidence for evolution in the \\msigma\\ relation at large cosmological distances (Shields \\etal\\ 2003; Treu \\etal\\ 2004; Walter \\etal\\ 2004; Peng \\etal\\ 2005). While undoubtedly this is a vital direction to pursue, there are a number of compelling reasons to study the local AGN \\msigma\\ relation as well.  For one thing, while various methods are available to characterize the bulge potential, virial mass estimation (e.g.,~Ho 1999; Wandel \\etal\\ 1999; Kaspi \\etal\\ 2000), where the dense broad-line region (BLR) gas is assumed to be on Keplerian orbits around the central BH, is currently the most widely utilized tracer of BH mass. In the absence of a detailed model of the BLR, the zeropoint for the virial BH mass scale is set through a direct comparison with stellar velocity dispersions for a small sample of local AGNs (Gebhardt \\etal\\ 2000b; Ferrarese \\etal\\ 2001; Nelson \\etal\\ 2004; Onken \\etal\\ 2004).  The virial BH masses considered in these studies are remarkably consistent with the \\msigma\\ relation of inactive galaxies, suggesting virial masses are reliable.  But we must be cautious.  For instance, there are compelling reasons to believe the BLR is actually a disk-wind (e.g.,~Murray \\& Chiang 1997; Proga et al. 2000; Proga \\& Kallman 2004) whose kinematics depend on both the mass and the accretion rate onto the BH.  In this scenario, the virial mass calibration would depend systematically on BH mass and luminosity.  We require objects spanning a wide range of AGN properties to properly calibrate the primary rung in our AGN BH ``mass ladder.''  At the same time, we may hope to learn about evolution of the \\msigma\\ relation by looking at the full distribution of local BHs (Robertson \\etal\\ 2005), and AGNs in particular (e.g.,~Yu \\& Lu 2004), in the \\msigma\\ plane. ",
        "conclusions": "We have established, with statistical confidence for the first time, that the \\msigma\\ relation of local AGNs, while generally similar to that of inactive galaxies, shows some significant differences.  We find evidence for a lower zeropoint, shallower slope, and (probably) larger scatter.  There are many competing effects, relating both to BLR physics and galaxy evolution,  \\hskip 10mm \\psfig{file=table2_v2.epsi,width=6cm,angle=0} \\vskip +3mm \\noindent that may contribute to the observed differences.  If we posit that AGNs obey the relation of inactive galaxies exactly, then the scatter and zeropoint offset may be attributed to variations in the geometry of the BLR. This is the assumption made by Onken \\etal\\ (2004), who derive a statistical offset in the factor $f$ that scales all BH masses upwards relative to the case of a spherical BLR. If, however, as suggested by disk-wind models (e.g.,~Proga \\& Kallman 2004), the geometry of the BLR depends on physical properties of the AGN such as the BH mass and the Eddington ratio, then virial mass estimates will not scatter randomly about one fixed value of $f$.  Rather, $f$ will change systematically with the state of the system, and its average value for any given sample will depend on the range of parameter space spanned by the objects used to derive it.  It is important to recognize that all work so far---including ours---has considered a rather limited range in \\mbh\\ and \\lledd, and so may lead to a biased value of $f$. Apart from the Greene \\& Ho (2004) objects studied by Barth \\etal\\ (2005), the practical challenge of detecting stellar absorption features to measure \\sigmastar\\ inevitably biases the final sample toward relatively low Eddington ratios and BH masses.  Excluding the Barth et al. objects, the sample summarized in Figure 1 has a median \\lledd\\ = $6\\times10^{-2}$ and $10^6$ \\lax\\ \\mbh\\ \\lax\\ $10^8$~\\msun.  While we expect that secondary parameters (e.g., \\lledd) may ultimately help to account for the overall scatter in the AGN \\msigma\\ relation, we refrain from discussing this issue here because we believe that a proper analysis would require a larger and more complete sample than is currently available.  Apart from systematic uncertainties in virial masses, evolution of the \\msigma\\ relation with cosmic time (e.g.,~Shields \\etal\\ 2003; Treu \\etal\\ 2004; Walter \\etal\\ 2004; Peng \\etal\\ 2005) also imprints scatter and skewness into the local relation as a function of \\mbh\\ (Robertson \\etal\\ 2005).  However, it is unclear at this stage how much scatter can be attributed to the fact that BHs in AGNs, in so far as they are radiating and thus accreting, are still gaining mass. Any differential growth between BHs and bulges will introduce additional scatter to the AGN \\msigma\\ relation if AGN accretion and star formation are not precisely synchronized (Ho 2005; Kim et al. 2005).  If we ascribe all the observed scatter to differential growth (taking the observed scatter of $\\sim 0.27$ dex in the relation for inactive objects from Tremaine \\etal), there can be no more than $\\sim$0.3 dex scatter (factor of 2) introduced by relative BH-bulge growth. This level of growth has only modest fuel requirements that are easy to sustain in nearby galaxies; a $10^6$ \\msun\\ BH requires only $0.022$~\\msun\\ yr$^{-1}$ to double its mass in a Saltpeter time ($4.5 \\times 10^{7}$ yr).  Finally, taking our entire data set at face value, we do find evidence for a shallower slope than the inactive \\msigma\\ relation.  Whether this is the result of flattening at low mass or a different slope in AGNs is difficult to determine at the present time, in the absence of more AGNs with \\mbh\\ $> 10^8$~\\msun.  In this regard, Wyithe (2005) argues that the slope for the inactive sample steepens considerably when the four smallest \\sigmastar\\ values are removed, independently suggesting flattening at low mass.  Flattening at low mass may be a result of the changing efficiency in AGN feedback in low-mass halos. For instance, the formalism of Vittorini \\etal\\ (2005) finds that a combination of decreased optical depth (proportional to galaxy radius) and cooling times shorter than a dynamical time (small, dense halos) allows these BHs to grow to larger relative masses without feedback limitations.  Alternatively, different growth modes between spheroidal and disk-dominated systems may result in a different final position on the \\msigma\\ plane.  On the other hand, we cannot preclude the possibility that a change in slope in the BLR radius-luminosity relation at low luminosity is responsible for the change in slope in the \\msigma\\ relation.  However, we note that a radius-luminosity slope closer to the theoretically preferred value of 0.5 (e.g.,~Kaspi \\etal\\ 2000) would only increase the observed discrepancy in slope. Reverberation mapping of low-mass BHs would be required to address this issue.  In summary, although we have increased the population of AGNs with \\sigmastar\\ measurements by a factor of nearly 6, we do not yet have a large enough sample, or adequate coverage of parameter space, to uniquely identify the cause or causes responsible for the observed differences between the \\msigma\\ relation of local active and inactive galaxies.  Future effort should focus on (1) further enlarging the sample of AGNs with robust \\sigmastar\\ measurements, (2) pushing the samples to the extremes of the mass distribution, particular above $10^8$~\\msun, (3) extending the luminosity coverage to include objects with a wider range in Eddington ratios, and (4) better characterizing the BLR radius-luminosity relation over a broader range of AGN properties than is currently available."
    },
    "0512/astro-ph0512182_arXiv.txt": {
        "abstract": "{We present numerical models of hydromagnetic instabilities under the conditions prevailing in a stably stratified, non-convective stellar interior, and compare them with previous results of analytic work on instabilities in purely toroidal fields. We confirm that an $m=1$ mode (`kink') is  the dominant instability in a toroidal field in which the field strength is proportional to distance from the axis, such as the field formed by the winding up of a weak field by differential rotation. We measure the growth rate of the instability as a function of field strength and rotation rate $\\Omega$, and investigate the effects of a stabilising thermal stratification as well as magnetic and thermal diffusion on the stability. Where comparison is computationally feasible, the results agree with analytic predictions. \\keywords {instabilities -- magnetohydrodynamics (MHD) -- stars: magnetic fields}} ",
        "introduction": "Magnetic fields probably play a significant role in the internal rotation of stars. Even a relatively weak magnetic field is sufficient to couple different parts of the star and maintain a state of nearly uniform rotation. For the interior of the present Sun, for example, a field of less than 1 gauss would be able to transmit the torque exerted by the solar wind through the interior (\\cite{Mestel:1953}). The observed rotation in the core of the Sun (\\cite{Chaplinetal:2001}) is quite uniform, suggesting that a magnetic field of this order or larger may actually be present. The progenitors of white dwarfs and supernovae go through giant stages in which the envelope rotates very slowly. The degree of coupling between core and envelope by a magnetic field in this stage will determine whether the rotation rates of pulsars and white dwarfs are just a remnant of the initial rotation of their progenitors, or if a secondary process must be responsible (\\cite{SprandPhi:1998}, \\cite{Spruit:1998dn}).  Models of gamma-ray bursts in which the central engine derives from the rapidly spinning core of a massive star (\\cite{Woosley:1993}, \\cite{Paczynski:1998}) also depend on the ability of the core to keep its high angular momentum for a sufficient period of time, in the face of magnetic spindown torques exerted by the slowly rotating envelope (\\cite{Hegeretal:2000}).  The uniform rotation of the solar core may be due to a magnetic field, but this field's origin, configuration and strength are not known. By analogy with the magnetic A-stars, one might speculate that a `fossil' magnetic field could exist in the core of the Sun {\\mk(\\cite{Cowling:1945}, \\cite{BraandSpr:2004}).} Since no net field is seen at the surface (averaged over the solar cycle), the radial component of such a fossil would however have to be weak -- of the order of a gauss or less. A field weaker than about this 1 G will quickly wrap up into a predominantly toroidal field, under the action of the remaining (weak) differential rotation in the core.  The predominantly toroidal magnetic field resulting from this process will not increase in strength arbitrarily. Eventually, the energy density in the field will become large enough that a magnetic instability will set in.  Analytic work, (e.g. \\cite{Tayler:1973}), shows that any purely toroidal field should be unstable to instabilities on the magnetic axis of the star (pinch-type instabilities, under the influence of the strongly stabilising stratification in a radiative stellar interior, or `Tayler instabilities' hereafter). The growth rates of these instabilities are expected to be of the order of the time taken for an Alfv\\'{e}n wave to travel around the star on a toroidal field line. This is very short compared to the evolutionary timescale of the star. In a star like the Sun, for example, with a field of $1000$ gauss, the growth timescale $r\\sqrt{4\\pi\\rho}/B$ would be of the order of years, if $r=R_\\odot/2$ is taken, and $\\rho=1.3 \\mathrm{g/cm}^3$.  {\\mk A magnetic field of this type can also be subject to other instabilities, such as the magnetic buoyancy (\\cite{Parker:1955}, \\cite{Gilman:1970}, \\cite{Acheson:1978}) and magnetic shear instabilities (\\cite{Velikhov:1959}, \\cite{Acheson:1978}, \\cite{BalandHaw:1992}). As was shown by Spruit (1999), the Tayler instability will be the first to appear as the strength of the toroidal field is increased. This is because with the magnetic buoyancy instability, as with all instabilities where displacements in the vertical are necessary, the stratification provides a strong stabilising force. The same is the case with the magnetic shear instability, whose effect in a stellar interior is very limited in comparison to its effect in accretion discs. In contrast, the Tayler instability occurs on the magnetic axis, where the magnetic field is perpendicular to gravity and the displacements caused by the instability are also perpendicular to gravity.}  In this paper we aim to test numerically the instability mechanism, and to verify that the predictions of the analytic work are relevant: that they cover all instabilities actually present in a system consisting of a predominantly toroidal field in a stable stratification. Much of the analytical stability analyses have been done under a local approximation. This can be shown to be exact for the case of adiabatic instabilities in a non-rotating star, but not for the more interesting cases in which rotation and the effects of magnetic and thermal diffusion are taken into account. Though it is not expected that major instabilities have been missed, numerical simulations can provide an important check.  It is much less certain how the magnetic field evolves once instability has set in. In a scenario developed by Spruit (2002), it is argued that the instability will lead to self-sustained dynamo action. The field remains predominantly toroidal, subject to decay by Tayler instability, but is continuously regenerated by the winding-up of irregularities produced by the instability. This scenario has been applied in stellar evolution calculations of the internal rotation of massive stars by Heger et al. (2003) and Maeder and Meynet (2003).  The balance between wrapping-up by differential rotation on the one hand and the destruction of the toroidal field by Tayler instabilities on the other determines the strength and configuration of the field, and the the rate at which it transports angular momentum through the star.  {\\mk The long-term goal is therefore to investigate the non-linear development of the instability. With 3-D numerical simulations we can determine how quickly an initial toroidal field decays (by reconnection across the magnetic axis), and to determine the type of magnetic field that is maintained by differential rotation in a stably stratified star, under the action of magnetic instabilities, and to develop from this a quantitative theory for the angular momentum transport by magnetic fields in stars. This is beyond the scope of this paper, but is looked at by Braithwaite (2005), where the operation of this differential-rotation driven magnetic dynamo is demonstrated.} ",
        "conclusions": "We set out in this study to verify, with numerical methods, the results of the existing analytical work on the instabilities of toroidal fields in stars, and to check that these results were relevant, e.g. that a predicted instability is not drowned out by stronger unstable modes that might have been missed, for example due to simplifying assumptions. The results have been largely positive.  Previous analytic work makes predictions of the dependence of instability conditions and growth rates on the parameters of an azimuthal magnetic field. These are the dependence of the field strength on distance $\\varpi$ from the axis (the index $p={\\rm d}\\ln B/{\\rm d}\\ln\\varpi$), the stability of the stratification (measured by the buoyancy frequency $N$, or equivalently the acceleration of gravity $g$), the rotation rate $\\Omega$, the effects of magnetic diffusion (diffusivity $\\eta$), and the reduction of buoyancy by thermal diffusion on small scales (diffusivity $\\kappa$).  First we checked the analytic prediction of which azimuthal orders $m$ should be unstable, depending on the value of $p$ (Eq. (\\ref{eq:mandp})). We have confirmed this prediction (see Sect. \\ref{sec:pandm}) for the cases $p=1$ and $p=2$, the former being considered the most important as it is this magnetic field which could plausibly be produced by the winding-up of a seed field by differential rotation. The dominant mode has $m=1$, a `kink' instability.  The theory predicts that in the adiabatic, unstratified case ($\\kappa=\\eta=N=0$) there is no threshold for instability, and that the growth rate $\\sigma$ is of the order of the angular frequency $\\omega_{\\rm A}=r/v_{\\rm A}$ of an Alfv\\'en wave travelling around the star on an azimuthal field line. The field strength and hence $\\omega_{\\rm A}$ varies through the computational volume, but since the instability is a local one, the growth should be dominated by the largest value of $\\omega_{\\rm A}$ in the volume, after an initial transient. The numerical results reproduce this well. In the best studied case, for example, a value $\\sigma/ \\omega_{\\rm A}=0.92$ was measured.  In the adiabatic case the prediction (Pitts \\& Tayler 1986) is that the instability is suppressed when the rotation rate exceeds the Alfv\\'en frequency $\\omega_{\\rm A}$. This was also verified, give or take a few percent at the most. The adiabatic case is somewhat singular with respect to the effect of rotation, however. Theory predicts that in the presence of strong thermal diffusion (${\\kappa/\\eta} \\rightarrow\\infty$) the threshold for instability disappears again, and that the growth rate is then of the order $\\sigma=\\omega^2_{\\rm A}/\\Omega$. We have not tested this dependence, since the calculations required for this limiting case are computationally rather demanding.  The effect of gravity and both magnetic and thermal diffusion were investigated. The effect of gravity was as expected -- above a certain value of $g$ the initial equilibrium is stable at a given vertical wavenumber $n$. This translates into a minimum unstable wavenumber which increases with increasing $g$.  The effect of magnetic diffusion on the shorter wavelengths was not exactly the same as that expected. It was found that at a given wavenumber $n$, a value of $\\eta$ of the order suggested by Acheson (1978) and Spruit (1999) did not kill the instability entirely, rather it reduced its growth rate by about half. An increase in $\\eta$ beyond this reduced the growth rate still further, but not to zero. Therefore it seems that magnetic diffusion alone cannot suppress the instability. It is possible that the discrepancy arises because the case $\\eta\\rightarrow \\infty$ may have a singular limit. The theoretical prediction was made for the case $\\kappa=0$, but the relevant parameter is actually $\\kappa/\\eta$. By analogy with other double-diffusive systems, instabilities must exist also in the case $\\kappa<\\eta$ which do not appear when $\\kappa=0$ is set from the beginning. Since this case $\\kappa<\\eta$ is not of much astrophysical relevance, we have not pursued this further.  Finally, the effect of thermal diffusion was tested. In the runs executed, it was expected that values of $\\kappa$ of the order of $9\\times10^{-5}$ and $2.7\\times10^{-4}$ would be needed to make a run with $g=g_{\\rm crit}$ behave like the runs with $g=0.8g_{\\rm crit}$ and $g=0.6g_{\\rm crit}$ respectively. This appears to be correct, given that these were only order-of-magnitude approximations.  The non-linear effect of the instability in all the above cases was found to be rather simple. The net effect is similar to that of an enhanced magnetic diffusivity: the field configuration spreads horizontally, while the mean azimuthal magnetic flux decreases due to effective reconnection across the magnetic axis. Toroidal fields in stars are therefore predicted to decay quickly by Tayler instability, once conditions for instability are satisfied, unless regenerated by differential rotation. {\\mk Work continues on the non-linear evolution of this instability, and some first results can be found in Braithwaite (2005).}"
    },
    "0512/astro-ph0512527_arXiv.txt": {
        "abstract": "{ We consider the sample of weakly active galaxies situated in 'Local Universe' collected in the paper of Pellegrini (2005) with inferred accretion efficiencies from $10^{-2}$ to $10^{-7}$. We apply a model of spherically symmetrical Bondi accretion for given parameters ($M_{BH}$,$T_{\\infty}$,$\\rho_{\\infty}$,) taken from observation. We calculate spectra emitted by the gas accreting onto its central objects using Monte Carlo method including synchrotron and bremsstrahlung photons as seed photons. We compare our results with observed  nuclear X-ray luminosities $L_{X,nuc}$ (0.3-10 keV) of the sample. Model is also tested for different external medium parameters ($\\rho_{\\infty}$ and $T_{\\infty}$) and different free parameters of the model. Our model is able to explain observed nuclear luminosities $L_X$ under an assumption that half of the compresion energy is transfered directly to the electrons. ",
        "introduction": "It is believed that most of the galactic nuclei host supper-massive black holes in their centers. The masses of these central objects estimated by various methods are of the order of $10^6-10^9 \\rm{M_{\\odot}}$. The radiation emerging from the vicinity of these objects is produced by an accretion process. We believe that Quasars and Seyfert galaxies are powered by the flows with high angular momentum and are characterized by high accretion efficiency. Corresponding models consisting of thin accretion disk surrounded by some kind of a corona are able to explain all spectral features like big blue bump, iron line and short time-scale variability. The accretion luminosity in these sources is of order of a few percent of the Eddington limit or more. Nevertheless there exist a number ($40\\%$ of nearby galaxies, Ho 2003) of galactic nuclei which are very faint objects although they contain black holes of masses of the same order as those in AGN. For example Galactic Center hosts a super-massive object of mass $\\sim 10^6 \\rm{M_{\\odot}}$ (Genzel et al. 2003). The X-ray luminosity of this object is about $10^{-9}$ times smaller than its Eddington limit. The density and the temperature near the capture radius were estimated from Chandra data (Baganoff et al. 2003), which allowed to compute the expected mass accretion rate, and calculate the dynamics of the accreting gas using for example the simplest Bondi model (1952) of the steady spherically symmetric flow. The efficiency of this accretion flow, as estimated from comparison of expected accretion rates and observed luminosities, is about $\\eta= 10^{-5}$, which is very small in comparison to the typical active quasar value where efficiency is $\\eta=0.1$. Galactic Center is an example of the extremely faint source, sometimes it is even called inactive (Nayakshin 2004). In paper of Pellegrini (2005, hereafter P2005) there are collected more examples of faint galactic nuclei sources (LLAGN, Low Luminosity Active Galactic Nuclei) located in the 'Local Universe'. The masses of black holes in centers of these objects obtained by various methods (for details see P2005) lay in the same range as typical $M_{BH}$ AGN. Their luminosities are of the order of $10^{-2}-10^{-7}$ of the Eddington luminosity. The efficiency of the accretion (assuming Bondi mass accretion rate) is thus also low. LLAGN have no big blue bump (Quataert et al. 1999), weak iron lines and no short time-scale variability (Ptak et al. 1998) and this is an evidence of nonexistence of the accretion disk inside the flow (Nayakshin 2004 and references therein). To build the consistent picture we would like those objects to be quiescent phases of long-term evolution of quasars (Nayakshin 2004). To explain such a low activity usually we assume one of the RIAF (radiatively inefficient accretion flow) solutions. Examples of RIAFs include Bondi flow with no angular momentum, low angular momentum flows (Proga 2003) and ADAF with high angular momentum (Narayan 2002). Other possibilities are spherical flows with magnetic fields, convection dominated flows (Narayan 2002 and references therein) or jet-wind accretion flows (Yuan 2003).\\\\ In this paper, to explain low luminosities of the sample of LLAGN, we apply the spherically symmetrical accretion in the Newtonian regime (Bondi 1952). In this model accretion rate is determined by external medium conditions and remain constant in the whole radius range, thus in this model we do not include an outflow in any form. Bondi flow can be also treated as rough approximation of very low angular momentum flows in a steady phase. Plasma accreting onto a central object in our model is a source of synchrotron and bremsstrahlung photons. We assume that plasma has a two-temperature structure. Because in our model cooling of the flow occur by electrons, ions are much hotter than electrons near the horizon. In calculations of emerging radiation spectrum we include the density and the temperatures of ions and electrons as a function of radius. Radiative transfer of synchrotron and bremsstrahlung radiation through the plasma is calculated using Monte Carlo technique (Gorecki $\\&$ Wilczewski 1984, Pozdnyakov et al.1983). Because of the method of calculating the spectra, we neglect the radiation pressure influence on the flow dynamics.  In Sect. 2 we describe the model and the technique of calculating the comptonization in details. We present the result in Sect. 3. Discussion is given in Sect. 4. ",
        "conclusions": "In this paper we reconsider the sample of 17 LLAGNs (with measured temperature and density of surrounding medium) and GC collected in paper of P2005. Usually LLAGN are discussed in context of ADAF model (Quataert 1999, Di Matteo et al. 1999, Di Matteo et al. 2003, Loewenstein et al. 2001). In general ADAF over-predicts measured luminosities, if one assume that the mass accretion rate of the flow is equal to Bondi mass accretion rate. As it is pointed out by Narayan (2002) (also Quataert 2003), one should include into calculation mass accretion rate reduced by the $\\alpha$ factor (dimensionless viscosity parameter) so that $\\dot{M}_{ADAF}=\\alpha \\dot{M}_{Bondi}$, to make the model self-consistent. Although ADAF model, with proper assumption of $\\dot{M}$, predicts well X-ray luminosities in some cases, it can over-predicts emission in radio band (Loewenstein 2001). Also in case of the sources with jet structures, it is hard to include additional emission in radio band (Narayan 2002). The second point concerns estimation of the Bondi mass accretion rate $\\dot{M}_{Bondi}$. $\\dot{M}_{Bondi}$ may be an order of magnitude smaller depending on polytropic index $\\gamma$ and mean particles mass $\\mu$. E.g. for NGC1399 the mass accretion rate assuming $\\gamma=5/3$ and $\\mu=0.5$ is $\\dot{M}_{Bondi}=4.6 \\cdot 10^{-3} M_{\\odot}/yr$, but if we assume different $\\gamma=1.4$ (which is typical for the partially ionized ISM and usually assumed in ADAF models) and different $\\mu=1.0$ we obtain an order of magnitude higher mass accretion rate $\\dot{M}_{Bondi}=4.2 \\cdot 10^{-2} M_{\\odot}/yr$. One should pay a particular attention in choosing constants determining accretion rate, because emerging spectrum depends strongly on this quantity. In our calculations we take $\\gamma \\approx 5/3$ and $\\mu=0.5$.  Narayan (2002) also points out that LLAGN weak luminosities can be be fitted to the data assuming two-temperature Bondi model. In this paper we apply a spherically symmetrical Bondi model of accretion to the sample of sources from P2005, and we calculate the radiation spectra emerging from accreting plasma, including a full treatment of radiation transfer through the gas. Historically, the problem of spherical accretion onto a compact objects was present in many papers (for review see Nobili et al. 1991). Although model of steady spherical flow is only a mathematical model (not realized in nature because there is always some angular momentum present), it can be used as a rough approximation of an accretion with very small angular momentum (as we would expect to be present in elliptical galaxies). The similar model of spherical accretion was tested by Yim and Park (1995) in case of our Galactic Center. (Detailed comparison to Yim and Park(1995) is not possible because the paper is unavailable). Melia (1994) considered semi-spherical flow (with very low angular momentum) but with the disc inside. Authors in both papers conclude that spectrum can be reasonably explained with such assumptions.  Our calculations also show that spherical accretion can reconstruct observed X-ray luminosities in most cases of the LLAGN sources presented in P2005, within an order of magnitude error (10 cases in 17 galaxies in a sample), for parameter $\\delta$=0.5. The nuclear luminosities $L_X$ are very sensitive on the changes of the heating parameter $\\delta$. This parameter indicate how much accretion energy will be transfered to the electrons, and  determine the electron temperature profile. ($\\beta$ parameter does not influence the electron temperature significantly) For most of the sources X-ray luminosities are strongly under-predicted when $\\delta=0.001$. The obtained efficiency $\\eta$ of spherical accretion in all cases for this value of $\\delta$ is very small, of the order of $10^{-7}-10^{-9}$. If we assume $\\delta$=0.5, the electrons are heated in the same degree as ions. Half of the accretion energy is transfered to ions, the other half to the electrons. But the increase in X-ray luminosity is not caused by the growth of bremsstrahlung emissivity. When the temperature grow, the bremsstrahlung emission goes up like $T^{0.5}$ and it is more sensitive to the changes of density (like $\\rho^2$), than to the electrons temperature. So the effect if not very strong for the emissivity of bremsstrahlung radiation. The reason of increase of $L_X$ is that the synchrotron emissivities grows, and X-ray part of the spectrum is dominated by the Compton scattered synchrotron photons. The efficiency grows even 5 orders of magnitude for $\\delta$=0.5. Our calculations show that the assumption, that electrons are heated directly by accretion in the same degree as ions, is more proper that assuming $\\delta$ to be 0.001 because of mass ratio ($m_e/m_p$) (e.g. Esin et al. 1997, Manmoto 2000) This is in agreement with the results of Bisnovatyi-Kogan $\\&$ Lovelace (1997). The authors argue that in a presence of magnetic field and plasma instabilities, gravitational energy is transfered predominantly to electrons. Also in work of Yuan (et al.2003) the assumption of $\\delta=0.55$ allows to accommodate the spectrum of Sgr A* with ADAF model.  Bremsstrahlung emissivity depends strongly on external medium conditions ($T_{\\infty}$ and $\\rho_{\\infty}$) which control the mass accretion rate. Values of these quantities are estimated by observations, so they rather cannot be changed significantly in this approach. $L_X$ in case of over-predicted sources could be reduced and accommodated to observation with assumption of an accretion with outflows. It was considered by e.g. Quataert $\\&$ Narayan (1999). For one of the sources (Galactic Center), nuclear luminosity $L_X$ was strongly over-predicted by our model for $\\delta=0.5$. External medium measurements indicate (assuming Bondi flow) mass accretion rate for this source to be $\\sim 10^{-6} \\rm{M_{\\odot}/yr}$. On the other hand, we have also mass accretion rate limitations estimated by measurements of polarization of radiation near the black hole in Sgr A*. From Faraday rotation we obtain limit for $\\rm{\\dot{M}} \\sim 10^{-7} \\rm{M_{\\odot}/yr}$ (Atiken et al. 2000). This fact also additionally eliminate simple ADAF model for this source (Bower et al. 2005). Differences between $\\rm{\\dot{M}}$ near the capture radius and $\\rm{\\dot{M}}$ near horizon indicate that most of the accreting material do not reach a black hole. For Sgr A* the possible explanation could be the accretion flow disturbed by the outflow. For Sgr A* the radio part of a spectrum is too high in our model for $\\delta=0.5$, also it is higher than NIR limitations for mass accretion rate $\\dot{M} \\sim 10^{-6} \\rm{M_{\\odot}/yr}$. The outflow could be a possible explanation of spectrum and also reduce mass accretion rate. For $\\delta=0.001$ the emission was very low, the peak of the synchrotron radiation attain in $\\nu F_{\\nu}=10^{34}$ [ergs/s]. ADAF model in paper of  Narayan (et al. 1998) gives a more denser accretion flow than Bondi flow, since the radial velocity is smaller. Thus our results  for $\\delta$=0.001 are much below the observational points. In Yuan (et al. 2003) the emission is accommodated to X-ray emission also in assumption of higher $\\delta=0.55$ value  but for  $\\dot{M}=10^{-8} M_{\\odot}/yr$. The reason why our model for Sgr A* gives so poor constrains for $\\delta=0.5$ in comparison to Yuan (et al. 2003) is that we assume Bondi accretion rate which is two orders of magnitude larger than  $\\dot{M}=10^{-8} M_{\\odot}/yr$. Our model can reproduce better the Sgr A* spectrum if we allow for the specific adjustment of the parameters like $\\dot M$ and $\\delta$ since the synchrotron emissivity is very sensitive to the temperature changes (controlled mainly by these two parameters). Additionally, our model can also explain the soft spectral slope seen in the data (Baganoff et al. 2003) since the comptonized synchrotron component can extend to the X-ray band. However, the model never fits the radio frequencies below $\\sim 10^{11}$ Hz. This part of the spectrum can be only explained by models which include non-thermal populations of electrons, like e.g. ADAF-jet model of Yuan et al. (2002). The presence of a jet-like outflow also in other sources is supported by observations (see Tab.~2). Also the negligence of the spherically symmetric outflow is a serious weakness of our model since there are now observational evidences for Sgr A* that accretion rate is not constant with radius (e.g.Bower et al. 2003).  The effect of the angular momentum of the accretion flow also cannot be neglected (the low angular momentum was considered by Proga et al. 2003, but without any estimations of emerging spectrum). Bondi spherical accretion model is a mathematical model. In reality there is always angular momentum. We also know that standing shocks are the part of a low angular momentum flows (Das 2002, Das, Pendharkar $\\&$ Mitra 2003). Also tubulent flow may create shocks. In our model we assume the thermal distribution of electrons. This is a weakness of the model, since the shock produce some fraction of non-thermal electrons. Thanks to studies of Sgr A* (Mahadevan 1999, Ozel et al. 2000, Yuan et al. 2003) with hybrid (thermal + non-thermal power-law tail) population of electrons we know that observed low-frequency radio spectrum of Sgr A* can be explained if small fraction of electrons is non-thermal. Close to a black hole turbulence and magnetic reconnection can accelerate electrons. Synchrotron emission from this electrons and Compton scattering on them can be a reason of flaring in broadband spectrum. In all calculations of spectra for M87 the low-frequency radio emission wasn't reconstructed, this may be a reason why non-thermal electrons should be included into calculations, moreover scatterred nonthermal photons can affect the X-ray band of the spectrum.  Our model is also not consistent with recent results of modeling spherical accretion with magnetic fields. Because we include synchrotron emission in calculating the radiation spectrum we assume that there is some random magnetic field in plasma surrounding a black hole. Spherical flows including magnetic field are unstable (e.g. Igumenshchev et al. 2002), thus we are able to obtain only some mean luminosity of spherical accretion. Our model is not able to reconstruct any detailed features of time-dependent spectrum like e.g. flares. Another disadvantage of our model is that the equations of spherically symmetrical flow are calculated in Newtonian regime. Including relativistic effects may be important very close to a horizon of a black hole, where in our calculations most of the synchrotron radiation comes from."
    },
    "0512/hep-ph0512068_arXiv.txt": {
        "abstract": "Several phenomenological and cosmological aspects of a minimal extension of the Georgi-Glashow model, where the Higgs sector is composed by $\\bm{5}_H$, $\\bm{15}_H$, and $\\bm{24}_H$, are studied. It is shown that the constraints coming from the unification of gauge interactions up to two-loop level predict light scalar leptoquarks. In this GUT scenario, the upper bound on the total proton decay lifetime is $\\tau_p \\leq 1.4 \\times 10^{36}$ years. The possibility to explain the matter-antimatter asymmetry in the universe through the decays of $SU(2)_L$ scalar triplets is also studied. We find that a successful triplet seesaw leptogenesis implies an upper bound on the scalar leptoquark mass,  $M_{\\Phi_b} \\lesssim 10^{6-7}$~GeV. We conclude that this GUT scenario can be tested at the next generation of proton decay experiments and future colliders through the production of scalar leptoquarks. ",
        "introduction": "Grand unified theories (GUTs) based on the $SO(10)$ gauge symmetry~\\cite{SO(10)1,Fritzsch:1974nn} are usually considered as the most appealing candidates for the unification of electroweak and strong interactions. They offer a number of advantages over $SU(5)$ theories~\\cite{GG}: (i) They provide a natural explanation of the smallness of neutrino masses through the seesaw mechanism~\\cite{seesaw}; (ii) they accommodate all fermions of one generation into one representation; (iii) they represent, in their minimal form, the most promising theory of fermion masses (For realistic grand unified theories based on the $SO(10)$ gauge symmetry see e.g. Refs.~\\cite{Babu:1998wi,Aulakh:2003kg,Babu:2005gx}.).  Nevertheless, it is well known that the only promising way to test the idea of grand unification is through nucleon decay. Therefore, it is very important to investigate the simplest realistic grand unified theory where proton decay can be well predicted. This crucial issue brings us back to non-supersymmetric GUT scenarios, since the unification scale is rather low $(M_{GUT} \\approx 10^{14}~ \\textrm{GeV})$. In particular, we focus on the simplest realistic $SU(5)$ theories, where the unification scale can be accurately predicted. Even though $SU(5)$ possesses uncorrelated regions in the Yukawa sector, the simplicity of the Higgs sector in the non-supersymmetric case offers a hope that the theory can be verified in near future.  In a recent work~\\cite{Dorsner:2005fq}, some of us argued that the simplest realistic extension of the Georgi-Glashow (GG) model is the one containing the $\\bm{5}_H$, $\\bm{15}_H$ and $\\bm{24}_H$ representations in the Higgs sector. The purpose of this paper is to demonstrate that the next generation of collider and proton decay experiments will refute or verify this minimal $SU(5)$ scenario. A first attempt was made in Ref.~\\cite{Dorsner:2005fq} in this direction. Here we offer the full two-loop treatment of the gauge coupling unification and discuss the constraints on the Higgs sector. To show the testability of the model, we include all the presently available experimental limits in our discussion. The upper bound on the total proton decay lifetime in our GUT model is corrected, and we investigate the possibility to explain the baryon asymmetry observed in the universe through the decays of $SU(2)_L$ scalar triplets living in $\\bm{15}_H$, showing how important the constraints coming from leptogenesis turn out to be. The latter, when combined with the unification constraints, lead us to conclude that the present GUT scenario could be tested at the next generation of collider experiments through the production of light leptoquarks. ",
        "conclusions": "\\label{conclusions}  We have investigated in detail the constraints coming from unification of gauge interactions in the minimal extension of the Georgi-Glashow model, where the Higgs sector is composed by $\\bm{5}_H$, $\\bm{15}_H$ and $\\bm{24}_H$. We have shown that the scalar leptoquark $\\Phi_b$ has to be light in order to achieve unification in agreement with all experimental constraints. Using the constraints coming from triplet seesaw leptogenesis, the upper bound on the leptoquark mass is $M_{\\Phi_b} \\lesssim 10^{6-7}$ GeV. Therefore there is a hope that our scenario could be tested at the next generation of collider experiments through the production of these light leptoquarks.  We have also predicted an upper bound on the total proton decay lifetime which is $\\tau_p \\leq \\ 1.4 \\times 10^{36}$ years. Since at the next generation of proton decay experiments the bounds are expected to be improved by a few orders of magnitude, this minimal non-supersymmetric $SU(5)$ model will be certainly tested or ruled out.  The upper bound on the proton decay lifetime and the exciting possibility to verify the model at future collider experiments make our GUT scenario an appealing candidate for the testability of the idea of grand unification."
    },
    "0512/astro-ph0512241_arXiv.txt": {
        "abstract": "Non-thermal TeV $\\gamma$-ray emission within a multiparsec has been observed from the center region of our Galaxy. We argue that these $\\gamma$-rays are the result of transient activity of the massive black hole Sgr A$^*$ that resides at the Galactic Center. Several thousand years ago, the black hole may have experienced an active phase by capturing a red giant star and forming an accretion disk, temporarily behaving like an active galactic nucleus. A powerful jet, which contains plenty of high speed protons, was launched during the process. These runaway protons interact with the dense ambient medium, producing TeV $\\gamma$-ray emission through the $\\pi^0$-decay process. We show that the total energy deposited in this way is large enough to account for observations. The diffusion length of protons is also consistent with the observed size of the TeV source. ",
        "introduction": "It is known that there are many remarkable high-energy sources harbored in the Galactic center (GC) region \\citep{Mel01}. Recently, TeV $\\gamma$-ray emission from the direction of the GC has been reported by three independent groups, Whipple \\citep{Kos04}, CANGAROO (Collaboration of Australia and Nippon for a Gamma Ray Observatory in the Outback; \\citet{Tsu04}), and HESS(High Energy Stereoscopic System; \\citet{Aha04}). At least four potential candidates are suggested for this TeV $\\gamma$-ray emission, which are the black hole Sgr $A^*$ \\citep{Ahar05, Ato04, Lev00}, the compact and powerful young supernova remnant (SNR) Sgr A East \\citep{Cro05}, the dark matter halo \\citep{Hor05, pro05, Gne04,Ell02}, and the whole diffusion {\\rm 10\\,pc} region \\citep{Aha05}. Interestingly, the angular scale of the TeV source was determined by HESS to be less than a few arc-minutes, indicating that this $\\gamma$-ray source is located in the central $\\leq 10\\,pc$ region \\citep{Aha04}. It suggests that the black hole Sgr $A^*$, with a mass of $2.6\\times 10^6M_\\odot$ \\citep{Sch02}, should be involved \\citep{Ahar05}.  Recently, a hadronic origin of TeV $\\gamma$-rays that is linked to the massive black hole has been addressed in detail by \\citet{Aha05}. They argued that the TeV $\\gamma$-rays are produced indirectly through the processes of $\\pi^0$-decay when relativistic protons are injected into the dense ambient gas environment. The flux of this radiation component depends on the density of the target, the diffusion speed of protons in the interstellar medium, and the injection rate of protons. As a result, a dense gas target and an extremely large proton flux is required. However, since the black hole Sgr $A^*$ currently only emits faint electromagnetic radiation, it is largely uncertain how the protons could be accelerated to relativistic speeds.  It is well known that jets associated with accretion disks surrounding black holes are efficient in accelerating particles. For example, TeV $\\gamma$-rays observed from several BL Lac objects (a subclass of active galactic nuclei) are argued to originate from relativistic jets \\citep{Pia98}. The jet model is also used to explain the production of TeV $\\gamma$-rays from microquasars \\citep{bos05}. However, such a jet model cannot be applied directly to the black hole Sgr $A^*$.  An extraordinarily low bolometric luminosity of $\\sim 10^{36}\\,ergs\\,s^{-1}$ has been estimated for the black hole through multi-wavelength observations, which indicates that Sgr $A^*$ is in its quiescent dim state, and no powerful jet exists at present.  We propose that the black hole Sgr $A^*$ could be re-activated and produce a powerful jet by capturing a red giant (RG) star, temporarily behaving like an active galactic nucleus (AGN). This may have happened thousands of years ago, naturally providing a mechanism to generate the proton flux required by the hadronic model of Aharonian \\& Neronov. The structure of our paper is as follows. The capture process is described in Sect.2. We estimate the model parameters and the energy of the proton flux available for the production of TeV $\\gamma$-rays in Sect.3. A brief discussion and conclusion are presented in Sect.4 ",
        "conclusions": "A strong TeV $\\gamma$-ray source has been detected at the Galactic center,  whose size is probably less than 10 pc. It should be closely related to the black hole Sgr A$^*$ of our Galaxy. Aharonian \\& Neronov (2005b) suggested a hadronic origin for these $\\gamma$-rays; i.e., they are produced through $\\pi^0$-decay process when a strong flow of relativistic protons interacts with the dense ambient gas. However, the nature of this proton flow is largely unknown. In this paper, we show that the required proton outflow could have been reasonably produced when the black hole Sgr A$^*$ captured a RG star and formed an accretion disk around it in the past. The whole process can be divided into three phases. In phase 2, which lasts for $\\sim 10^4$ yr, the black hole becomes active and luminous thanks to a relatively high accretion rate, temporarily behaving like an AGN. A relativistic outflow of protons can be ejected in this phase, whose total energy is as high as $E_{j,tot} \\sim 2.76 \\times 10^{51}$\\,ergs, large enough to meet the requirement of the Aharonian \\& Neronov model.  According to our calculations, the propagation of protons in the target gas is through the ECP scenario. We show that the injected protons have diffusiond into a volume of $\\sim 9.78\\,pc$, consistent with the observationally inferred size of the TeV source. Furthermore, when the 10 TeV protons diffuse in the target gas, their escape time is comparable with the characteristic time of proton-proton interactions, implying that the spectrum of the observed TeV $\\gamma$-rays should be similar to that of the injected protons \\citep{Aha05}. Since the $\\gamma$-ray spectrum observed by HESS is, $J(E)=(2.5\\pm0.21)\\times 10^{-12}E^{- 2.21}\\,photon\\,(cm^2\\,s\\,TeV)^{-1}$ \\citep{Aha04}, we can conjecture that the initial spectrum of the injection protons should be $Q(E) \\propto E^{-2.2}\\exp({-E/E_0})$. Aharonian \\& Neronov (2005b) pointed out that the cutoff energy is $E_0=10^{15}\\,eV$.  In our study, we mainly consider the case in which the captured star is a RG star. If the captured star is a MS one, things will be much different. In Table 1, we have listed the key quantities calculated for a MS star capture, comparing them directly with those of a RG star capture. Interestingly enough, we find that the energy injection by a MS star capture also meets the requirement of energetics. However, we note that $t_{crit}>t_{acc}$ for a MS star; this means that the jet formed by the capture of a MS star terminates much earlier. In this situation, the diffusion timescale of the injected protons is $t_{dif}\\sim t_{acc}+ 300\\,yr$, which is much smaller than the diffusion time of a RG capture event. Consequently, within the jet lifetime, the protons originating from a MS star capture will diffusion to \\begin{eqnarray} R=\\sqrt{4D(10\\,TeV)t_{dif}} &\\approx& \\left\\{\\begin{array} {r@{\\;,\\quad}l} 1.05\\,pc& {\\rm for\\;the\\; ECP\\;}, \\,\\,\\,\\\\ 16.4\\,pc & {\\rm for\\;the\\;KTT\\;}, \\,\\,\\,\\\\ 105 \\,pc & {\\rm for\\;the\\;BD\\;}, \\,\\,\\,\\nonumber \\end{array} \\right. \\end{eqnarray} which is inconsistent with the size ($\\sim 10$ pc) of the TeV source inferred from current observations. This is the main reason that we prefer a RG star capture rather than a MS star capture in our framework.  However, to determine the size of the TeV source observationally is not an easy task because of its extended nature. We note that the HESS collaboration has not yet published the final analysis of its observations of the Galactic center on larger scales, so there is still a lack of information on TeV emission beyond the 10 pc scale at the area. It is thus possible that the actual size of the TeV source may be larger. In addition, the reduction of TeV emission may also be due to the decrease of the density of the gas at larger distances, but not the lack of an ultra-relativistic proton flow. Taking into account these factors, the MS star capture scenario still cannot be completely excluded. In fact, a comprehensive analysis of the TeV radiation and a thorough investigation of the environment within $\\sim 100$ pc around Sgr A$^*$ are necessary for us to understand the history of the activities in the GC region.   We thank the anonymous referee for valuable comments and suggestions that lead to an overall improvement of this study. We are grateful to S. N. Zhang for very valuable discussions and thoughtful comments. Thanks also goes to W. Wang for helpful discussions. This research was supported by a RGC grant of Hong Kong government, by the National Natural Science Foundation of China (Grants 10273011, 10573021, 10433010, 10233010, and 10221001), by the Special Funds for Major State Basic Research Projects, and by the Foundation for the Author of National Excellent Doctoral Dissertation of P. R. China (Project No: 200125."
    },
    "0512/gr-qc0512030_arXiv.txt": {
        "abstract": "The experimental evidence that the equation of state (EOS) of the dark energy (DE) could be evolving with time/redshift (including the possibility that it might behave phantom-like near our time) suggests that there might be dynamical DE fields that could explain this behavior. We propose, instead, that a variable cosmological term (including perhaps a variable Newton's gravitational coupling too) may account in a natural way for all these features. ",
        "introduction": "The accelerated expansion of the universe is nowadays one of the central issues of observational and theoretical cosmology. The usual paradigm assumes the existence of the so-called {\\em dark energy} (DE)-- a mysterious cosmic component with negative pressure. An obvious candidate for the role of DE is the cosmological constant (CC) \\cite{CCP}. Others are the dynamical DE models\\,\\cite{PSW}, e.g. quintessence \\cite{Quintessence}, phantom energy \\cite{phantom} etc.  Here we generalize the cosmological constant concept allowing the CC term $\\rho_{\\Lambda}$ (with dimensions of energy density) and possibly the gravitational coupling ($G$) being variable with the cosmic time, both assumptions compatible with the cosmological principle. The support for such a generalization comes from quantum field theory on the curved space-time \\cite{JHEPCC1,RGTypeIa,ShapSolNPB,RGTypeIa2,Babic,SSS} and/or quantum gravity approaches \\cite{Reuter}. ",
        "conclusions": "We have shown that a model with variable $\\rho_{\\Lambda}$ (and may be also with variable $G$) generally leads to a non-trivial effective EOS; the model can effectively appear as quintessence, and even as phantom energy, without need of invoking any combination of fundamental quintessence and phantom fields. This possibility should be taken into account in the next generation of high precision cosmology experiments aiming to determine the effective EOS of the DE.  \\ack This work has been supported in part by MEC, FEDER and DURSI. It is my pleasure to thank H. \\v{S}tefan\\v{c}i\\'{c} for a very nice collaboration and discussion of these matters."
    },
    "0512/astro-ph0512594_arXiv.txt": {
        "abstract": "Selective amplification of the line and continuum source by microlensing in a lensed quasar can lead to changes of continuum spectral slopes and  line shapes in the spectra  of the quasar components. Comparing the spectra of different components of the lensed quasar and the spectra of an image observed in different epochs one can infer the presence of millilensing, microlensing and intrinsic variability. Especially, microlensing can be used for investigation of the unresolved broad line (BLR) and continuum emitting region structure in active galactic nuclei (AGN). Therefore the spectroscopic monitoring of selected lensed quasars with 3D spectroscopy open new possibility for investigation of the  BLR structure in AGN. Here we discuss  observational effects that may be present during the BLR microlensing in the spectra of lensed QSOs ",
        "introduction": "\\label{sec:1}    Gravitational lensing is in general achromatic (the deflection angle of a light ray does not depend on its wavelength); however, the wavelength-dependent geometry of the different emission regions may result in chromatic effects (see Popovi\\'c \\& Chartas 2005, and references therein). Studies aimed at determining the influence of microlensing on the spectra of lensed quasars (hereafter QSOs) need to account for the complex structure of the QSO central emitting region. Since the sizes of the emitting regions are wavelength-dependent, microlensing by stars in a lens galaxy will lead to a wavelength-dependent magnification. The geometries of the line and the continuum emission regions are in general different and there may be a variety of geometries depending on the type of AGN (i.e. spherical, disc-like, cylindrical, etc.). Observations and modeling of microlensing of the broad-line region (BLR) of lensed QSOs are promising, because the study of the variations of the broad emission-line shapes in a microlensed QSO image could constrain the size of the BLR and the continuum region.   Our knowledge of the inner structure of quasars is very limited and largely built on model calculations. Continuum-line reverberation experiments with low-redshift QSOs tell us that the broad-emission line region (BLR) is significantly smaller than earlier assumed, and it is typically several light days up to a light year across (e.g., Kaspi et al.\\ 2000). It means that the BLR radiation could be  significantly amplified due to microlensing  by  (star-size) objects in an intervening galaxies (Abajas et al. 2002). Hence, gravitational lensing can provide an additional method for studying the inner structure high-redshift quasars for several reasons:  (i) the extra flux magnification, from a few to 100 times, provided by the lensing effect enables us to obtain high signal-to-noise ratio (S/N) spectra of distant quasars with less observing time;   (ii) the magnification of the spectra of the different images may be chromatic (as was noted in Wambsganss \\& Paczy\\'nski 1991, Wisotzki et al. 2003, Wucknitz et al. 2003, Popovi\\'c \\& Chartas 2005) because of the line and continuum emitting region are different in sizes and geometrically complex and/or complex gravitational potential of lensing galaxy; (iia) consequently, microlensing events lead to wavelength-dependent magnifications of the continuum that can be used as indicators of their presence (Wisotzki et al. 2003, Popovi\\'c \\& Chartas 2005);  (iii) gravitational microlensing can also change the shape of the broad lines (see Popovi\\'c et al. 2001, Abajas et al. 2002, Popovi\\'c \\& Chartas 2005),  the deviation of the line profile depends on the geometry of the BLR.  Finally, the monitoring  of lensed QSOs in order  to investigate the effect of lensing on the spectra  can be useful not only for constraining the unresolved structure of the central regions of QSOs, but also for providing insight to the complex structure of the lens galaxy. ",
        "conclusions": ""
    },
    "0512/astro-ph0512307_arXiv.txt": {
        "abstract": "% We study the colors, structural properties, and star formation histories of a sample of $\\sim1600$ dwarfs over look-back times of $\\sim3$ Gyr ($z=0.002-0.25$). The sample consists of 401 distant dwarfs drawn from the Galaxy Evolution from Morphologies and SEDs (GEMS) survey, which provides high resolution {\\it Hubble Space Telescope (HST)} Advanced Camera for Surveys (ACS) images and accurate redshifts, and of 1291 dwarfs at 10--90 Mpc compiled from the Sloan Digitized Sky Survey (SDSS). We find that the GEMS dwarfs are bluer than the SDSS dwarfs, which is consistent with star formation histories involving starbursts and periods of continuous star formation. The full range of colors cannot be reproduced by single starbursts or constant star formation alone. We derive the star formation rates of the GEMS dwarfs and estimate the mechanical luminosities needed for a complete removal of their gas. We find that a large fraction of luminous dwarfs are likely to retain their gas, whereas fainter dwarfs are susceptible to a significant gas loss, {\\it if} they would experience a starburst. ",
        "introduction": "The evolution of dwarfs is a complex problem, where evolutionary paths may depend on a variety of external and internal factors. Our knowledge of the local volume ($<8$ Mpc) has deepened, in particular due to strong efforts in determining distances to many nearby galaxies \\citep[and references therein]{kar03}. However, it is still unclear what governs the evolution of dwarfs in low density regions and how the different morphological types form.  Here, we present a study of the properties of dwarf galaxies over the last 3 Gyr ($z=0.002-0.25$)\\footnote{We assume a flat cosmology with $\\Omega_M = 1 - \\Omega_{\\Lambda} = 0.3$ and $H_{\\rm 0}$=70~km~s$^{-1}$~Mpc$^{-1}$.} drawn from GEMS \\citep{rix04} and SDSS \\citep{aba04}. ",
        "conclusions": ""
    },
    "0512/astro-ph0512131_arXiv.txt": {
        "abstract": "Photometric UBVI CCD photometry is presented for NGC 188 and Berkeley 17.  Color-magnitude diagrams (CMDs) are constructed and reach well past the main-sequence turn-off for both clusters.  Cluster ages are determined by means of isochrone fitting to the cluster CMDs.  These fits are constrained to agree with spectroscopic metallicity and reddening estimates.  Cluster ages are determined to be $7.0 \\pm 0.5\\,$ Gyr for NGC 188, and $10.0 \\pm 1.0\\,$Gyr for Berkeley 17, where the errors refer to uncertainties in the relative age determinations. These ages are compared to the ages of relatively metal-rich inner halo/thick disk globular clusters and other old open clusters. Berkeley 17 and NGC 6791 are the oldest open clusters with an age of 10 Gyr.  They are 2 Gyr younger than the thick disk globular clusters. These results confirm the status of Berkeley 17 as one of the oldest known open cluster in the Milky Way, and its age provides a lower limit to the age of the Galactic disk. ",
        "introduction": "Galaxy formation and evolution theory remains one of the great outstanding problems in contemporary astrophysics.  Although considerable progress has been made in this field over the past decades, there are many unanswered questions regarding the formation of galaxies like the Milky Way exists.  Observations within the Milky Way galaxy can be used to probe galactic evolution.  In particular, determining the relative ages of the different stellar populations in the Milky Way --- the halo, thick disk, thin disk, and bulge --- by dating open and globular stellar clusters provides significant insight into the chronology of Galaxy formation \\citep{liu00,sal04}.  Owing to the great importance of the old open cluster population in probing the chemical and dynamic evolution of the Galaxy, considerable effort has been devoted to determining the physical parameters of these clusters \\citep[e.g.][]{phe94,jan94,sco95}.  However, relatively few of the oldest open clusters have received commensurate attention in astrometric, photometric, and spectroscopic studies, despite the great promise they hold in determining the age of the Galactic disk (notable exceptions include NGC 188 and NGC 6791).  This paper presents accurate photometric UBVI photometry for old open clusters NGC 188 and Berkeley 17.  A number of excellent studies have previously been performed for NGC 188 that have accurately determined its physical parameters \\citep{von98, sar99, pla03,mic04,ste04,van04}. NGC 188 is used as a fiducial cluster, and a precise, relative age is obtained for Berkeley 17 relative to NGC 188.  Additionally, by directly comparing it with previously published data, the NGC 188 photometry serves as an independent test of the accuracy of the photometric calibration from the instrumental to the standard system.  Since the first photometric study of NGC 188 by \\cite{san62} demonstrating that it belonged to the oldest open clusters in the Galactic disk, the cluster has been the subject of numerous studies. As the highest priority cluster in the WIYN Open Cluster Study, excellent, multi-color CMDs and proper-motion data have been published for this cluster \\citep{von98,sar99,pla03}.  In their UBVRI photometric study, \\cite{sar99} found that NGC 188 has an age of $7.0 \\pm 0.5\\,$Gyr, a reddening of E(B--V) = 0.09 $\\pm$ 0.02, and a distance modulus of (m--M)$_V$ = 11.44 $\\pm$ 0.08.  These values are in general agreement with other recent photometric studies of the cluster \\citep{cap90,twa89}.  \\cite{ste04} obtained new data on NGC 188 and did a comprehensive review of existing data in the literature resulting in a large, homogeneous photometric database.  These data were used by \\cite{van04} to determine an age of $6.8\\pm 0.7\\,$Gyr for NGC 188.  Berkeley 17 was discovered by \\cite{set62}, and is located at a low Galactic latitude in the direction of the Galactic anticenter.  As a consequence of the cluster's location as well as its large distance, the field in the direction of Berkeley 17 is highly reddened and greatly contaminated by field stars.  Since the extensive open cluster study by \\cite{phe94}, a general agreement has prevailed that Berkeley 17 is indeed the oldest known open cluster in the Galaxy.  Soon after the Phelps et al.~study, \\cite{kal94} performed the first BVI photometry of the cluster, and established that Berkeley 17 is as old or somewhat older than the previously oldest known open cluster NGC 6791, based on comparisons of the morphologies of the two clusters' CMDs.  However, the Kaluzny study was hampered by poor weather, thereby preventing absolute determinations of Berkeley 17's physical parameters.  The first age determinations for Be 17 based on isochrone fitting were made by \\cite{phe97} who found an age of 10--13 Gyr, a metallicity of --0.30 $<$ [Fe/H] $<$ 0.00, a reddening of 0.52 $<$ E(B--V) $<$ 0.68 and 0.61 $<$ E(V--I) $<$ 0.71, and a distance modulus of (m--M)$_V$ = 14.05 $\\pm$ 0.25.  Although other attempts at establishing the age of Berkeley 17 have been carried out \\citep{car99b,sal04}, no further observational data has been published for the cluster.  Section \\ref{section2} describes the observations and the data reduction process, and \\S \\ref{n188photometry} presents cluster CMDs.  Section \\ref{section4} describes our stellar models, isochrone fits, and the age determinations of the clusters.  These ages are compared to the ages of other old open clusters and relatively metal-rich globular clusters in \\S \\ref{compare}. Section \\ref{section5} evaluates these results in the context of the formation of the Milky Way. ",
        "conclusions": "\\label{section5}  This paper confirms the well-established age of old open cluster NGC 188 at $7.0 \\pm 0.5\\,$Gyr.  Be 17 is found to be $10.0 \\pm 1.0\\,$Gry old.  As a result of the internal consistency of the ages obtained for the two clusters, the age difference of Be 17 relative to NGC 188 is highly robust: Be 17 is $3.0 \\pm 1.1\\,$Gyr older than NGC 188.  These ages were determined using the metallicity scale of \\cite{fri02}. There are some indications that the \\cite{fri02} [Fe/H] measurements are too low by about 0.1 dex as high resolution spectroscopy studies of M67 \\citep{hob91,tau00} find [Fe/H] values 0.1 dex higher than \\cite{fri02}.  If the \\cite{fri02} metallicities are 0.1 dex too low, then our ages need to be revised downward by about 8\\%.  Using the same age determination technique, the \\cite{fri02} metallicity measurements and photometry from the literature, ages of 16 other old, open clusters were determined.  Be 17 was found to be the oldest open cluster, with the same age as the well studied cluster NGC 6791.  The ages of 13 relatively metal-rich globular clusters were determined using the same methodology as the open cluster ages.  Table \\ref{allages} summarizes the age determinations.  These ages were determined using isochrone fitting in a self-consistent manner.  This leads to precise relative ages, but the absolute ages have considerably larger uncertainties.  We note that the globular cluster ages we derive are $\\sim 8\\%$ larger than the accurate absolute ages found by \\cite{kra03}.  This is the same age reduction estimated for the open clusters, based upon the suggestion that the \\cite{fri02} metallicity values are too low by 0.1 dex.  The globular clusters have ages in the range 11.8 -- 14.0 Gyr, with no evidence for an age range.  The thick disk globular clusters (the most metal-rich clusters in our sample) were found to have an average age of $12.3\\pm 0.4\\,$Gyr, while the two oldest open clusters in our sample have an age of $10.0\\pm 0.7\\,$Gyr.  Thus, the oldest open clusters in the thin disk are found to be $2.3\\pm 0.8\\,$Gyr younger than the thick disk globular clusters.  In contrast, \\cite{sal04} find no significant age difference between the thin and thick disks in the Galaxy.  However, their age determinations generally had larger error bars than our determinations, and \\cite{sal04} determined the age of two thick disk globular clusters, while we determine the age of three thick disk globular clusters.  Combined, these two effects made it difficult for \\cite{sal04} to determine if age differences less than 3 Gyr existed. In summary, we find that the oldest open clusters imply that the thin disk started forming $2.3\\pm 0.8\\,$Gyr after the formation of the thick disk."
    },
    "0512/astro-ph0512217_arXiv.txt": {
        "abstract": "High-energy photons from pair annihilation of dark matter particles contribute to the cosmic gamma-ray background (CGB) observed in a wide energy range. Since dark matter particles are weakly interacting, annihilation can happen only in high density regions such as dark matter halos. The precise shape of the energy spectrum of CGB depends on the nature of dark matter particles --- their mass and annihilation cross section, as well as the cosmological evolution of dark matter halos. In order to discriminate between the signals from dark matter annihilation and other astrophysical sources, however, the information from the energy spectrum of CGB may not be sufficient. We show that dark matter annihilation not only contributes to the mean CGB intensity, but also produces a characteristic {\\it anisotropy}, which provides a powerful tool for testing the origins of the observed CGB. We develop the formalism based on a halo model approach to analytically calculate the three-dimensional power spectrum of dark matter clumping, which determines the power spectrum of annihilation signals. We show that the expected sensitivity of future gamma-ray detectors such as the Gamma Ray Large Area Space Telescope (GLAST) should allow us to measure the angular power spectrum of CGB anisotropy, if dark matter particles are supersymmetric neutralinos and they account for most of the observed mean intensity of CGB in GeV region. On the other hand, if dark matter has a relatively small mass, on the order of 20 MeV, and accounts for most of the CGB in MeV region, then the future Advanced Compton Telescope (ACT) should be able to measure the angular power spectrum in MeV region. As the intensity of photons from annihilation is proportional to the density squared, we show that the predicted shape of the angular power spectrum of gamma rays from dark matter annihilation is different from that due to other astrophysical sources such as blazars and supernovae, whose intensity is linearly proportional to density. Therefore, the angular power spectrum of the CGB provides a ``smoking-gun'' signature of gamma rays from dark matter annihilation. While the mean CGB intensity expected from dark matter halos with smooth density profiles is smaller than observed, the dark matter substructure within halos may provide the origin of additional ``boost'' factors for the annihilation signal. Our formalism can be used for any other radiation processes that involve collision of particles. ",
        "introduction": "\\label{sec:Introduction}  The energy density of the universe is dominated by invisible components: dark matter and dark energy, both of which are of unknown origin. Remarkable progress in both the theoretical and observational studies of the cosmic microwave background (CMB) anisotropy, Type Ia supernovae, and large-scale structure has allowed us to precisely determine what fraction of the total energy these dark components convey, $\\Omega_\\chi = 0.23$ and $\\Omega_\\Lambda = 0.73$ \\cite{WMAP}, where the subscripts $\\chi$ and $\\Lambda$ denote dark matter particles and dark energy, respectively.  While the nature of dark energy is a complete mystery, there are several candidates for dark matter particles. Of which, the most popular candidate is the stable supersymmetric neutralino with mass on the order of GeV to TeV, which can explain the observed dark matter mass density today. Although very weak, it is expected that dark matter particles may interact with the usual matter via scattering, interact with themselves, and/or annihilate into gamma rays, positrons, and neutrinos. The direct detection of particle dark matter is, therefore, under intensive and extensive efforts of both physicists and astronomers \\cite{DMReview1,DMReview2,DMReview3}.  High-energy photons from annihilation of dark matter particles provide indirect means to probe the properties of dark matter. Annihilation signatures, especially gamma rays, have been searched for in regions where the dark matter density is expected to be high, as annihilation rate is proportional to the density squared, $\\rho_\\chi^2$. An obvious site is the central region of our Galaxy. Strong gamma-ray emission has been detected toward the Galactic center over a wide energy range, and the nature and origin of this emission have been investigated by many researchers \\cite{DMGammaGC1,DMGammaGC2,DMGammaGC3,DMGammaGC4,DMGammaGC5,DMGammaGC6,DMGammaGC7,DMGammaGC8,DMGammaGC9}. Although both the spectrum and angular distribution of gamma rays from the Galactic center are well observed, it is still quite difficult to distinguish a dark matter component from the other possibilities such as emission from ``ordinary'' astrophysical objects; even future gamma-ray detectors and telescopes do not have sufficient resolution to remove all the other sources from the Galactic central image.  Another possibility is the extragalactic background light, the cosmic gamma-ray background (CGB), which has been measured in a wide energy range \\cite{GeVCGB1,GeVCGB2,MeVCGB1,MeVCGB2,MeVCGB3}. It has been speculated that some fraction of the CGB may originate from annihilation of dark matter particles in halos distributed over cosmological distances \\cite{CGBDM1,CGBDM2,CGBDM3,CGBDM4,CGBDM5,A05,OTN05,AK1,AK2}. The dark matter contribution to the CGB depends on the nature of dark matter particles, such as their mass and annihilation cross section, as well as the cosmological evolution of dark matter halos. While it is likely that the dominant part of the CGB spectrum in the GeV region, detected by the Energetic Gamma Ray Experiment Telescope (EGRET) \\cite{GeVCGB1,GeVCGB2}, comes from unresolved blazars, i.e., beamed population of Active Galactic Nuclei (AGN), annihilating dark matter might give large contribution at some specific energy range.\\footnote{As has been known and is shown below, with the canonical choice of relevant parameters, dark matter annihilation cannot give enough contribution to the mean CGB intensity; one still needs large boost by some mechanism. In addition, its contribution is strongly constrained by gamma-ray observations toward the Galactic center \\cite{A05}. However, one might avoid these constraints by invoking gamma-ray emission from sub-halos having $M\\sim 10^{-6} M_\\odot$ within host halos \\cite{OTN05}, or from density cusps forming around intermediate-mass black holes \\cite{IMBH}.} Possible candidates for this case include the supersymmetric neutralino, as well as the Kaluza-Klein dark matter predicted by theories of universal extra-dimension (e.g., \\cite{DMReview3}) and heavy relic neutrinos \\cite{Neutrino1,Neutrino2}. On the other hand, the origin of the CGB in MeV region is much less understood than in GeV region. The soft gamma-ray spectrum in 1--20 MeV cannot fully be attributed to either AGN or Type Ia supernovae or a combination of the two \\cite{SNIa1,SNIa2}; thus, annihilation of dark matter particles is one of the most viable explanations for the CGB in MeV region, provided that the dark matter mass is around 20 MeV \\cite{AK1,AK2}. Such a light dark matter particle was originally introduced to explain the origin of the 511 keV emission line from the Galactic center \\cite{MeVDM,BBB}, detected by the International Gamma-Ray Astrophysics Laboratory (INTEGRAL) \\cite{INTEGRAL1,INTEGRAL2}, which is otherwise difficult to explain. Therefore, the MeV dark matter is an attractive candidate that satisfies the observational constraints from both the Galactic center and the CGB,\\footnote{After this paper has been submitted, Ref.~\\cite{beacom/yuksel:2005} has appeared on the preprint server. They claim that the mass of MeV dark matter should be less than 3 MeV from a more accurate treatment of relativistic positrons annihilating with electrons in the interstellar medium. While we use $m_\\chi=20$~MeV throughout this paper, one can easily extend our calculations to arbitrary dark matter masses.} while its particle physics motivation is less clear than for the neutralinos.  Dark matter annihilation may be a viable explanation for the CGB, but how do we know for sure that the CGB does come from annihilation? What would be a smoking-gun signature for the annihilation signal? We argue that {\\it anisotropy} of the CGB may provide a smoking-gun signature. Although the CGB is isotropic at the leading order, anisotropy should also exist if the CGB originates from cosmological halos. The future gamma-ray detectors with an enhanced sensitivity and angular resolution, such as the Gamma Ray Large Area Space Telescope (GLAST) or the Advanced Compton Telescope (ACT), should be able to see such anisotropy. We believe that the CGB anisotropy is going to be the key to discriminating between the dark matter annihilation signal and the other sources.  In this paper, we develop the formalism to analytically calculate the angular power spectrum of the CGB from dark matter annihilation. We calculate the angular power spectrum in GeV region (for supersymmetric neutralinos) as well as in MeV region (for MeV dark matter). We then discuss the detectability of CGB anisotropy in GeV region by GLAST and in MeV region by ACT, showing that the predicted anisotropy can be easily measured by 1-year operation of these experiments. The formalism given in this paper can also be used to evaluate the CGB anisotropy due to the other astrophysical objects with some modification. Note that the case of Type Ia supernovae has been discussed in Ref.~\\cite{SNIaAnisotropy}, being complementary to the present study.  This paper is organized as follows. In Sec.~\\ref{sec:Cosmic gamma-ray background: Isotropic component}, we review the mean intensity of the CGB from dark matter annihilation. In Sec.~\\ref{sec:Cosmic gamma-ray background anisotropy}, we develop the formalism to calculate the power spectrum of CGB anisotropy in the context of the halo models. Results of the angular power spectrum are then shown in Sec.~\\ref{sec:Angular power spectrum} for both the neutralino and MeV dark matter. We compare the predictions with the expected sensitivities of the future gamma-ray detectors. In Sec.~\\ref{sec:Discussion}, we discuss energy dependence of the obtained anisotropy (Sec.~\\ref{sub:Dependence on gamma-ray energy}), other astrophysical sources which might contaminate the cosmological CGB (Sec.~\\ref{sub:Other astrophysical sources}), and the effects of dark matter substructures (Sec.~\\ref{sub:Substructure of dark matter halos}). In Sec.~\\ref{sec:Conclusions}, we conclude the present paper with a brief summary. ",
        "conclusions": "\\label{sec:Conclusions}  The CGB provides indirect means to probe the nature of dark matter particles via high-energy photons from dark matter annihilation. The dark matter annihilation has occurred in all the past halos, and now contributes to the CGB flux at some level; its contribution might be dominant if the flux is boosted by some mechanism such as tidally survived dark matter substructure. Therefore, revealing the origin of the CGB is an very important problem that is potentially connected to the dark matter properties and halo substructure as well as ordinary astrophysical objects. In order to achieve this purpose, while the energy spectrum of the mean intensity of CGB has been investigated by many researchers, anisotropy of CGB  from dark matter annihilation has been entirely neglected. In this paper we have calculated the angular power spectrum of CGB from dark matter annihilation for the first time, using an analytical halo model approach. As for dark matter candidates, we have discussed two possibilities: one is the supersymmetric neutralino, one of the most popular candidates today, which contributes to the CGB in GeV region. The other is the MeV dark matter, a dark matter species first introduced to explain the 511 keV emission line from the Galactic center by annihilation into $e^-e^+$ pair, which contributes to the CGB in MeV region.  Since the gamma-ray intensity from annihilation is proportional to the density squared, $\\rho_\\chi^2$, we have derived the 3D power spectrum of the quantity $f = \\delta^2 - \\langle \\delta^2 \\rangle$, $P_f(k)$. The calculation involves the Fourier transformation of the four-point correlation function of underlying mass density fluctuations, $\\xi^{(4)}$, which has been shown to dominate over the two-point correlation contribution, $\\xi^{(2)2}$. This $P_{f,4} (k)$ includes the 1-halo and 2-halo terms, the former containing two points within the same halo, and the latter containing two points in two different halos. The analytical expressions for $P_f(k)$ from each term are given in Eqs.~(\\ref{eq:P_f,2}), (\\ref{eq:PS 1-halo term}), and (\\ref{eq:PS 2-halo term}), and the results are plotted in Figs.~\\ref{fig:Delta_2nd}, \\ref{fig:Delta}, and \\ref{fig:Delta_mass}. At all scales the four-point contribution totally dominates the signal; at small scales, the 1-halo term dominates.  We note that our formalism can also be used for any other emission processes that involve collisions of two particles. For example, one may use this to compute the power spectrum of free-free or bound-free emission from the ionized gas in halos. For this application one needs to replace the dark matter density profile with the gas density profile.  Using Eq.~(\\ref{eq:C_l}) that connects the 3D power spectrum, $P_f(k)$, to the angular power spectrum, $C_l$, we have calculated the CGB angular power spectrum as a function of multipoles, $l$, for both the neutralino ($m_\\chi = 100$ GeV) and MeV dark matter ($m_\\chi = 20$ MeV) at various gamma-ray energies. We have also compared the predicted signals with the expected sensitivity of future gamma-ray detectors --- GLAST in GeV region (for neutralinos) and ACT in MeV region (for MeV dark matter). The results are shown in Figs.~\\ref{fig:C_l_cut6}--\\ref{fig:C_l_MeV_cut-6}. For both cases, we have found that these detectors will have sufficient sensitivity to measure the angular power spectrum with reasonable accuracy.  We have studied the effects of the minimum mass, $M_{\\rm min}$, on the predicted angular power spectrum in detail. While the 1-halo contribution, which dominates at small angular scales (large $l$), decreases for smaller $M_{\\rm min}$, the 2-halo contribution, which dominates at large angular scales (small $l$), is virtually unaffected by $M_{\\rm min}$. This property results in a peculiar dependence of the shape of $C_l$ on $M_{\\rm min}$, which may be used in combination with the information from the energy spectrum of the CGB to determine $M_{\\rm min}$. As $M_{\\rm min}$ depends on the radiation processes of gamma rays from annihilation as well as the survival of micro halos contributing to the CGB, the shape of $C_l$ provides a powerful tool for determining these properties, which are otherwise difficult to probe. Our conclusion about prospects for measuring $C_l$ of CGB anisotropy are robust regardless of $M_{\\rm min}$.  By applying our formalism with some modification, we have shown that the other astrophysical sources such as blazars would reveal a different shape of the angular power spectrum, as these contributions are linearly proportional to density fluctuations. The shape of $C_l$ at large $l$ might further change when we take into account the dark matter substructure, but the result in this paper sets a lower bound on anisotropy at all multipoles, which provides excellent prospects for detection of CGB anisotropy by future gamma-ray detectors.  We conclude that the angular power spectrum of CGB provides a smoking-gun signature for gamma-ray emission from annihilation of dark matter particles, which would be a powerful tool for understanding the nature of dark matter particles."
    },
    "0512/astro-ph0512021_arXiv.txt": {
        "abstract": "The secular evolution of the orbital angular momentum (OAM), the systemic mass $(M=M_{1}+M_{2})$ and the orbital period of 114 chromospherically active binaries (CABs) were investigated after determining the kinematical ages of the sub-samples which were set according to OAM bins. OAMs, systemic masses and orbital periods were shown to be decreasing by the kinematical ages. The first order decreasing rates of OAM, systemic mass and orbital period have been determined as $\\dot J = 3.48 \\times 10^{-10}~yr^{-1}$ per systemic OAM, $\\dot M = 1.30 \\times 10^{-10}~yr^{-1}$ per systemic mass and $\\dot P = 3.96\\times 10^{-10}~yr^{-1}$ per orbital period respectively from the kinematical ages. The ratio of $d \\log J/ d \\log M  = 2.68$, which were derived from the kinematics of the present sample, implies that there must be a mechanism which amplifies the angular momentum loss $\\bar A = 2.68$ times in comparison to isotropic angular momentum loss of hypothetical isotropic wind from the components. It has been shown that simple isotropic mass loss from the surface of a component or both components would increase the orbital period. ",
        "introduction": "The observational evidence of decaying rotation rate for stars with spectral types later than F in the stellar evolution was well documented by Skumanich\\ (1972) by studying projected equatorial speeds $(v \\sin i)$ of the late-type stars in open clusters of different ages. Such a decay process in stellar rotation is explained in the terms of angular momentum loss (AML) through magnetically driven stellar winds, also called the magnetic braking (cf. Schatzman\\ 1959, Kraft\\ 1967 and Mestel\\ 1968). For the tidally locked binaries with late-type components, such an AML is known to be provided by the reservoir of the orbital angular momentum (OAM). Therefore, the AML from the components of spin-orbit coupled binaries causes the orbit to shrink. Spin-orbit coupling and a shrinking orbit, then, imposes spin-up the rotation of the components, which is different from single stars slowing down rotation. This mechanism is considered as the main way to form W UMa-type contact binaries from systems initially detached (cf. Huang\\ 1966, Okamoto \\& Sato\\ 1970, van't Veer\\ 1979, Vilhu \\& Rahunen\\ 1980, Mestel\\ 1984, Guinan \\& Bradstreet\\ 1988, Maceroni \\& van't Veer\\ 1991, Stepien\\ 1995, Demircan\\ 1999).  The period evolution and the time scale of forming the contact binaries from the detached progenitors were estimated differently among the various authors. In the work of Guinan \\& Bradstreet\\ (1988), the AML of a component star was computed directly from Skumanich's law ($V_{rot} \\sim t^{1/2}$), which is derived from the relatively slow rotating stars ($V_{rot}\\leq 17$ km s$^{-1}$), and then AML from the two components made equal to the orbital AM change. Orbits evolve initially almost with constant periods until the very end where the orbits shrink sharply to form contact binaries. For the braking law of binary orbits, van't Veer \\& Maceroni\\ (1988, 1989) gave a period evolution function which was predicted from the initial (Abt \\& Levy\\ 1976, Abt\\ 1983) and the present day (Farinella et al.\\ 1979) period distributions of G-type binaries corrected for selection and detectability effects. Stepien\\ (1995) derived a formula with no free parameters for the AML via a magnetized wind and calibrated it from the spin down of single stars. He concludes that usually the orbital periods which are less than 5 days can form contact systems within their main-sequence lifetime.  What is common among those works is the spin-orbit coupling set by synchronization. In other words, the above mechanisms do not work for asynchronous binaries as van't Veer \\& Maceroni\\ (1988) stated that the wider systems without spin-orbit coupling will not change the orbital period. Therefore, Demircan\\ (1999) studied the orbital AM distribution of 40 well-known CABs only with $P_{orb}<10$ days, and found an observational estimate of the rate of orbital AML and the braking law from the upper boundary of the orbital AM distribution.  Usage of small number statistics and saturation in activity-rotation relation may be responsible for the weakness among those studies. Perhaps, the upper boundary does not represent all systems, and the mass loss and the AML may be age-dependent quantities. In order to better understand the orbital period evolution of detached active binaries, systems with different ages should be considered. Although the ages of binaries can be estimated only by detailed evolutionary isochrones, the sub-group ages can be found kinematically. Recently, after having more accurate data on a large sample of detached CABs, Karata\\c{s} et al.\\ (2004) became successful in breaking up the sample into kinematically distinct sub-samples. Karata\\c{s} et al.\\ (2004) initially divided the whole sample of 237 systems into two groups: the possible moving group (MG) members (95) and the older field binaries (142). A comparison of the total mass, the orbital period, the mass ratio and the orbital eccentricity of these two groups revealed clear observational evidence that detached CABs lose mass and angular momentum so that their orbits circularize and shrink. Related orbital period decreases, unfortunately, are not detectable on commonly used O--C diagrams formed by the eclipse minimum times. This is due to: 1) very short time span covered by existing O--C data (at most 100 years) in comparison with durations implied by the predicted kinematical ages which are at the order of $10^9$ years; 2) large scatter of unevenly distributed data (especially visual and photographic observations) in the present O--C diagrams; 3) existence of complicated larger amplitude short time scale fluctuations caused by many different effects (cf Kreiner, Kim \\& Nha\\ 2001; Demircan\\ 2000, 2002). Being independent of the physical cause, the O--C diagrams are commonly used in the study of orbital period changes of binaries in general. At present, mean minimum time deviations as small as 10 seconds for binaries with sharp eclipses are detectable. Nevertheless, the minimum detectable period variation on O--C diagrams depends on the time span of observations over the actual variations.  This classical approach with O--C diagrams, thus, are not suitable to obtain slower secular orbital decreases implied by the wind driven mass loss of our sample of detached CABs. Consequently, our aim in the present work is to further investigate the mass-loss, the angular momentum loss, orbital period decrease and then determine the rates of change of those parameters statistically from the new and more accurate data (absolute dimensions and kinematical data) of Karata\\c{s} et al.\\ (2004). The present approach, apparently, is more advantageous in detecting changes at the order of evolutionary time scales, where O--C diagrams becomes insufficient, and also more practical to give no distinction weather the involved binaries are eclipsing or not. ",
        "conclusions": "The well known spin-down of single stars due to AML requires orbital shrinkage (period decrease) in close binary systems. For this process to be effective, the spin-orbit coupling $(P_{s} \\cong P_{orb})$ consequence of tidal locking was already suggested to be a necessary condition (Guinan \\& Bradstreet\\ 1988; van't Veer \\& Maceroni\\ 1988, 1989; Stepien\\ 1995). Thus, long period binaries with no tidal locking are not expected to evolve into shorter period systems since tidal locking would be ineffective in transferring orbital angular momentum to the spinning components where magnetic braking operates.  With a shrinking orbit, the more massive component of a short period binary may fill its Roche lobe faster and start transferring mass to the other component, while the system is still evolving towards shorter periods under the wind-driven mass loss and spin-orbit coupling mechanism. Only after Roche lobe overflow starts, AM evolution of the binary becomes dominated by the mass transfer. Until the mass-ratio reversal, the orbital period should be decreasing. But during the second stage, after the mass ratio reversal, the orbital period is expected to increase, that is, opposite to shrinking, and the orbit starts to enlarge.  The sample of this study contains detached CABs with orbital period greater than 0.479 days (XY UMa). A direct period limit to tidal locking is not known. However, to obtain undistorted statistical results, lower limits are applied (2.4 days by van't Veer \\& Maceroni\\ 1988, and 10 days by Demircan\\ 1999). This limit appears to be the function of the total mass, the orbital period, as well as age of the system (see, Tassoul\\ 2000). By comparing orbital and rotation periods in Table 1 and Table 2, we estimated that it is not less then about $\\sim 70$ days in the field CABs and around 10 days in the MG CABs. However, the MG CABs are not fully effective to give age dependent variations in $J$, $M$ and $P$ and there are only two systems with $P>70$ days in the field CABs (see Table 2). Thus, we just ignored the tidal locking limit to the orbital period of field CABs.  Nevertheless, the tidal or the magnetic locking, in principle, does not involve a mass loss directly. But, it is a mechanism which transfers OAM to the components, where the AM is lost by magnetically driven stellar winds. The magnetic field lines, especially the ones which are perpendicular to the stars surface and could reach up to the Alfven radius, enforces plasma to co-rotate. As long as the mass is lost from the Alfven radius, not from the surface directly as in the case of isotropic winds, angular momentum loss per particle appears to be amplified. Since OAM loss becomes associated with the mass loss and there is a mechanism which amplifies angular momentum loss per particle, the condition derived from eq. (12) will be valid.  The amplification factor could differ from one system to another. However the general trend, as implied by the plots in Fig. 2, indicates that the average amplification ($\\bar A$) must be bigger than 5/3 for the majority (perhaps all) of CABs so that the decrease of $P$ together with the decrease of $J$ and $M$ by age is observed. Otherwise, if $\\bar A<5/3$, the orbital periods would have increased despite the mass and OAM loss. The average amplification factor for the present sample could be determined as  \\begin{eqnarray} {\\bar A} = {{dJ \\over dt}{1\\over J} \\over {dM \\over dt}{1\\over M}} = {{\\alpha _{J} \\over \\alpha _{M}} = 2.68}. \\end{eqnarray}  If such an amplification mechanism operates among the CABs, one must expect the rate coefficients $\\alpha _{J}$, $\\alpha _{M}$ and  $\\alpha _{P}$ must hold a relation  \\begin{eqnarray} 3{\\alpha _{J}} = 5{\\alpha _{M}} + {\\alpha _{P}} \\end{eqnarray} according to eq. (6). Replacing $\\alpha _{J}$ by $\\bar A \\alpha _{M}$ according to eq. (23), eq. (24) reduces to  \\begin{eqnarray} {(3\\bar A -5)}{\\alpha _{M}}= {\\alpha _{P}}, \\end{eqnarray} which implies that the decreasing rate of orbital periods can also be predicted from the amplification parameter and the mass loss rate if they are known.  The mean amplification factor $\\bar A$ = 2.68, and the mean mass loss rate $\\alpha_{M}=1.30\\times 10^{-10}yr^{-1}$ for the present sample of CABs require the mean rate of orbital period decrease to be $3.95 \\times 10^{-10}yr^{-1}$ according to above relation. It is indeed interesting that this computed value agrees with the value $3.96 \\times 10^{-10}yr^{-1}$ which was found independently from the regression analysis of linear line fitting to $P$ data in Table 3. So we can conclude that the present data confirm the orbital period decrease as a cause of the mass loss from CAB systems with a mechanism which is sufficient to draw 2.68 times more OAM than isotropic mass loss from the component surfaces.  As for the direct confirmation of the predicted continuous orbital period decreases by O--C diagrams, it is known in general that the changes as small as about $10^{-9}$ days in systems with periods around one day would be observable over a timespan of a century. However, in the case of CABs, because of the large scatter and fluctuations due to magnetic activity or due to third body with light-time effect in the O--C diagrams, it may not possible to detect our prediction of continuous orbital period decreases which operate in a much larger time-scale which is comparable to main-sequence evolutionary times."
    },
    "0512/astro-ph0512398_arXiv.txt": {
        "abstract": "A new type of cosmic-ray experiment is under construction in the Pyh\\\"asalmi mine in the underground laboratory of the University of Oulu, Finland. It aims to study the composition of cosmic rays at and above the \\textit{knee} region. The experiment, called EMMA, will cover approximately 150 m$^2$ of detector area. The array is capable of measuring the multiplicity and the lateral distribution of underground muons, and the arrival direction of the air shower. The full-size detector is expected to run by the end of 2007. ",
        "introduction": "%  Due to a slight change observed in the cosmic-ray energy spectrum in the energy interval of 10$^{15}$ -- 10$^{16}$ eV, it is believed that the origin, modification in the chemical composition, acceleration mechanism or propagation of cosmic rays (or a combination of these) changes. Up to this energy, so called \\textit{knee} region, most cosmic rays are supposed to be produced inside the galaxy, and are also confined by the galactic magnetic field.  At these high energies the source cannot be observed directly and the cosmic-ray composition is used as a tool to investigate the origin of the cosmic radiation. The direct composition measurements are no longer practical at or above the knee and the method is solely based on the measurement of extensive air showers, i.e. the secondary particles created in the atmosphere and detected by large arrays on the ground.  The origin of the knee has been one of the fundamental problems of cosmic-ray physics, and it has been discussed for decades. Several models have been presented predicting different composition at the knee energies, and could only be identified by the experimental evidence on the composition. Some new experimental efforts have been devoted to the study of cosmic rays in recent years. These experiments are based, for example, on multi-parameter measurement of extensive air-showers, on shower maximum measurement by \\v Cerenkov or fluorescence detectors and on underground multimuon measurements (see, for example, Ref. \\cite{Hor05} and references therein). Their conclusions, however, have so far been diverse, implying the need for further studies, especially using different approaches. The results are also known to be strongly model dependent.  EMMA (Experiment with MultiMuon Array) uses a different approach. It is not the first underground cosmic-ray experiment (see, for example, Refs. \\cite{Ceb90,Kas97,Ava03,Gru03,EAS04}), but it differs significantly from previous underground experiments with its ability to measure the lateral distribution function of underground muons. In EMMA the composition analysis is based on the lateral distribution of high-energy muons and on their multiplicity. The muons detected by EMMA are generated in the upper part of the air shower close to the primary interaction. ",
        "conclusions": "A new underground cosmic ray experiment EMMA is under construction and it is expected to start recording data in the full scale by the end of 2007. With a partial-size array the data recording can be started by the middle of 2006. The analysis of simulated air showers shows that the primary cosmic-ray composition could be resolved (with a two-component model) at and above the knee energies. A possibility for the model-independent way of the determination of the muon lateral distribution would allow to improve high-energy interaction models. Due to new method used, the EMMA experiment could provide comprehensive (and perhaps new) information on the composition of the cosmic rays at the knee region within the next few years.   \\ack  The support from the Magnus Ehrnrooth Foundation, the Jenny and Antti Wihuri Foundation, the Finnish Academy of Science and Letters (V\\\"ais\\\"al\\\"a Foundation), and the Finnish Cultural Foundation is acknowledged. The work is funded by the European Union Regional Development Fund and it is also supported by the Academy of Finland (projects 108991, 7108875 and 7106570)."
    },
    "0512/astro-ph0512351_arXiv.txt": {
        "abstract": "The L1551 molecular cloud, unlike most of the Taurus Molecular Complex, is undergoing a long and sustained period of relatively high efficiency star formation. It contains two small clusters of Class 0 and I protostars, as well as a halo of more evolved Class II and III YSOs, indicating a current and at least one past burst of star formation.  We present here new, sensitive maps of 850 and 450 \\mum\\  dust emission covering most of the L1551 cloud, new CO J=2-1 data of the molecular cloud, and a new, deep, optical image of \\sii\\ emission (6730\\AA).  We have detected all of the previously known Class 0 and I YSOs in L1551, and no new ones.  Compact sub-millimetre emitters are concentrated in two sub-clusters:  IRS5 and L1551NE, and the HL~Tauri group.  Both stellar groups show significant extended emission and outflow/jet activity.  A jet, terminating at HH~265 and with a very weak associated molecular outflow, may originate from LkH$\\alpha$~358, or from a binary companion to another member of the HL~Tauri group. Several Herbig Haro objects associated with IRS5/NE were clearly detected in the sub-mm, as were faint ridges of emission tracing outflow cavity walls.  We confirm a large-scale molecular outflow originating from NE parallel to that from IRS5, and suggest that the ``hollow shell'' morphology is more likely due to two interacting outflows.  The origin of the E-W flow east of HH~102 is undetermined.  We confirm the presence of a prestellar core (L1551-MC) of mass 2-3 M$_{\\odot}$ north-west of IRS5.  The next generation cluster may be forming in this core.  The L1551 cloud appears cometary in morphology, and appears to be illuminated and eroded from the direction of Orion, perhaps explaining the multiple episodes of star formation in this cloud.  The full paper (including figures) can be downloaded at http://www.jach.hawaii.edu/$\\sim$gms/l1551/l1551-apj641.pdf, or viewed at http://www.jach.hawaii.edu/$\\sim$gms/l1551/ . ",
        "introduction": "In stark contrast to most of the Taurus-Auriga star formation complex, which is characterized by very low efficiency formation of isolated stars (Palla \\& Stahler 2002), the L1551 molecular cloud has sustained a long and continuing period of star formation of relatively high efficiency. The cloud is surrounded by a halo of at least 30 classical and weak T Tauri stars (cTTs, wTTs) and proto-brown dwarfs (Cudworth \\& Herbig 1979; Jones \\& Herbig 1979; Feigelson \\& DeCampli 1981; Feigelson \\& Kriss 1983; Wood \\etal\\ 1984; Feigelson \\etal\\ 1987; Gomez \\etal\\ 1992; Carkner \\etal\\ 1996; Brice{\\~n}o, {\\em et al.} 1998; Brice{\\~n}o, Stauffer \\& Kirkpatrick 2002; Favata \\etal\\ 2003; G{\\aa}lfalk \\etal\\ 2004), indicating star formation activity over the past few million years. Embedded within the cloud are two known groups of active star formation, one comprised of IRS5 and L1551-NE (hereafter NE), and the other the HL~Tau group, located $\\sim$5\\arcmin\\ north of IRS5 and consisting of HL~Tau, XZ~Tau, LkH$\\alpha$ 358, and HH~30*.   (In the following discussion, HH~30 refers to the Herbig-Haro object(s), while HH~30* (also known as V1213 Tau) refers to the driving star.) The question remains, will this sustained star formation continue beyond the present epoch, and what is the reason for this sustained activity?  LDN 1551 (Lynds 1962), a modest-sized ($M \\approx$ 50 \\msol, diameter $D \\sim$ 1 pc) (Moriarty-Schieven \\& Snell 1988) dark cloud, lies directly behind the Hyades star cluster in Taurus. Its distance was recently constrained using Hipparchos observations by Bertout, Robichon \\& Arenou (1999), who determined that the average distance to the Taurus-Auriga complex was 139$^{+10}_{-9}$ pc, but found the distance to T Tauri (which is much nearer to L1551 than the rest of the Tau-Aur complex) to be 168$^{+12}_{-28}$ pc. Extended optical emission from L1551 was first discovered on photographic plates and cataloged as the diffuse \\Hii\\ region S239 (Sharpless 1959). This emission was later found to exhibit characteristics of a Herbig-Haro object (shock excited optical emission powered by outflows from young stars) and re-categorized as HH~102 (Strom, Grasdalen, \\& Strom 1974). Three bright and compact HH objects, HH~28, HH~29, and HH~30, were also found in the region (Herbig 1974). They were at first mis-identified as high velocity stars due to their high proper motions (Luyten 1963, 1971). Recent deep images of the L1551 cloud have led to the discovery of additional Herbig-Haro objects (Graham \\& Heyer 1990; L\\'{o}pez \\etal\\ 1995; Riera \\etal\\ 1995).   New aspects of the L1551 star forming region continue to be uncovered as deep optical images, high angular resolution radio continuum maps, and other studies are obtained.  High angular resolution radio continuum maps have shown that IRS~5, an {\\em FU Ori} type Class I protostar (Sandell \\& Weintraub 2001), is a binary (Rodriguez \\etal\\ 1986, 2003a) with a separation of about 50 AU (Looney, Mundy \\& Welch 1997; Rodriguez \\etal\\ 1998), total mass $\\sim$1.2$_{\\odot}$, and period $\\sim$260yr (Rodr{\\'i}guez \\etal\\ 2003a). Each star in this system drives a distinct stellar jet (Snell \\etal\\ 1985; Fridlund \\& Liseau 1998; Rodr{\\'i}guez \\etal\\ 2003b) with orientations that are misaligned with the CO outflow lobes, discovered by Snell, Loren \\& Plambeck (1980), and thought to be powered by IRS~5.  Deep optical imaging, proper motion studies, and spectroscopy have provided evidence that the nearby YSO L1551~NE, a Class 0 protostar (Moriarty-Schieven, Butner \\& Wannier 1995) discovered by Emerson \\etal\\ (1984), located about 1\\arcmin\\ to the east of IRS~5, may power the bright Herbig-Haro object HH~29 and possibly HH~28 (Devine, Reipurth, \\& Bally 2000; Hartigan et al. 2000) along with a highly collimated 1.5 parsec-long chain of fainter Herbig-Haro objects that are superimposed on the main CO lobes of the L1551 outflow complex.  The NE source is also a multiple system (Rodriguez, Anglada \\& Raga 1995; Moriarty-Schieven \\etal\\ 2000), and possesses a molecular outflow (Moriarty-Schieven, Butner \\& Wannier 1995), although the outflow is aligned with and hence confused by that of IRS5.  Associated with IRS5/NE is another low-velocity molecular outflow, the L1551 E-W flow (Moriarty-Schieven \\& Wannier 1991; Pound \\& Bally 1991), which extends $\\sim$30\\arcmin nearly due west of IRS5/NE, but whose origin is unknown.  Near HL~Tauri is a compact group of four or five YSOs,  of which all but one have outflows and/or jets (Mundt, Buhrke \\& Ray 1988). Deep images obtained through narrow-band interference filters transmitting the light of \\Ha\\ and \\sii\\ show that the nearby sources HL and XZ~Tau and the source of the HH~30 stellar jet power optically visible Herbig-Haro flows which also criss-cross the IRS~5 outflow lobes (Devine, Reipurth, \\& Bally 1999).  Several molecular outflows have also been detected toward this group (Monin, Pudritz \\& Lazareff 1996).      Most of the surveys for dense cores in the L1551 cloud have concentrated on the IRS5 region (e.g. Moriarty-Schieven \\etal\\ 1987 for CS, Chandler \\& Richer 2000 for dust continuum), while Onishi \\etal\\ (2002) neglected L1551 in their systematic H$^{13}$CO$^{+}$ survey of cores in Taurus.  In this work we present 850 \\mum\\  and 450 \\mum\\  dust continuum images of a region covering most of the L1551 cloud.  In addition, we present new CO J=2-1 maps of the cloud, in order to search for new outflows and disentangle those already known.  The full paper (including figures) can be downloaded at  \\noindent http://www.jach.hawaii.edu/$\\sim$gms/l1551/l1551-apj641.pdf,  \\noindent or viewed at http://www.jach.hawaii.edu/$\\sim$gms/l1551/ . ",
        "conclusions": "In this paper, we present new sub-millimeter dust continuum maps and J=2-1 CO images of the L1551 cloud in the Taurus cloud complex.   These data are combined with deep visual wavelength images and analyzed in the context of the extensive literature on this nearby star forming cloud.  All previously known Class 0 and Class I YSOs in L1551 are detected in the 850 $\\mu$m sub-mm continuum images. Furthermore, no new compact sub-mm sources were found. Thus, the inventory of highly embedded young stellar objects appears to be complete.  However, it is possible that the known Class 0/I protostars contain additional members which remain unresolved in our data.  The compact sub-mm emitters are concentrated in two sub-clusters.  The binary system IRS5, the most luminous and massive YSO in L1551,  and the L1551-NE multiple system, are located near the southeastern end of the cloud.  Both members of IRS5 and one member of the NE system power nearly parallel jets which appear to be collectively responsible for powering the main CO outflow in this cloud.  While the IRS5 binary and its HH~154 jet clearly power HH~102 and the northern rim of the CO outflow complex, NE is responsible for the highly collimated HH~454 jet and \\sii\\ knots, the bright HH object HH~29, possibly HH~28, the eastern portion of HH~262, and possibly the distant bow shock HH~286. Thus, one member of the NE system is responsible for the southern portion of the main CO outflow lobes.  The driving source of the large redshifted east-west CO lobe originating near HH~102 remains unclear. Our SCUBA observations did not find any additional YSO candidates along the axis of this outflow lobe.  However, the orientation of this limb-brightened low-velocity outflow lobe suggests that the driving source may one member of the L1551-NE multiple star syst.  The second sub-cluster of sub-mm emitters consists of the sources clustered around HL~Tau (HH~30*, XZ~Tau, and LkH$\\alpha$~358).  All four of these YSOs are associated with 850 $\\mu$m emission with HL~Tau being by far the brightest.  LkH$\\alpha$~358 (or perhaps the binary companion of XZ~Tau or a currently unknown binary companion of HL~Tau) may be the origin of a jet terminating at HH~265.  A very weak molecular outflow was detected associated with the end of this jet.  In addition to the compact sources, the sub-mm continuum emission exhibits diffuse components.  Ridges of 850 $\\mu$m emission extends from IRS5 to HH~102, from IRS5 towards the axis of the southwestern CO lobe, from IRS5 towards NE, and from from NE towards HH~262.  The brightest parts of both HH~102 and HH~262 are detected at 850 $\\mu$m. Although CO J = 3-2 emission may make a small contribution to this emission, most of the flux is likely to be produced by warm dust entrained by outflows.  A complex halo of dust emission also surrounds the HL~Tau region.  This extended emission may compromise the use of this YSO as a secondary calibrator for sub-mm continuum observations.  The L1551 cloud is cometary, indicating progressive erosion from the general direction of Orion, located to the southeast. Deep wide-field visual wavelength images and CO maps show that the L1551 cloud has a cometary morphology consisting of a dense head near the IRS5 / NE and HL~Tau region, and a diffuse tail extending more than a degree towards the northwest.  The densest part of this diffuse tail can be clearly seen in our \\sii\\ image (Figure 4).  Diffuse 850 $\\mu$m dust emission is produced by this region, dubbed L1551-MC.  This object is likely to be a pre-stellar core showing signs of undergoing the first phases of gravitational collapse (Swift, Welch, \\& Di Francesco 2005).  We do not see any signs of fragmentation of the clump.  Assuming a dust temperature of 9 K, our sub-mm continuum flux implies a total mass of order 3 M$_{\\odot}$.  The wide-field \\sii\\ image (Figure 4) shows that the southeastern rim of the L1551 cloud is illuminated from the southeast. The CO maps show that the cloud has a sharp edge in this direction. Many background galaxies are visible beyond the southeastern edge of the cloud, indicating that this portion of the cloud has been dispersed.  The locations of the older Class II and Class III YSOs born from the L1551 cloud indicate that stars formed relatively recently southeast of the current cloud edge (Figure 4). The cometary shape of the L1551 cloud, the projected distribution of older YSOs, the faint glow along the southeastern rim of L1551, and the location of L1551-MC northwest of the recently formed Class 0/I YSOs indicate that the L1551 cloud has been irradiated and eroded from the general direction of Orion.  The currently active sites of star formation in Orion are located between 380 and 470 pc from the Sun.  However, the older sub-groups of the Orion OB association are located closer with some members of the 1a and 1b subgroup as near as 300 pc, about 150 pc from L1551 (de Zeeuw et al 1999).  The closest portion of the Orion-Eridanus supershell, the H$\\alpha$ feature known as the Eridanus Loop (Reynolds, \\& Ogden 1979; Boumis et al. 2001), has been estimated to be only 160 pc from the Sun  (Burrows \\& Guo 1996), which places it within 50 pc of L1551.  Portions of the 21 cm HI supershell driven  by the Orion OB association and 100 $\\mu$m dust emission traced by the IRAS satellite are located well outside the H$\\alpha$ loops illuminated by Orion's massive stars (e.g. Brown, Hartnamm, \\& Burton 1995; Heiles, Haffner, \\& Reynolds 1999).  The two well-known massive stars, Betelgeuse and Rigel, are located in the closer-parts of the Orion-Eridanus supershell and southeast of L1551.  Betelgeuse, is about 130 pc from the Sun and has an 18 \\kms\\ proper motion towards PA $\\sim$ 68\\arcdeg . Rigel is located 240 pc from the Sun and illuminates a small complex of star-forming molecular clouds associated with IC 2118 in the southwester portion of Orion (Kun et al. 2004).  The proper motion of Betelgeuse (7.6 milliarcsec /year) would place it southeast of L1551 several million years ago. Thus, both of these massive stars could have been within 100 pc of the L1551 cloud within the last several million years and may have contributed to the photo-erosion of this cloud.  In summary, distribution of YSOs, the cometary shape, the external illumination, and location of the pre-stellar core L1551-MC indicate that the evolution of the L1551 cloud has been influenced by massive stars and photo-erosion from the general direction of Orion.  The closest parts of the Orion OB associations, and massive stars in the direction of Orion may be responsible for the evolution of the L1551 cloud."
    },
    "0512/astro-ph0512484_arXiv.txt": {
        "abstract": "A key science goal of upcoming dark energy surveys is to seek time evolution of the dark energy. This problem is one of {\\em model selection}, where the aim is to differentiate between cosmological models with different numbers of parameters.  However, the power of these surveys is traditionally assessed by estimating their ability to constrain parameters, which is a different statistical problem.  In this paper we use Bayesian model selection techniques, specifically forecasting of the Bayes factors, to compare the abilities of different proposed surveys in discovering dark energy evolution. We consider six experiments --- supernova luminosity measurements by the Supernova Legacy Survey, SNAP, JEDI, and ALPACA, and baryon acoustic oscillation measurements by WFMOS and JEDI --- and use Bayes factor plots to compare their statistical constraining power.  The concept of Bayes factor forecasting has much broader applicability than dark energy surveys. ",
        "introduction": "Uncovering the nature of dark energy in the Universe is perhaps the greatest challenge facing cosmologists in coming years. In recent months many proposed experiments to probe dark energy have been defined, especially in response to a call for white papers by the Dark Energy Task Force set up jointly in the US by the NSF, NASA and DOE. These propose a variety of techniques to constrain dark energy parameters, including the luminosity distance--redshift relation of Type Ia supernovae (SNe-Ia), the angular-diameter distance--redshift and expansion rate--redshift relations measured by baryon acoustic oscillations, and use of weak gravitational lensing to probe the growth rate of structures.  Following on from heritage of CMB anisotropy studies, the standard tool used to illustrate the power of a given instrument or survey is a plot of the projected parameter errors around one or more fiducial models, estimated using a Fisher information matrix approach or likelihood analysis of Monte Carlo simulated data (Knox 1995; Jungman et al.~1996; Zaldarriaga, Spergel \\& Seljak 1997; Bond, Efstathiou \\& Tegmark 1997; Efstathiou \\& Bond 1999).  Typically, a projection of the parameter uncertainties onto a two-parameter equation-of-state model for dark energy is deployed, showing how tightly parameters are expected to be constrained around, for instance, the cosmological constant model. The implication is intended to be that if the true values lie outside those error ellipses, then the survey will be able to exclude the cosmological constant model.  However, the principal goal of such surveys is usually identified as being the discovery of dark energy evolution. This is not a parameter estimation question, but rather one of {\\em model selection} (Jeffreys 1961; MacKay 2003; Gregory 2005), where one seeks to compare cosmological models with different numbers of variable parameters. Within the framework of Bayesian inference, the statistical machinery to make such comparisons exists, and is based around statistics known as the Bayesian evidence and the Bayes factor. The Bayes factor has the literal interpretation of measuring the change in relative probabilities of two models in light of observational data, updating the prior relative model probabilities to the posterior relative model probabilities.  In this paper we use Bayesian model selection tools to assess the power of different proposed experiments. Our method is related to the Expected Posterior Odds (ExPO) forecasting recently developed by Trotta (2005).  The main difference is that he takes the present observational constraints on the extended model, and seeks to estimate the fraction of that parameter space within which that model can be distinguished from a simpler embedded model. By contrast, we take a theoretically-motivated view of the parameter space of interest, and seek the locations within that parameter space corresponding to dark energy models which are distinguishable from a cosmological constant by a given experiment.  We also differ computationally, in that as well as using approximate techniques, we use the nested sampling algorithm of Skilling (2004), as implemented by Mukherjee, Parkinson \\& Liddle (2006), to compute the evidences accurately numerically.  The paper is organized as follows. In Section~\\ref{bsf} we introduce model selection in the Bayesian framework.  Section~\\ref{surveys} describes the dark energy surveys we make model selection forecasts for, and Section~\\ref{results} presents the results.  We conclude in Section~\\ref{Conclusions}. We consider some additional technical details and review the standard parameter forecast procedure in two Appendices. ",
        "conclusions": "\\label{Conclusions}  In this paper we have introduced the Bayes factor plot as a tool for assessing the power of upcoming experiments. It offers a full implementation of Bayesian model selection as a forecasting tool. As compared to the traditional parameter error forecasting technique, it offers a number of advantages enumerated in Section~\\ref{ss:list}. Amongst those, perhaps the most important are that observational data is simulated at each point in the plane, rather than at a small number of fiducial models, and that the Bayes factor plot properly captures the experimental motivation as being one of model selection rather than parameter estimation.  As a specific example, we have used the Bayes factor plots to examine a number of proposed dark energy surveys, concentrating on their ability to distinguish between the $\\Lambda$CDM model and a two-parameter dark energy model. Figure~\\ref{fig:all} indicates the region of parameter space outside which the true model has to lie, in order for the experiment to have sufficient statistical power to exclude $\\Lambda$CDM using model selection statistics.  An important caveat is that our plots do not show the effects of systematics, which are likely to be the dominant uncertainty for many of the experiments. This drawback is shared by parameter error forecasts, and it is more or less the nature of systematic uncertainties that they cannot be usefully modelled in advance of actual observational data being obtained. In judging the true merit of an experimental proposal, it is therefore essential to judge how well structured it is for optimal removal of systematics, as well as looking at its raw statistical power.  While we have focussed on dark energy as a specific application, the concept of the Bayes factor plot has much broader applicability, and is suitable for deployment in a wide range of cosmological contexts."
    },
    "0512/astro-ph0512437_arXiv.txt": {
        "abstract": "The colours of high redshift Type II QSOs are synthesized from observations of moderate redshift systems. It is shown that Type II QSOs are comparable to starbursts at matching the colours of $z_{850}$-dropouts and $i_{775}$-drops in the Hubble Ultra Deep Field, and more naturally account for the bluest objects detected. Type II QSOs may also account for some of the $i_{775}$-drops detected in the GOODS fields. It is shown that by combining imaging data from the {\\it Hubble Space Telescope} and the {\\it James Webb Space Telescope}, it will be possible to clearly separate Type II QSOs from Type I QSOs and starbursts based on their colours. Similarly, it is shown that the UKIDSS $ZYJ$ filters may be used to discriminate high redshift Type II QSOs from other objects. If Type II QSOs are prevalent at high redshifts, then AGN may be major contributors to the re-ionization of the Intergalactic Medium. ",
        "introduction": "\\label{sec:introduction}  The band-dropout method has proven an impressive means for discovering high redshift objects. Pioneered by Guhathakurta \\etal (1990), Bithell (1991), and Steidel \\& Hamilton (1992, 1993), the method exploits the downward step in flux shortward of the Lyman limit for discovering starburst galaxies. Absorption by intervening hydrogen in the Intergalactic Medium (IGM) will similarly produce extremely red optical colours for objects at $z\\gsim3$, a characteristic signature of a high redshift system (Bithell 1991; Madau 1995). Following Steidel \\& Hamilton's discovery of a population of Lyman break galaxies at $z\\approx3$, the method has evolved into a standard tool for identifying even higher redshift objects in recent very deep surveys like the Hubble Deep Field and the Ultra-Deep Field. While follow-up spectroscopy has demonstrated that most of the objects are galaxies, a few dozen reveal the emission line signatures of Active Galactic Nuclei (AGN) (Steidel \\etal 2002), about half of which have the narrow lines characteristic of Type~II systems. These AGN have made possible the first measurement of the faint end of the luminosity function of Quasi-Stellar Objects (QSOs) at $z\\approx3$ (Hunt \\etal 2004), critical for assessing the poorly constrained contribution of QSOs to the UV ionizing background at this epoch (Meiksin 2005a).  Since then, band-dropout objects have been discovered in a host of surveys, most recently pushing to $z\\approx6$ (Stanway, Bunker \\& McMahon 2003; Bouwens \\etal 2004; Bunker \\etal 2004; Yan \\& Windhorst 2004; Giavalisco \\etal 2004). The objects have generally been modelled as young star-forming galaxies and used to infer the cosmic star formation rate of the universe and its evolution. They have also been used to assess the contribution of galaxies to the budget of ionizing photons required to re-ionize the universe, concluding their numbers are either too few (Bunker \\etal 2004) or easily adequate (Stiavelli \\etal 2004), depending on model assumptions.  In this paper, it is suggested that a portion of the band-dropout objects may be Type II QSOs.\\footnote{There is not universal consensus on the definition of a Type II QSO. In this paper, a definition similar to that of Zakamska \\etal (2003) is assumed, that the restframe FWHM of hydrogen lines be less than 2000\\kms. An alternative x-ray motivated definition is based on x-ray spectral evidence for a large obscuring \\HI column density local to the active nucleus, for example, Gandhi \\etal (2004).} The prospects of detecting Type II QSOs in deep surveys has not previously been explored. These objects have been of considerable recent interest because they are predicted in unification models of AGN (Antonucci 1993). Discovering examples at higher redshifts would help to further elucidate their properties and their connection to Type I QSOs. As sources of energetic photons, they are also candidate sources for high energy background radiations. Obscured QSOs have been postulated as sources of the hard x-ray background (Madau, Ghisellini \\& Fabian 1994), and Type II QSOs are natural candidates for such sources. While Type I QSOs may dominate the UV ionizing metagalactic background at $z\\lsim3.5$, their numbers appear inadequate at higher redshifts (eg, Meiksin 2005a). Although galaxies are possible sources, the abundance of high redshift Type II QSOs is far too uncertain to rule out an ionizing background dominated by AGN sources. If their numbers are sufficiently great at high redshifts, Type II QSOs may also have contributed substantially to the re-ionization of the IGM. Indeed, if the IGM were re-ionized at $z=6-8$, consistency with measurements of the optical depth of the IGM at $z<6$ favour hard ionizing sources like Pop~III stars and AGN (Meiksin 2005a).  The Sloan Digital Sky Survey (SDSS) has discovered nearly 300 Type II QSOs over the redshift range $0.3<z<0.83$, with the rest wavelengths of the measured spectra restricted to being longward of about 2100~\\AA\\ (Zakamska \\etal 2003). These elusive narrow-lined objects have proven even more difficult to discover at higher redshifts, and their UV spectral properties are consequently still poorly understood. To date, only four Type II QSOs are known to exist at $z>1.5$ with optical spectra showing narrow emission lines, two discovered through observations with the {\\it Chandra X-Ray Observatory} (Norman \\etal 2002; Mainieri \\etal 2005), one discovered as an optical band-dropout (Stern \\etal 2002), and one discovered through optical spectroscopy (Jarvis \\etal 2005). A fifth candidate was found by Dawson et al (2001), but with incomplete emission line results reported. The discovery of high redshift counterparts would add substantially to our knowledge of the nature and origin of these systems. ",
        "conclusions": "Spectra for Type II QSOs are synthesized based on observed spectra and photometric measurements, and colours predicted at higher redshifts. The possibility that some Type II QSOs have high equivalent widths suggests unusual colours may be expected for these objects, distinct from Type I QSOs and starbursts. It is shown that the colours of the $z_{850}$-dropouts discovered in the UDF by Bouwens \\etal are matched by Type II QSOs, as are some of the $i_{775}$-drops found in the UDF and GOODS. In the case of $i_{775}$-drops, the large but measurable $(i_{775}-z_{850})_{\\rm AB}$ colour is achievable in only narrow redshift windows, one at $z\\gsim6$ when the objects abruptly become $i_{775}$-dropouts, and a second near $z=4.6$ for which $(i_{775}-z_{850})_{\\rm AB}>1.3$.  It may seem unlikely to find such objects given their rarity at lower redshifts. It should, however, be noted that even at low redshifts, Type II AGN appear to be a factor of a few to as much as an order of magnitude more abundant than Type I AGNs at low luminosities (Mart\\'inez-Sansigre \\etal 2005; Simpson 2005). The luminosity function of low-luminosity Type I QSOs at high redshifts is unknown. An estimate is made by Meiksin (2005a), under the assumptions of either pure luminosity evolution (PLE) or pure density evolution (PDE), based on the low luminosity QSO counts of Hunt \\etal (2004) at $z\\approx3$ and the bright QSOs detected at high redshift ($3.6\\lsim z\\lsim6$) by the SDSS (Fan \\etal 2001, 2004). Using the maximum likelihood PLE model of Meiksin (2005a) (as given in his Table~1), the predicted number of Type I QSOs at $z>5.7$ in the $11.5\\,{\\rm arcmin}^2$ of the UDF with $z^\\prime<28$ is about 0.07 (0.002 in the PDE model, as the UDF should probe well below the knee in the luminosity function in this model). Thus the number of detectable Type II QSOs in the UDF is expected to lie in the range $0.01-1$. Within the $165\\,{\\rm arcmin}^2$ of the GOODS-S ACS field, the number of $z^\\prime<27$ Type I QSOs at $z>5.7$ is predicted in the PLE model to be 0.6 (0.03 in the PDE model), and a number per unit redshift of $dN/dz\\approx1$ for $z\\approx4.7$ (0.1 in the PDE model). Thus detecting at least one Type II QSO is possible, although several may be unlikely. But this is one reason why searching for them is of such interest. For instance, if most galaxies went through an AGN phase early in their histories, then, in AGN unification scenarios, most would appear as dim Type II objects. If they went through the AGN phase at $z\\gsim6$, then their numbers may be larger than the above estimates, and a few to several may be detectable in the UDF and GOODS fields, sufficient to provide a subtantial fraction of the photons required to re-ionize the IGM (Meiksin 2005a).  Although the UKIDSS Ultra Deep Survey is a somewhat shallower survey, the relatively large survey area of $0.78\\,{\\rm deg}^2$ yields a detectable number of Type I QSOs:\\ about six are predicted at $z>5.7$ for $J<26$ by the PLE model, and $dN/dz\\approx13$ at $z\\approx4.7$. Several Type II QSOs may thus be detectable at these redshifts. While the predicted numbers of Type I QSOs are substantially reduced under the PDE model, about 0.4 at $z>5.7$ and $dN/dz\\approx1.6$ at $z\\approx4.7$, the Ultra Deep Survey should prove an effective means for determining which of these models more closely describes the actual evolution of QSOs.  Without follow-up spectroscopy, it is difficult to distinguish a starburst from an AGN. The images of the band-dropout objects tend to be compact (eg, Stanway \\etal 2004a), as would be expected for \\HII regions, but also for an AGN, so images may not readily be used to distinguish starbursts from AGNs. Extended emission would not necessary preclude an AGN nature either, both because the emitting regions extend from scales of several to a few tens of kiloparsecs in size (Hines \\etal 1999) and because some may be embedded in starbursts and so have multiple bright emitting regions; both AGN and starbursts are generally believed to be triggered by merger activity. Even with spectroscopy, high equivalent width Type II QSOs at $z\\approx4.6$ could still masquerade as starbursts, with \\CIV emission being mistaken for \\Lya at $z\\approx6.0$. Strong \\Ha emission, falling at $\\lambda=3.6\\mu {\\rm m}$, could also mimic the 4000\\AA~ break of a starburst at $z\\approx6.0$. Even if \\Lya emission is correctly identified, the absence of \\NV emission does not preclude a Type II QSO, as it would a Type I QSO, since the \\Lya to \\NV ratio in Type II QSOs is observed sometimes to be very large.  The absence of x-ray emission also may not necessarily exclude an AGN nature. The expected x-ray flux appears to cover a broad range for a given UV flux. The ratio of restframe 5.4~keV to 1915\\AA~ fluxes ($\\nu f_\\nu$) is 0.45 for CXO~52, and the ratio of restframe 25.7~keV to 1915\\AA~ fluxes is 4.6 (Stern \\etal). The corresponding values for CDFS-263 are an order of magnitude lower, with a ratio of restframe 5.8~keV to 1957\\AA~ fluxes of 0.05 and of restframe 27.9~keV to 1957\\AA~ fluxes of 0.16 (Mainieri \\etal). The 2~Ms deep exposure by {\\it Chandra} has a lower flux limit for the detection of sources of $1.9-9.3\\times10^{-17}\\,{\\rm erg\\, cm^{-2}\\, s^{-1}}$ in the $0.5-2$~keV (soft) band and $1.1-7.5\\times10^{-16}\\,{\\rm erg\\, cm^{-2}\\, s^{-1}}$ in the $2-8$~keV (hard) band, depending on position (Alexander private communication as quoted in Stanway \\etal 2004b). At $z\\approx6$, this implies a source like CXO~52 would need to have an AB magnitude at 1.3~$\\mu$m (observed) of about $26.5-28$ to be detectable in the x-ray, while a source like CDFS-263 would need to be 2.5 or more magnitudes brighter. This is comparable to or somewhat brighter than the $J$- and $H$-band detection limits of the deep surveys used to detect the band-dropout objects.  Although genuine high redshift band-dropout objects are almost certainly dominated by starbursts, the possibility that a few are AGN can have major implications for the re-ionization of the IGM. The potentially important role played by AGNs in re-ionizing the universe may be demonstrated by comparing the numbers of UV ionizing photons generated by black holes and by stars through the lifetime of the universe. The specific luminosity of an AGN at the Lyman edge is $L_L\\simeq0.1L_{\\rm bol}/ \\nu_L$, where $L_{\\rm bol}$ is the bolometric luminosity of the AGN and $\\nu_L$ is the frequency of the Lyman edge. The production rate of ionizing photons is then $\\dot N_{\\rm bh}\\simeq0.2L_{\\rm bol}/(h\\nu_L)$ (Meiksin 2005a), where $h$ is the Planck constant. Assuming a conversion efficiency $\\epsilon_{\\rm bh}$ of mass into energy for accretion onto a black hole, the global comoving number density of ionizing photons produced by black hole accretion in the universe changes at the rate \\begin{equation} \\dot n_{\\rm bh} \\simeq \\frac{0.2\\epsilon_{\\rm bh}\\dot\\rho_{\\rm bh} c^2} {h\\nu_L}, \\label{eq:nbdot} \\end{equation} where $\\dot\\rho_{\\rm bh}$ is the rate at which the average comoving mass density of black holes in the universe increases with time (assumed largely due to accretion). The efficiency of mass conversion is unknown, but phenomenological estimates are in the range $0.1\\lsim\\epsilon_{\\rm bh}\\lsim0.3$ (eg, Yu \\& Tremaine 2002).  The production rate of \\HI ionizing photons by stars is sensitive to the initial mass function (IMF) of the stars and to their metallicities.  For a Salpeter IMF with metallicity 20\\% of solar, the results of Smith, Norris \\& Crowther (2002) correspond to a production rate of ionizing photons per solar mass of stars formed of $dN_*/dM\\approx10^{61}\\,{\\rm ph\\, M_\\odot^{-1}}$. For solar metallicity, the production rate is a factor of 3 smaller, while it may be somewhat larger for Pop~III stars. The comoving number density of ionizing photons produced by stars then changes at the rate \\begin{equation} \\dot n_* \\simeq\\frac{dN_*}{dM}\\dot\\rho_*, \\label{eq:nsdot} \\end{equation} where $\\dot\\rho_*$ is the rate at which the average comoving mass density of stars in the universe increases with time.  The number of ionizing photons available to re-ionize the IGM is reduced by internal absorption, both in an AGN and in a star-forming galaxy. Denoting the escape fractions from AGN and galaxies by $f_{\\rm bh, esc}$ and $f_{\\rm *, esc}$, respectively, the ratio of the ionization rate of the IGM by black holes to that by stars is \\begin{equation} \\frac{\\dot n_{\\rm bh}}{\\dot n_*}\\approx330 \\left(\\frac{\\epsilon_{\\rm bh}}{0.2}\\right) \\left(\\frac{dN_*/dM}{10^{61}\\,{\\rm ph\\,M_\\odot^{-1}}}\\right) \\left(\\frac{f_{\\rm bh, esc}}{f_{\\rm *, esc}}\\right) \\left(\\frac{\\rho_{\\rm bh}}{\\rho_*}\\right). \\label{eq:ndotratio} \\end{equation} Here, the approximation is made that $\\dot\\rho_{\\rm bh}/\\dot\\rho_*\\approx\\rho_{\\rm bh}/\\rho_*$, averaged over the age of the universe, where $\\rho_{\\rm bh}$ and $\\rho_*$ are the respective mass densities of QSO black holes and in stars in the universe today.  This seems a reasonable approximation to make since the inferred growth time of a central massive black hole in a galaxy is comparable to that of the stars in the bulge over a broad range of bulge masses, at least in the present-day universe (Heckman \\etal 2004). Estimates for the current mass densities in QSO black holes and stars are $\\rho_{\\rm bh}\\approx10^5{\\rm M_\\odot\\,Mpc^{-3}}$ (Yu \\& Tremaine 2002) and $\\rho_*\\approx3\\times10^8{\\rm M_\\odot\\,Mpc^{-3}}$ (Baldry \\& Glazebrook 2003), both for $h=0.7$. The escape fraction of ionizing photons from AGN is unknown, but their spectra suggest at least half escape, so $f_{\\rm bh, esc}\\approx0.5$ is assumed. The observational upper limit on the escape fraction of ionizing photons from galaxies is $f_{\\rm *, esc}<0.04$ (Fernandez-Soto, Lanzetta \\& Chen 2002). This then gives for the ratio of IGM ionizing photons produced by black holes to that produced by stars \\begin{equation} \\frac{n_{\\rm bh}}{n_*}\\gsim \\left(\\frac{\\epsilon_{\\rm bh}}{0.2}\\right) \\left(\\frac{dN_*/dM}{10^{61}\\,{\\rm ph\\,M_\\odot^{-1}}}\\right) \\left(\\frac{f_{\\rm bh, esc}}{0.5}\\right) \\left(\\frac{f_{\\rm *, esc}}{0.04}\\right)^{-1}. \\label{eq:nratio} \\end{equation} Despite the low number of black holes compared with stars, their higher mass-to-energy conversion efficiency compared with stellar nuclear fusion and the larger escape fraction from AGN compared with galaxies make black holes competitive with stars as candidate sources for the photons which re-ionized the IGM.  The relative ionization rates may be related to the relative fractions of objects detected in deep surveys as follows. Defining $f_{\\rm AGN}$ and $f_{\\rm SB}$ to be the fraction of band-dropout objects that are AGNs and starbursts, respectively, $f_L/f_M$ the intrinsic ratio of flux densities at the Lyman edge and a fiducial frequency $\\nu_M$ normalising the counts of objects, and $\\alpha^{\\rm eff}_{\\rm AGN}$ and $\\alpha^{\\rm eff}_{\\rm SB}$ the respective effective spectral indices of AGN and starbursts shortward of the Lyman edge, then the ratio of ionizing photon rates of AGN to starbursts injected into the IGM is \\begin{equation} \\frac{\\dot n_{\\rm b}}{\\dot n_*}=\\frac{f_{\\rm AGN}}{f_{\\rm SB}} \\frac{f_{\\rm AGN, esc}}{f_{\\rm SB, esc}}\\left(\\frac{f_L}{f_M}\\right)_{\\rm AGN} \\left(\\frac{f_L}{f_M}\\right)_{\\rm SB}^{-1}\\frac{\\alpha^{\\rm eff}_{\\rm SB}} {\\alpha^{\\rm eff}_{\\rm AGN}}. \\label{eq:ndotratio_obs} \\end{equation} Here, $f_{\\rm AGN, esc}$ and $f_{\\rm SB, esc}$ are the respective escape fractions of ionizing photons from the observed AGN and starbursts. The mean observed AGN escape fraction in particular is expected to be much smaller than the mean black hole escape fraction above because the contribution from Type II AGN is expected to be very small. The only ionizing photons that may be observed from Type II AGN are those re-emitted from the emission line gas illuminated by the central engine:\\ the direct ionizing photons will be obscured. If only Type I AGN are counted in Eq.~\\ref{eq:ndotratio_obs}, then $f_{\\rm AGN, esc}$ may be taken to refer only to these objects, so that $f_{\\rm AGN, esc}\\approx f_{\\rm bh, esc}$, although it should be noted that the amount of re-radiated ionizing photons from Type II AGN is unknown and may not be completely negligible. It will also be assumed that $f_{\\rm SB, esc}\\approx f_{\\rm *, esc}$. For Pop~III stars and hard spectra AGN, $\\alpha^{\\rm eff}\\approx0.5$ is expected, while for Pop~II stars, $\\alpha^{\\rm eff}_{\\rm SB}\\approx1.8-2.3$ is expected (Meiksin 2005a), while soft spectra AGN may have $\\alpha^{\\rm eff}_{\\rm AGN}\\approx1.8$. Further uncertainty arises from the ratio of the flux $f_L$ at the Lyman limit to the flux $f_M$ measured most near the Lyman limit. But these likely combine to an uncertainty of a factor of only a few. The dominant uncertainty is $f_{\\rm bh, esc}f_{\\rm AGN}/ f_{\\rm *, esc}f_{\\rm SB}$. For escape fractions of $f_{\\rm bh, esc}\\approx0.5$ and $f_{\\rm *, esc}\\approx0.04$, only about one object in a dozen, with an uncertainty of a factor of at least a few, need by an AGN for AGN to compete with starbursts as the dominant source of ionizing photons. A further implication of a substantial AGN contribution to the re-ionization of the IGM is that the post-reionization temperature will be substantially boosted if the AGN are sufficiently hard to ionize \\HeII to \\HeIII as well, which current AGN number counts suggest they may be able to do by $z\\lsim5.5$ (Meiksin 2005a).  Ultimately, combining {\\it HST} observations of $z_{850}$-dropouts with future follow-up {\\it JWST} imaging to detect high redshift Type II AGN may be the best means of settling the question of how prevalent Type II QSOs are at high redshifts."
    },
    "0512/astro-ph0512571_arXiv.txt": {
        "abstract": "{ Previously, large discrepancies have been found between theory and observation for \\ion{Fe}{xv} emission line ratios in solar flare spectra covering the 224--327~\\AA\\ wavelength range, obtained by the Naval Research Laboratory's S082A instrument on board {\\em Skylab}. These discrepancies have been attributed to either errors in the adopted atomic data or the presence of additional atomic processes not included in the modelling, such as fluorescence. However our analysis of these plus other S082A flare observations (the latter containing \\ion{Fe}{xv} transitions between 321--482~\\AA), performed using the most recent \\ion{Fe}{xv} atomic physics calculations in conjunction with a {\\sc chianti} synthetic flare spectrum, indicate that blending of the lines is primarily responsible for the discrepancies. As a result, most \\ion{Fe}{xv} lines cannot be employed as electron density diagnostics for solar flares, at least at the spectral resolution of S082A and similar instruments (i.e. $\\sim$0.1~\\AA). An exception is the intensity ratio I(3s3p $^{3}$P$_{2}$--3p$^{2}$ $^{3}$P$_{1}$)/I(3s3p $^{3}$P$_{2}$--3p$^{2}$ $^{1}$D$_{2}$) = I(321.8~\\AA)/I(327.0~\\AA), which appears to provide good estimates of the electron density at this spectral resolution. ",
        "introduction": "Emission lines of \\ion{Fe}{xv} are prominent features of the solar extreme ultraviolet (EUV) spectrum, with numerous transitions present in the $\\sim$220--500~\\AA\\ wavelength interval (see, for example, Dere 1978; Thomas \\& Neupert 1994). It has long been known that these lines provide useful electron density diagnostics for the emitting plasma (Bely \\& Blaha 1968), but to date there have been relatively few detailed analyses of the \\ion{Fe}{xv} solar EUV spectrum. Probably the most comprehensive have been those of Young et al. (1998) and Keenan et al. (2005),  both of which employed high resolution solar active region spectra obtained by the Solar EUV Research Telescope and Spectrograph (SERTS), and generally found good agreement between theory and observation.  However the situation for solar flares is very different. The most detailed study of \\ion{Fe}{xv} flare lines was undertaken by Dufton et al. (1990), using spectra from the S082A instrument on board {\\em Skylab}. These authors found major discrepancies between theory and observation, which they attributed to possible errors in the adopted atomic data, especially for higher-lying levels, which would be more populated in flares than in active regions due to the higher electron densities. Alternatively, Dufton et al. suggested that some additional atomic process could be in operation, which may be more important in flares than in active regions, such as fluorescence.  In this paper we use the most recent atomic physics calculations for \\ion{Fe}{xv}, in conjunction with a synthetic flare spectrum generated with the {\\sc chianti} database (Dere et al. 1997; Young et al. 2003), to investigate if the discrepancies found between theory and observation by Dufton et al. (1990) may be resolved. In addition, we perform an analysis of other S082A flare spectra which contain \\ion{Fe}{xv} emission lines not considered by these authors. ",
        "conclusions": "As noted in Section 1, Dufton et al. (1990) found large discrepancies between theory and observation for \\ion{Fe}{xv} line ratios measured from solar flare spectra obtained with the S082A instrument on board {\\em Skylab}. Specifically, they noted that: \\begin{itemize} \\item the intensities of the 3s3p $^{3}$P$_{J}$--3s3d $^{3}$D$_{J^\\prime}$ transitions between $\\sim$224--235~\\AA, and that of the 3s3p $^{1}$P$_{1}$--3s3d $^{1}$D$_{2}$ line at 243.8~\\AA, when ratioed against the 3s3p $^{3}$P$_{2}$--3s3d $^{3}$D$_{3}$ feature at 233.9~\\AA, give values of line ratios which imply electron densities far lower (by up to an order of magnitude) than those indicated by other diagnostics, \\item the observed ratios for the intensity of the 243.8~\\AA\\ line against those of the 3s3p $^{3}$P$_{J}$--3s3d $^{3}$D$_{J^\\prime}$ transitions all lie well outside the range of values allowed by theory. \\end{itemize}  Dufton et al. (1990) suggested that these discrepancies were due to either (i) errors in the adopted atomic data, arising for example from the neglect of fine-structure in the excitation cross section calculations (which were generated in LS-coupling), (ii) the limited number of fine-structure levels (14) included in the diagnostic calculations, as for high electron density flares excitation to higher-lying levels followed by cascades may be important, (iii) the possibility of additional atomic processes being responsible for the \\ion{Fe}{xv} line emission, such as fluorescence. However, more recently Kastner \\& Bhatia (2001) noted that blending in the observational data may be responsible, although they did not investigate this possibility.  The present theoretical line ratios for \\ion{Fe}{xv} discussed in Section 3 have been calculated using fully relativistic electron excitation rates, as opposed to those generated in LS-coupling employed by Dufton et al. (1990). Additionally, the ratios have been produced with a more extensive model ion than that considered by Dufton et al., with 53 levels as opposed to 14. Nevertheless, as noted in Section 3, the present theoretical ratios are quite similar to those of Dufton et al. Crucially, they do not resolve the discrepancies between theory and observation found for the {\\em Skylab} flare spectra, indicating that problems with the theoretical ratios are unlikely to be the source of the problem.  We therefore consider blending in the observations. Firstly, Dere (1978) point out that the 233.9~\\AA\\ line of \\ion{Fe}{xv} is partially blended with a \\ion{Ni}{xviii} transition. However, Dufton et al. (1990) note that the \\ion{Ni}{xviii} blend can be allowed for, while Keenan et al. (1997) showed that the \\ion{Fe}{xv} and \\ion{Ni}{xviii} features can be well resolved and reliable line intensities determined. To investigate this, we have generated a synthetic solar flare spectrum  using the Spectral Synthesis Package option (ch$\\underline{~~}$ss) within Version 4.2 of the {\\sc chianti} database. Details of {\\sc chianti} and instructions on its use may be found in Dere et al. (1997) and Young et al. (2003), and also on the website http://wwwsolar.nrl.navy.mil/chianti.html. In our {\\sc chianti} calculations, we adopted a constant electron density of 10$^{11}$ cm$^{-3}$, the ionization balance of Mazzotta et al. (1998), and the solar flare differential emission measure (DEM) option. However we note that varying these parameters, even by relatively large amounts, does not significantly affect the following discussions nor the conclusions drawn as to the identities of blending species. An example of this is provided below in our analysis of the \\ion{Fe}{xv} 227.2~\\AA\\ transition. In addition, we stress that we are not attempting to accurately assess the amount of blending in the \\ion{Fe}{xv} lines. Given the uncertainties in both the observed and theoretical \\ion{Fe}{xv} line ratios, plus possible errors in the calculated {\\sc chianti} flare spectrum, this would be very difficult. Instead, we are simply investigating if blending species can be identified, and if the predicted (approximate) level of blending is consistent with what is observed, hence providing an explanation for the discrepancies found between theory and experiment for the \\ion{Fe}{xv} flare spectra.  The synthetic {\\sc chianti} flare spectrum indicates that no other feature blends with the 233.9~\\AA\\ transition at greater than the 1\\%\\ level, so that the line intensity should be well determined, although of course unidentified features could be present. Unfortunately, the same argument does not apply to the other \\ion{Fe}{xv} transitions considered by Dufton et al. (1990). The 243.8~\\AA\\ line is badly blended with an \\ion{Ar}{xiv} feature (Dere 1978), which has been confirmed by Keenan et al. (2003) in their analysis of \\ion{Ar}{xiv} line ratios from the S082A flare spectra. Indeed, Keenan et al. have shown that \\ion{Ar}{xiv} contributes typically around 40\\%\\ to the \\ion{Fe}{xv}/\\ion{Ar}{xiv} 243.8~\\AA\\ blend. Similarly, the 3s3p $^{3}$P$_{0}$--3s3d $^{3}$D$_{1}$ transition of \\ion{Fe}{xv} at 224.8~\\AA\\ is actually listed as a \\ion{S}{ix} feature by Dere, and the {\\sc chianti} synthetic flare spectrum indicates that \\ion{S}{ix} is responsible for approximately 75\\%\\ of the measured line intensity.  In the case of the 3s3p $^{3}$P$_{1}$--3s3d $^{3}$D$_{2}$ line of \\ion{Fe}{xv} at 227.2~\\AA, {\\sc chianti} lists a \\ion{C}{v} transition which is predicted to contribute about 10\\%\\ to the total measured intensity. Although relatively small, this degree of blending is generally sufficient to explain the discrepancies between theory and observation for the R = I(227.2~\\AA)/I(233.9~\\AA) ratio, due primarily to the fact that R is not very sensitive to variations in the electron density and hence small changes in the measured ratio lead to large ones in the derived value of N$_{\\mathrm{e}}$. For example, for the 1973 August 9 flare, Dufton et al. measured R = 0.30 which implied N$_{\\mathrm{e}}$ = 10$^{9.6}$~cm$^{-3}$, much lower than that derived from line ratios in \\ion{Fe}{xiv} (N$_{\\mathrm{e}}$ = 10$^{10.9}$~cm$^{-3}$; Keenan et al. 1991a), which is formed at a similar electron temperature to \\ion{Fe}{xv}. However if one assumes that the {\\sc chianti} prediction for the strength of the \\ion{C}{v} blending line is correct, and hence reduce the measured R ratio by 10\\%\\ so that R$_{\\mathrm{deblended}}$ = 0.27, the present \\ion{Fe}{xv} line ratio calculations then indicate N$_{\\mathrm{e}}$ = 10$^{10.3}$~cm$^{-3}$, much closer to the \\ion{Fe}{xiv} value. Indeed, the \\ion{C}{v} contribution to the 227.2~\\AA\\ feature would only need to be increased to 17\\%\\ for the \\ion{Fe}{xv} electron density to be the same as that deduced from \\ion{Fe}{xiv}. We note that there are numerous uncertainties in the {\\sc chianti} synthetic spectrum, including for example errors in the adopted electron impact excitation rates or emission measure distribution, and possible variations in element abundances. Hence an increase in the \\ion{C}{v} contribution from 10\\%\\ to 17\\%\\ would not be unreasonable. For example, if we reduce the adopted electron density in the {\\sc chianti} synthetic spectrum to 10$^{10}$ cm$^{-3}$, and change the DEM from that for a flare to the quiet Sun, the predicted \\ion{C}{v} contribution to the 227.2~\\AA\\ blend increases from 10\\%\\ to 50\\%. This also illustrates our point above, namely that varying the adopted {\\sc chianti} parameters by even large amounts does not affect our conclusions regarding the identifications of line blends.  Similarly, for the 3s3p $^{3}$P$_{2}$--3s3d $^{3}$D$_{2}$ line of \\ion{Fe}{xv} at 234.8~\\AA, the {\\sc chianti} synthetic spectrum indicates blending with two \\ion{Ne}{iv} lines, with the \\ion{Fe}{xv} component only contributing about 75\\%\\ to the measured 234.8~\\AA\\ intensity. Reducing the experimental I(234.8~\\AA)/I(233.9~\\AA) intensity ratios to allow for this blending once again improves agreement between theory and observation. For example, Dufton et al. (1990) measured the I(234.8~\\AA)/I(233.9~\\AA) ratio for the 1974 January 21 flare to be 0.10, implying N$_{\\mathrm{e}}$ = 10$^{9.5}$~cm$^{-3}$. However removing the predicted \\ion{Ne}{iv} contribution to the 234.8~\\AA\\ intensity reduces the experimental ratio to 0.08, which then yields N$_{\\mathrm{e}}$ = 10$^{10.5}$~cm$^{-3}$. This is in very good agreement with the value of N$_{\\mathrm{e}}$ = 10$^{10.7}$~cm$^{-3}$ determined from \\ion{Fe}{xiv} (Keenan et al. 1991a).  In the case of the 3s3p $^{3}$P$_{1}$--3s3d $^{3}$D$_{1}$ feature of \\ion{Fe}{xv} at 227.7~\\AA, the {\\sc chianti} synthetic spectrum does not predict any significant blending species, with the strongest line (\\ion{Fe}{xii} 227.66~\\AA) only contributing about 3\\%\\ to the total 227.7~\\AA\\ intensity. However the discrepancy between theory and observation indicates the presence of a major blend. For example, in the 1973 December 17 flare, Dufton et al. (1990) measure R = I(227.7~\\AA)/I(233.9~\\AA) = 0.18, which gives N$_{\\mathrm{e}}$ = 10$^{9.3}$~cm$^{-3}$, compared to the \\ion{Fe}{xiv} density estimate of N$_{\\mathrm{e}}$ = 10$^{10.4}$~cm$^{-3}$ (Keenan et al. 1991a). The experimental R ratio needs to be reduced to 0.11 in order for the \\ion{Fe}{xv} density to match that deduced from \\ion{Fe}{xiv}, implying that the blending species in the 227.7~\\AA\\ feature contributes about 40\\%\\ to the total intensity. An inspection of line lists, such as the NIST Database (http:$/$$/$physics.nist.gov$/$PhysRefData$/$), reveals the only likely candidates for the blend to be the 2p$^{2}$ $^{3}$P$_{1}$--2p3s $^{3}$P$_{0}$ (227.63~\\AA) and 2p$^{2}$ $^{3}$P$_{2}$--2p3s $^{3}$P$_{1}$ (227.69~\\AA) transitions of \\ion{O}{v}. Other $\\Delta$n = 1 lines of \\ion{O}{v} have been detected in the S082A flare spectra between $\\sim$193--249~\\AA, arising from 2s2p--2s3s and 2s2p--2s3d transitions (Keenan et al. 1991b). Unfortunately, no atomic data exist for the 2p$^{2}$ $^{3}$P--2p3s $^{3}$P lines, and hence theoretical intensity ratios cannot be generated involving these features, to compare with the observations. However, Huang et al. (1988) have detected the 2p$^{2}$ $^{3}$P--2p3s $^{3}$P lines in a tokamak spectrum, and found their intensities to be comparable to those of the 2s2p--2s3s and 2s2p--2s3d features. We are therefore confident that the blending in the 227.7~\\AA\\ line arises from \\ion{O}{v}.  Based on the above discussions, we conclude that the \\ion{Fe}{xv} lines considered by Dufton et al. (1990) are blended in the S082A flare spectra, hence providing an explanation for the discrepancies between theory and observation found by these authors. As a result, it is clear that the \\ion{Fe}{xv} features in the short wavelength S082A plates cannot be employed as electron density diagnostics for solar flares. However, does the same argument apply to the \\ion{Fe}{xv} lines detected on the long wavelength S082A plates? An inspection of Table 1 reveals that all of the observed line ratios are much larger than allowed by theory. For example, in the 1973 August 9 flare, the measured intensity ratio I(321.8~\\AA)/I(417.3~\\AA) = 1.0, while from Fig. 3 the theoretical high density limit is 0.14. Similarly the experimental I(481.5~\\AA)/I(417.3~\\AA) ratio for this flare is 0.26, compared to the theoretical high density limit of 0.13. The {\\sc chianti} synthetic flare spectrum indicates that the 417.3~\\AA\\ transition should be free from blends, with no nearby features predicted to contribute more than about 0.5\\%\\ to the total measured line intensity. However, unfortunately the 481.5~\\AA\\ feature is very badly blended with \\ion{Ne}{v} transitions, which according to {\\sc chianti} are responsible for around 80\\%\\ of the total line flux.  In the case of the 321.8 and 327.0~\\AA\\ lines of \\ion{Fe}{xv}, {\\sc chianti} indicates either the presence of no blending species (327.0~\\AA), or only a small ($\\sim$20\\%) contribution from \\ion{Fe}{x} transitions (321.8~\\AA). A search of line lists also reveals no likely blends. Furthermore, intensity ratios involving these lines in active region spectra show good agreement between theory and observation (Keenan et al. 2005), and one would actually expect \\ion{Fe}{x} to make a smaller contribution to the 321.8~\\AA\\ line intensity in a flare than in an active region, given that \\ion{Fe}{xv} is more highly ionized and hence should be relatively much stronger in the former. (Indeed, {\\sc chianti} predicts that \\ion{Fe}{x} contributes around 50\\%\\ of the 321.8~\\AA\\ line intensity in an active region). Given these facts, it is difficult to believe that blending is responsible for the discrepancies between theory and observation for the I(321.8~\\AA)/I(417.3~\\AA) and I(327.0~\\AA)/I(417.3~\\AA) intensity ratios. Significant problems with the measurements of the 417.3~\\AA\\ line can also be ruled out, as the experimental I(321.8~\\AA)/I(327.0~\\AA) ratios similarly show discrepancies with theory, being up to $\\sim$75\\%\\ larger than the theoretical high density limit of 0.55. We therefore believe that the discrepancies are due to intensity calibration uncertainties in the S082A observations, most probably arising from the fact that the 321.8 and 327.0~\\AA\\ lines lie close to the short wavelength cutoff of the plates, and hence are more susceptible to abberation effects (see Section 2). Support for this comes from a spectrum of a solar flare in the 280--330~\\AA\\ region (Zhitnik et al. 2005), obtained by the RES--C spectroheliograph on the CORONAS--F orbital station, at a resolution of 0.1~\\AA. Although most of the \\ion{Fe}{xv} lines in the RES--C flare spectrum are blended (as in the S082A data), a reliable measurement is available for the \\ion{Fe}{xv} intensity ratio I(321.8~\\AA)/I(327.0~\\AA) = 0.48$\\pm$0.05. This implies N$_{e}$ = 10$^{10.5\\pm0.4}$~cm$^{-3}$ (Keenan et al. 2005), in excellent agreement with typical flare densities derived from, for example, \\ion{Fe}{xiv} line ratios (Keenan et al. 1991a).  In Table 2 we summarise our identifications of lines which blend with the \\ion{Fe}{xv} transitions in the S082A flare spectra. On the basis of this, we may confidently state that most of the \\ion{Fe}{xv} lines in solar flare spectra are blended (at least at spectral resolutions of $\\sim$0.1~\\AA). There is hence no need to invoke errors in atomic data or other reasons to explain discrepancies between theory and observation, as suggested by Dufton et al. (1990). As a consequence, \\ion{Fe}{xv} flare lines do not provide useful electron density diagnostics for the emitting plasma, with the exception of the I(321.8~\\AA)/I(327.0~\\AA) intensity ratio.  \\begin{table} \\caption{Summary of blending species for \\ion{Fe}{xv} transitions in S082A flare spectra} \\label{table:1} \\begin{tabular}{l l} \\hline \\ion{Fe}{xv} line (\\AA) & Possible blending species \\\\ \\hline 224.8 & Mostly due to \\ion{S}{ix} 224.8~\\AA \\\\ 227.2 & \\ion{C}{v} 227.2~\\AA\\ contributes 10--20\\%\\ to the blend \\\\ 227.7 & Major blend with \\ion{O}{v} 227.6 and 227.7~\\AA \\\\ 233.9 \\\\ 234.8 & \\ion{Ne}{iv} 234.7~\\AA\\ contributes about 20\\%\\ to the blend \\\\ 243.8 & Badly blended with \\ion{Ar}{xiv} 243.8~\\AA \\\\ 321.8 & \\ion{Fe}{x} 321.8~\\AA\\ contributes about 20\\%\\ to the blend \\\\ 327.0 \\\\ 417.3 \\\\ 481.5 & Badly blended with \\ion{Ne}{v} 481.4~\\AA \\\\ \\hline \\end{tabular}  \\end{table}"
    },
    "0512/hep-th0512236_arXiv.txt": {
        "abstract": "The Bethe-Salpeter equation in a strong magnetic field is studied for positronium atom in an ultra-relativistic regime, and a (hypercritical) value for the magnetic field is determined, which provides the full compensation of the positronium rest mass by the binding energy in the maximum symmetry state. The compensation becomes possible owing to the falling to the center phenomenon. The relativistic form in two-dimensional Minkowsky space is derived for the four-dimensional Bethe-Salpeter equation in the limit of an infinitely strong magnetic field, and used for finding the above hypercritical value. Once the positronium rest mass is compensated by the mass defect the energy barrier separating the electron-positron system from the vacuum disappears. We thus describe the structure of the vacuum in terms of strongly localized states of tightly mutually bound (or confined) pairs. Their delocalization for still higher magnetic field, capable of screening its further growth, is discussed. ",
        "introduction": "In the present paper we are considering, in the framework of quantum electrodynamics, the phenomenon of falling to the center in a system of two charged particles caused by the ultraviolet singularity $1/x^2$ of the photon propagator - which mediates the interaction between them - on the light-cone, $x^2=x^2_0-{\\bf x}^2\\simeq 0$. Here $x_0$ is the time, and $\\bf x$ is the space coordinate. This phenomenon occurs in a number of problems. In some of cases listed below the matter comes to the one-dimensional \\Sch ~with the singular attractive potential $U(r)=-\\beta/r^2$, $0<r<\\infty$, the falling to the center depending on the value of the constant $\\beta$ that fixes the strength of the coupling \\cite{QM}.  Long ago it was established  that the falling to the center is inherent in \\BS ~for electron-positron system treated in a fully relativistic way and taken in the ladder approximation, provided that the fine structure constant $\\alpha$ exceeds a critical value $\\alpha_{\\rm cr}$, $(1/137)\\ll\\alpha_{\\rm cr}<1$. We refer to the so-called Goldstein solution \\cite{goldstein} of the Bessel differential equation in the variable $s=\\sqrt{-x^2},~$ the space-like interval between the particles, to which the problem is reduced in the ultra-relativistic limit. F\\al~ implies that the singular attraction is so strong that the energy spectrum becomes unlimited from below, and hence - once the singular problem is duely defined - the binding energy may completely cancel the rest mass, so that the gap between positive and negative continua disappears. This reminds to some extent the situation in the problem of an electron in an external electric field, occupying the whole space, where the tunnelling through the gap leads to the vacuum instability with respect to $e^+e^--$pair creation known as Schwinger effect, although the latter does not have a threshold character. The value $\\alpha=\\alpha_{\\rm cr}$ may delimit the range of self-consistency of the theory known as quantum electrodynamics, at least in its customary form.  Another well-known theater \\cite{popov}, \\cite{greiner}, \\cite{shabad2} where the f\\al~ acts is the point-like nucleus with its charge $Z$ exceeding the critical value $Z_{\\rm cr}=\\alpha^{-1}=137$, into whose Coulomb field a relativistic electron described by the Dirac equation is placed. After this problem is properly regularized, the electron energy level reaches the border $\\varepsilon=-m$ of the lower continuum as the charge $Z$ grows, which gives rise to electron-positron pair production, the free positrons leaving the atom. This mechanism leads to diminishing the charge and restores it at the critical value that cannot thus be exceeded.  If an electron in an atom is treated within the \\Sch, or a pair of particles is addressed in a semi-relativistic way with the aid of \\BS, wherein the famous Salpeter's equal-time Ansatz is made (resulting in the disregard of the retardation effects), the f\\al ~does not take place. The matter is that the singularity in the point $r=0$, $r=\\sqrt{{\\bf x}^2}$ of the Coulomb potential $U(r)=-\\alpha Z/r$ originating in this case from the light-cone singularity of the photon propagator, mentioned above, is not sufficiently strong, and the energy level remains shallow. On the contrary, in the above two cases the f\\al~ may be attributed to relativistic enhancement of the Coulomb force - the second-order differential equation to which the spherically symmetric Dirac equation may be reduced by excluding the second spinor component, contains the sufficiently singular attractive term  $-(\\alpha Z/r)^2$ in the effective potential \\cite{popov}.  The situation changes drastically when a strong external magnetic field  ${\\bf B}$ is applied to the system. Already the \\Sch ~for a particle situated in the combination of the Coulomb and strong magnetic fields possesses the f\\al ~caused by the singularity $1/|z|$ in the differential \\Sch ~defined on the whole axis $-\\infty<z<\\infty$. Here $z$ is the electron coordinate along the axis parallel to the magnetic field. The energy spectrum of this equation contains large negative values tending to minus infinity \\cite{loudon} as the magnetic field $B\\rightarrow\\infty$. The reason lies in the dimensional reduction: for strong magnetic field the electron is restricted to the lowest Landau level, consequently its wave function obeys a two-dimensional (one space and one time coordinates) differential equation. The reduction of the number of degrees of freedom causes the effective strengthening of the attraction, like in the quark confinement problem. Analogous strengthening takes place when the electron-positron system is described by  \\BS ~studied using the semi-relativistic Ansatz that implies the non-relativistic character of the relative motion of the constituent particles \\cite{leinson1} -% ~\\cite{leinson2}. Again, the energy spectrum of the corresponding two-dimensional equation is unlimited from below as the magnetic field grows. The total rest mass of the ground state is compensated  when the magnetic field reaches the value \\cite{ShUs}, whose order of magnitude is determined by the large factor $\\exp (\\rm{const}/\\alpha)$. % As pointed in \\cite{ShUs} % those results, however, are not reliable, since they depend on the unrighteous extrapolation of non-relativistic procedure to the essentially relativistic region of large mass defect.  The relativistic treatment of an electron in a combination of external Coulomb and magnetic field was given in \\cite{semikoz}. The result was that if the magnetic field is reasonably strong ($\\sim 10^{15}$ G) the critical value of the nucleus charge $Z$ may correspond to stable elements with $Z\\sim 90$.  In the present paper we consider the system of two charged relativistic particles - especially the electron and positron - in interaction with each other, when placed in a strong constant and homogeneous magnetic field $B$. To this end we use \\BS ~ in the ladder approximation without exploiting any non-relativistic assumption. We derive its limiting - when $B\\rightarrow\\infty$ - form, which comes out to be a Bethe-Salpeter equation in two-dimensional Minkowsky space-time, covariant under the corresponding Lorentz subgroup - the boost along the magnetic field. Stress, that the two-dimensionality holds only with respect to the degrees of freedom of charged particles, while the photons remain 4-dimensional in the sense that  the singularity of the photon propagator is determined by the inverse d'Alambert operator in the 4-dimensional, and not two-dimensional Minkowsky space. (Otherwise it would be weaker). The term responsible for interaction with an arbitrary electric field directed along ${\\bf B}$ is also included and does not lay obstacles to the dimensional reduction. Throughout the paper we set $\\hbar=c=1$ and use the Heaviside-Lorentz units, where $\\alpha=e^2/4\\pi.$ $\\left.B\\right|_{\\rm Gaussian}=\\sqrt{4\\pi} \\left.B\\right|_{\\rm Heaviside-Lorentz}.$  We make sure that in the case under consideration the critical value of the coupling constant is zero, $\\alpha_{\\rm cr}=0$, i.e., the f\\al ~is present already for its genuine value $\\alpha=1/137$, in contrast to the no-magnetic-field case, where $\\alpha_{\\rm cr}>1/137$. If the magnetic field is large, but finite, the dimensional reduction holds everywhere except a small neighborhood of the singular point $s=0$, wherein the mutual interaction between the particles dominates over their interaction with the magnetic field. The dimensionality of the space-time in this neighborhood remains to be 4, and its  size is determined by the Larmour radius $L_B=(eB)^{-1/2}$ that is zero in the limit $B=\\infty$. The latter supplies the singular problem with a regularizing length. The larger the magnetic field, the smaller the regularizing length, the deeper the level. We find the value of the magnetic field - we call it \\textit{first hypercritical field} - \\bee\\label{final0} B^{(1)}_{\\rm hpcr}=\\frac{m^2}{4e}\\exp\\left\\{\\frac{\\pi^{3/2}} {\\sqrt{\\alpha}}+2C_{\\rm E}\\right\\},\\quad \\eend where $C_{\\rm E}=0.577$ is the Euler constant, that provides disappearance of the center-of-mass energy of the electron-positron pair and of its center-of-mass momentum component along $\\bf B$. We refer to this situation as a collapse of positronium.  In discussing the physical consequences of the f\\al ~we appeal to the approach recently developed by one of the present authors as applied to the \\Sch ~with singular potential \\cite{shabad} and to the Dirac equation in supercritical Coulomb field \\cite{shabad2}. Within this approach the singular center looks like a black hole. The solutions of the differential equation that oscillate near the singularity point are treated as free particles emitted and absorbed by the singularity. This treatment becomes natural after the differential equation is written as the generalized eigenvalue problem with respect to the coupling constant. Its solutions make a Hilbert space and are subject to orthonormality relations with a singular measure. This singularity makes it possible for the oscillating solutions to be normalized  to  $\\delta$-functions, as free particle wave-functions should be. The nontrivial, singular  measure that appears in the definition of the scalar product of quantum states in the Hilbert space of quantum mechanics introduces the geometry of a black hole of non-gravitational origin and the idea of horizon. The deviation from the standard quantum theory manifests itself in this approach only when particles are so close to one another that the mutual Coulomb field they are subjected to falls beyond the range, where the standard theory may be referred to as firmly established \\cite{shabad2}.  Within this approach the regularizing  length provided by the Larmour radius is dealt with not as a cut-off, but as a lower border of the normalization volume, the event horizon in a way. Although the result (\\ref{final0}) is obtained following the concept of Refs. \\cite{shabad}, \\cite{shabad2}, it can be reproduced without essential alteration within the standard cut-off philosophy, too.  The most intriguing question is what happens after the magnetic field exceeds the first hypercritical value (\\ref{final0}). The solution of \\BS ~in two-dimensional space-time in the ultra-relativistic limit studied in the present paper corresponds to formation  of  special \"confined\" states in the kinematical domain called sector III in \\cite{shabad}, \\cite{shabad2}. (Within the standard approach these would be bound states, although this is less adequate). As the corresponding overall energy and momentum of such $e^+e^-$-state is zero, it is not separated from the vacuum by an energy barrier. Besides, this state is the one of maximum symmetry in the coordinate and spin space. Hence, it may be thought of as relating to the vacuum, as well, and describing its structure. The confined particles cannot escape to infinite distance from one another, on the contrary the probability density of the confined state is concentrated near the point $s=0$, behind the horizon - as distinct from the ordinary bound state.  The situation is expected to change as the magnetic field goes on growing. At a certain stage - we reserve the name \\textit{second hypercritical} for the corresponding value of the magnetic field - deconfinement of the above strongly localized states may occur. The corresponding solutions to the Bethe-Salpeter equation are not yet strictly obtained, which makes us describe the deconfinement more hypothetically. After the level deepens further, the center-of-mass 2-momentum  gets into sufficiently far space-like region, and solutions oscillating at large distances between electron and positron appear. Thus, the state delocalizes. The delocalized electron-positron pairs produced from the vacuum, each particle  on a Larmour orbit, should screen the magnetic field and stop its growing above the second hypercritical value, this value being the absolute maximum of the magnetic field admitted within quantum electrodynamics. Simultaneously, the space-like total momentum provides the lattice structure to the vacuum.  This resembles the case of supercritical nucleus where there are states (that belong to sector IV in the nomenclature of Refs.\\cite{shabad}, \\cite{shabad2}) admitting the leakage to infinity, which provides the mechanism for reducing the  charge of the nucleus below the critical value.  No sooner than the delocalized states are found in our present problem one may definitely claim the instability of the vacuum with the second hypercritical magnetic field or - which is the same - the instability of such field under the pair creation that might provide the mechanism for its diminishing. For the present, we state that the first hypervalue (\\ref{final0}) is such a value of the magnetic field, the exceeding of which would already cause restructuring of the vacuum and  demand a profound revision of quantum electrodynamics.  The paper is organized as follows. In Section 2 we revisit Goldstein's solution by referring to various possibilities in approaching the f\\al, especially the one invoked by the previous work \\cite{shabad}, \\cite{shabad2}.  In Section 3 we derive the ultimate two-dimensional form of \\BS ~in its differential version characteristic of the ladder approximation, when the magnetic field tends to infinity, with the help of expansion over the complete set of Ritus matrix eigenfunctions \\cite{ritus}. The latter accumulate the spacial and spinor dependence on the transversal-to-the-field degree of freedom. The Fourier-Ritus transform of the Bethe-Salpeter amplitude obeys an infinite chain of coupled differential equations that decouple in the limit of large $B$, so that we are left with one closed equation for the amplitude component with the Landau quantum numbers of the electron and positron both equal to zero, while the components with other values of Landau quantum numbers vanish in this limit. The resulting equation is a differential equation with respect to two variables that are the differences of the particle coordinates: along the time $t=x_0^\\rme-x_0^\\rmp$ and along the magnetic field $z=x_3^\\rme-x_3^\\rmp$. It contains only two Dirac matrices $\\gamma_0$ and $\\gamma_3$ and can be alternatively written using $2\\times 2$ Pauli matrices. Arbitrary external electric field $\\bf E$ along $\\bf B$ is also included, $E\\ll B$. By introducing different masses the resulting two-dimensional equation is easily modified to cover also the case of an one-electron atom in strong magnetic field and/or other pairs of charged particles.  In Section 4 the ultra-relativistic  solutions (possessing maximum symmetry) to the equation derived in Section 3 are depicted corresponding to the vanishing energy-momentum of the $e^+e^-$-state, and the first hypercritical magnetic field is found basing on the standing wave boundary condition imposed on the lower border of the normalizing volume - as prescribed by the theory in Refs. \\cite{shabad}, \\cite{shabad2}. Also the standard cut-off procedure of Refs. \\cite{QM}, \\cite{popov} is fulfilled to give practically the same value (\\ref{final0}). Further, we estimate possible modifications that might be introduced by radiative corrections to the mass operator, to find that these cannot change the conclusions any essentially, and discuss the deconfinement. ",
        "conclusions": "In Section 3 we derived  the fully relativistic two-dimensional form that the differential Bethe-Salpeter equation for the electron-positron system takes in the limit of infinite constant and homogeneous magnetic field imposed on the system. We studied the f\\al~ phenomenon inherent in this equation basing on exactly relativistic treatment of the relative motion of the electron and positron. Thanks to this phenomenon, at a certain finite value of the magnetic field (\\ref{final}) called here the \\textit{first hypercritical value}  the positronium level deepens so much that the rest energy of the system is completely compensated for by the mass defect. The most symmetrical solution of \\BS ~ corresponding to the center-of-mass momentum equal to zero may be attributed to the vacuum. In Subsection 4.1 we described the vacuum restructuring that takes place after the magnetic field exceeds the first hypercritical value in terms of formation of localized states of the pair, which are either \"confined\" or tightly bound - depending on whether the theory of the f\\al ~ in Refs. \\cite{shabad}, \\cite{shabad2} is appealed to or not. We estimate in Subsection 4.2 the modification of the first hypercritical value of the magnetic field that may be introduced by the mass corrections to the Dirac field propagator in the strong magnetic field. In Subsection 4.3 we discuss the \\textit{second hypercritical value} of the magnetic field where a lattice appears in the vacuum and the latter becomes unstable under the delocalization of the states of the pair, the delocalized charged particles on the Larmour orbits being capable of screening the external field and thus setting a limit to its growth.  The above limiting values are obtained within pure quantum electrodynamics. Up to now, it was accepted that the vacuum of this theory is stable with any magnetic field, contrary to electric field and contrary to non-Abelian gauge field theories like QCD. In spite of the huge values, expected  to be present, perhaps, only in superconducting cosmic strings \\cite{witten}, the values obtained may be important as setting the limits of applicability of QED.  As being due to the special, non-perturbational mechanism described above, the hypercritical field is determined by the inverse square root of the fine structure constant elevated to the exponent. This makes it hundred or so orders of magnitude  smaller than other known typical values \\cite{ritus2} of the magnetic field that may be expected to lead us beyond  the scope of coverage of QED owing to the lack of asymptotic freedom. For instance \\cite{shabad3}, the photon becomes a tachyon in the magnetic field of the order of $B_0\\exp (3\\pi/\\alpha)$."
    },
    "0512/hep-ph0512320_arXiv.txt": {
        "abstract": "We classify dark energy models in a plane of observables that correspond to the common parameterization of a non-constant equation of state, $w(a)=w_0+w_a(1-a)$, where $a$ is the scale factor of the universe. The models fall into four classes and only two of these classes have a region of overlap in the observable plane. We perform a joint analysis of all Type Ia supernova (SNIa) data compiled by the High-Z SN Search Team (HZT) and the Supernova Legacy Survey (SNLS) and find that no class of models is excluded by current SNIa data. However, an analysis of large scale structure, Ly$\\alpha$ forest and bias constraints from SDSS, the Gold SNIa data and WMAP data indicates that non-phantom barotropic models with a positive sound speed are excluded at the 95\\%~C.~L. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/nucl-th0512098_arXiv.txt": {
        "abstract": "We apply the recently developed $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$ test of compact star cooling theories for the first time to hybrid stars with a color superconducting quark matter core. While there is not yet a microscopically founded superconducting quark matter phase which would fulfill constraints from cooling phenomenology, we explore the hypothetical 2SC+X phase and show that the magnitude and density-dependence of the X-gap can be chosen to satisfy a set of tests: temperature - age ($\\mathrm{T}$-$\\mathrm{t}$), the brightness constraint, $\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$, and the mass spectrum constraint. The latter test appears as a new conjecture from the present investigation. ",
        "introduction": "\\label{intro}  Recently, preparations for terrestrial laboratory experiments with heavy-ion collisons have been started, where it is planned to access the high-density/ low temperature region of the QCD phase diagram and explore physics at the phase boundary between hadronic and quark matter, e.g., within the CBM experiment at FAIR Darmstadt.  Predictions for critical parameters in  this domain of the temperature-density plane are uncertain since they cannot be checked against Lattice-QCD simulations which became rather precise at zero baryon densities. Chiral quark models have been developed and calibrated with these results. They can be extended into the finite-density domain and suggest a rich structure of color superconducting phases. These hypothetical phase structures shall imply consequences for the structure and evolution of compact stars, where the constraints from mass and radius measurements as well as from the cooling phenomenology have recently reached an unprecedented level of precision which allows to develop decisive tests of models for high-density QCD matter.  Among compact stars (we will address them also with the general term {\\it neutron stars} (NSs)) one can distiguish three main classes according to their composition: hadron stars, quark stars (bare surfaces or with thin crusts), and hybrid stars (HyS). The latter are the subject of the present study.  Observations of the surface thermal emission of NSs is one of the most promising ways to derive detailed information about processes in interiors of compact objects (see \\cite{Page:2005fq,Page:2004fy,Yakovlev:1999sk} for recent reviews). In \\cite{Popov:2004ey} (Paper I hereafter) we proposed to use a population synthesis of close-by cooling NSs as an addtional test for theoretical cooling curves. This tool, based on calculation of the $\\mathrm{Log\\, N}$-$\\mathrm{Log\\,S}$ distribution, was shown to be an effective supplement to the standard $\\mathrm{T}$~-~$\\mathrm{t}$ (Temperature vs. age) test. In Paper I we used cooling curves for hadron stars calculated in \\cite{Blaschke:2004vq}. Here we study cooling curves of HyS calculated in \\cite{Grigorian:2004jq} (Paper II hereafter).  Except $\\mathrm{T}$~-~$\\mathrm{t}$ and $\\mathrm{Log\\, N}$-$\\mathrm{Log\\,S}$ we use also the brightness constraint test (BC) suggested in \\cite{Grigorian:2005fd}. We apply altogether three tests -- $\\mathrm{T}$-$\\mathrm{t}$, $\\mathrm{Log\\, N}$-$\\mathrm{Log\\,S}$, and BC  -- to five sets of cooling curves of HyS. In the next section we describe calculation of these curves. In Section III we discuss the population synthesis scenario. After that we present our results which imply the conjecture of a new mass spectrum constraint from Vela-like objects. In Section 5 we discuss the results and present our conclusions in Section ~6. ",
        "conclusions": "We made a preliminary study of cooling curves for HyS based on the approach in Paper II which suggests that if quark matter occurs in a compact star, it has to be in the 2SC+X phase, where the hypothetical X-gap still lacks a microscopical explanation. All three tests of the cooling behavior ($\\mathrm{Log\\, N}$-$\\mathrm{Log\\, S}$, $\\mathrm{T}$-$\\mathrm{t}$, BC) are applied. Four models defined by the two-parameter ansatz for the X-gap were calculated. Model II with a density-independent X-gap could directly be excluded since it was not able to explain some cooling data including Vela at all. Two of the models (I and III) successfully passed two tests and marginally the third. None of these models could explain explain the temperature-age value for Vela within a typical mass range, i.e. for a Vela mass below 1.5 M$_\\odot$. However, with a steeper density dependence of the X-gap than suggested in Paper II, we were able to fulfill all 4 constraints and exemplified this for model IV. To conclude, HyS with a 2SC+X quark matter core appear to be good candidates to explain the cooling behaviour of compact objects, although a consistent theoretical explanation of the hypothetical X-gap and its steep density dependence has still to be developed."
    },
    "0512/astro-ph0512086_arXiv.txt": {
        "abstract": "For the case of Tycho's supernova remnant (SNR) we present the relation between the blast wave and contact discontinuity radii calculated within the nonlinear kinetic theory of cosmic ray (CR) acceleration in SNRs. It is demonstrated that these radii are confirmed by recently published Chandra measurements which show that the observed contact discontinuity radius is so close to the shock radius that it can only be explained by efficient CR acceleration which in turn makes the medium more compressible. Together with the recently determined new value $E_\\mathrm{sn}=1.2\\times 10^{51}$~erg of the SN explosion energy this also confirms our previous conclusion that a TeV \\gr flux of $(2-5)\\times 10^{-13}$~erg/(cm$^2$s) is to be expected from Tycho's SNR. {\\it Chandra} measurements and the {\\it HEGRA} upper limit of the TeV \\gr flux together limit the source distance $d$ to $3.3\\leq d\\leq 4$~kpc. ",
        "introduction": "Cosmic rays (CRs) are widely accepted to be produced in SNRs by the diffusive shock acceleration process at the outer blast wave (see e.g. \\cite{drury,be87,bk88,jel91,mald01} for reviews). Kinetic nonlinear theory of diffusive CR acceleration in SNRs \\cite{byk96,bv97} couples the gas dynamics of the explosion with the particle acceleration. Therefore in a spherically symmetric approach it is able to predict the evolution of gas density, pressure, mass velocity, as well as the positions of the forward shock and the contact discontinuity, together with the energy spectrum and the spatial distribution of CR nuclei and electrons at any given evolutionary epoch $t$, including the properties of the nonthermal radiation. The application of this theory to individual SNRs \\cite{bkv02,bkv03,bpv03,vbkr} has demonstrated its power in explaining the observed SNR properties and in predicting new effects like the extent of magnetic field amplification, leading to the concentration of the highest-energy electrons in a very thin shell just behind the shock.  Recent observations with the {\\it Chandra} and {\\it XMM-Newton} X-ray telescopes in space have confirmed earlier detections of nonthermal continuum emission in hard X-rays from young shell-type SNRs. With {\\it Chandra} it became even possible to resolve spatial scales down to the arcsec extension of individual dynamical structures like shocks \\cite{vl03,long,bamba}. The filamentary hard X-ray structures are the result of strong synchrotron losses of the emitting multi-TeV electrons in amplified magnetic fields downstream of the outer accelerating SNR shock \\cite{vl03,bkv03,bv04b,vbk05}. Such observational results gain their qualitative significance through the fact that these effective magnetic fields and morphologies turned out to be exactly the same as predicted theoretically from acceleration theory.  This theory has been applied in detail to Tycho's SNR, in order to compare results with the existing data \\cite{vbkr,vbk05}.  We have used a stellar ejecta mass $M_{ej}=1.4M_{\\odot}$, distance $d=2.3$~kpc, and interstellar medium (ISM) number density $N_H=0.5$~ H-atoms cm$^{-3}$.  For these parameters a total hydrodynamic explosion energy $E_{sn}=0.27\\times 10^{51}$~erg was derived to fit the observed size $R_s$ and expansion speed $V_s$.  A rather high downstream magnetic field strength $B_d\\approx 300$~$\\mu$G and a proton injection rate $\\eta=3\\times 10^{-4}$ are needed to reproduce the observed steep and concave radio spectrum and to ensure a smooth cutoff of the synchrotron emission in the X-ray region. We believe that the required strength of the magnetic field, that is significantly higher than the MHD compression of a $5~\\mu$G ISM field, has to be attributed to nonlinear field amplification at the SN shock by CR acceleration itself.  According to plasma physical considerations \\cite{lb00,belll01,bell04}, the existing ISM magnetic field can indeed be significantly amplified at a strong shock by CR streaming instabilities.  After adjustment of the predictions of the nonlinear sphe\\-ri\\-cally-symmetric model by a physically necessary renormalization of the number of accelerated CR nuclei to take account of the quasi-perpendicular shock directions in a SNR, very good consistency with the existing observational data was achieved.  Using {\\it Chandra} X-ray observations \\cite{warren05} have recently estimated the ratio between the radius $R_\\mathrm{c}$ of the contact discontinuity (CD), separating the swept-up ISM and the ejecta material, and the radius $R_\\mathrm{s}$ of the forward shock. The large mean value $R_\\mathrm{c}/R_\\mathrm{s}=0.93$ of this ratio was interpreted as evidence for efficient CR acceleration, which makes the medium between those two discontinuities more compressible.  Here we present the calculations of the mean ratio $R_\\mathrm{c}/R_\\mathrm{s}$, which are the unchanged part of our earlier considerations \\cite{vbkr,vbk05}, and demonstrate that these results, which are in fact predictions, fit the above measurements very well.  Since our calculations have been made in spherical symmetry they concern a priori an azimuthally averaged ratio $R_\\mathrm{c}/R_\\mathrm{s}$. We shall extend them by taking the effects of the Rayleigh-Taylor (R-T) instability of the CD into account. We shall in addition discuss a physical mechanism that leads to the observed {\\it azimuthal} variations of $R_\\mathrm{c}/R_\\mathrm{s}$. Finally we shall take the recent determination of the mechanical energy output $E_\\mathrm{sn}\\approx 1.2 \\times 10^{51}$ \\cite{bad05} which results from the theory of delayed detonations in the physics of type Ia SN explosions \\cite{gam04,gam05}, in order to predict a range for the \\gr spectrum that is consistent with the existing upper limits of high energy \\gr fluxes from Tycho's SNR. ",
        "conclusions": "Fig.1 and partly Fig.2 show the calculations of shock and CD related quantities which were part of our earlier considerations \\cite{vbkr,vbk05}. The calculated shock as well as CD radii and speeds are shown as a function of time for the two different cases of interior magnetic field strengths $B_\\mathrm{d}=240$~$\\mu$G and $B_\\mathrm{d}=360$~$\\mu$G considered, together with the azimuthally averaged experimental data available at the time. \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{ksenofontov2_fig1.eps} \\caption{(a) Shock radius $R_\\mathrm{s}$, contact discontinuity radius $R_\\mathrm{c}$, shock speed $V_\\mathrm{s}$, and contact discontinuity speed $V_\\mathrm{c}$, for Tycho`s SNR as functions of time, including particle acceleration; (b) total shock ($\\sigma$) and subshock ($\\sigma_\\mathrm{s}$) compression ratios. The {\\it dotted vertical line} marks the current epoch. The {\\it solid and dashed lines} correspond to the internal magnetic field strength $B_\\mathrm{d}=240$~$\\mu$G and ($B_\\mathrm{d}=360$~$\\mu$G), respectively.  The observed mean size and speed of the shock, as determined by radio measurements \\cite{tg85}, are shown as well.} \\label{f1} \\end{figure}  According to Fig.1a Tycho is nearing the adiabatic phase. To fit the spectral shape of the observed radio emission we assumed a proton injection rate $\\eta=3\\times 10^{-4}$. This leads to a significant nonlinear modification of the shock at the current age of $t=428$~yrs. A larger magnetic field lowers the Alfv\\'enic Mach number and therefore leads to a decrease of the shock compression ratio, as seen in Fig.1b. The result is a total compression ratio $\\sigma=5.7$ and a subshock compression ratio $\\sigma_s=3.5$ for $B_\\mathrm{d}=240$~$\\mu$G. In turn $\\sigma=5.2$, $\\sigma_s=3.6$, for $B_\\mathrm{d}=360$~$\\mu$G.  Therefore, as can be seen from Fig.2, including CR acceleration at the outer blast wave, the calculated value of the ratio $R_\\mathrm{c}/R_\\mathrm{s}$ for $B_\\mathrm{d}=360$~$\\mu$G is slightly lower than for $B_\\mathrm{d}=240$~$\\mu$G. At the current epoch we have $R_\\mathrm{c}/R_\\mathrm{s}\\approx 0.90$ which is lower than the value $R_\\mathrm{c}/R_\\mathrm{s}=0.93$ inferred from the observations. Qualitatively our result goes in the same direction as calculations by \\cite{blon01} which modeled SNRs with a uniform specific heat ratio $\\gamma_\\mathrm{eff} < 5/3$, for the circumstellar medium and the ejecta material alike.  Projecting a highly structured shell onto the plane of the sky tends to favor protruding parts of the shell. Therefore the average radius measured in projection is an overestimate of the true average radius. Analysing the amount of bias from the projection for the shock and CD radii \\cite{warren05} found a corrected \"true\" value $R_\\mathrm{c}/R_\\mathrm{s}=0.93$ which is lower than their measured ``projected average'' value $R_\\mathrm{c}/R_\\mathrm{s}=0.96$, as a result of the above geometrical effect.  In turn, starting from a spherically symmetric calculation of the CD radius, as we do, one has to take into account that the actual CD is subject to the R-T instability. In the nonlinear regime it leads to effective mixing of the ejecta and swept-up ISM material with ``fingers'' of the ejecta on top of this mixing region, which extend farther into the shocked gas than the radius $R_\\mathrm{c}$ predicted when assuming spherical symmetry e.g. \\cite{chevetal92,dwark00,blon01,wang01}. Therefore our ratio $R_\\mathrm{c}/R_\\mathrm{s}=0.90$, calculated within the spherically symmetric approach, has to be corrected for this effect in order to compare it with the measured value $R_\\mathrm{c}/R_\\mathrm{s}=0.93$.  In the case when all the fingers have length $l$ and occupy half of the CD surface, one would have a mean CD size $R_c'\\approx R_c+0.5l$ which has to be compared with $0.93R_\\mathrm{s}$.  According to the numerical modelling of \\cite{wang01}, albeit without particle acceleration, the R-T instability allows fingers of ejecta to protrude beyond the spherically symmetric CD radius by 10\\%.  The longest fingers of size $l\\approx 0.1R_c$ occupy less than 50\\% of the CD surface.  However, in projection they stick out of the mixing region, whose thickness is roughly $0.5l$.  This leads to a rough estimate of the corrected CD radius $R'_c = 1.05R_c$ which has to be compared with the experimentally estimated value.  The comparison of the corrected values $R'_\\mathrm{c}/R_\\mathrm{s}$, according to our earlier calculations as well as for different assumptions (see below) about the explosion energy and source distance, with this experimentally estimated value $R_\\mathrm{c}/R_\\mathrm{s}=0.93$ (in Fig.2 we present that value with 2\\% uncertainties, according to \\cite{warren05}) shows quite good agreement (see Figs.2 and 3) even if one takes into account some uncertainty in the quantitative determination of our correction factor $R'_\\mathrm{c}/R_\\mathrm{c}$, which in our view lies in the range 1.03-1.07. \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{ksenofontov2_fig2.eps} \\caption{ The ratio $R_\\mathrm{c}/R_\\mathrm{s}$ of the radii of the contact discontinuity and the forward shock  as a function of time. {\\it Solid and dashed} lines correspond to the same two cases as in Fig.1. The lines of all other styles correspond to the SN explosion energy $E_\\mathrm{sn}=1.2\\times 10^{51}$~erg and four different distances (see the caption of Fig.3) {\\it Thin} lines represent the values calculated in the spherically symmetric model, whereas the {\\it thick} lines show the values $R'_\\mathrm{c}/R_\\mathrm{s}$ which contain the correction for the effect produced by the R-T instability. The experimental point is taken from \\cite{warren05}.} \\label{f2} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{ksenofontov2_fig3.eps} \\caption{{\\em Top panel}: the forward shock radius $R_\\mathrm{s}$ and {\\em bottom panel}: the ratio $R_\\mathrm{c}/R_\\mathrm{s}$ of the radii of the contact discontinuity and the forward shock as a function of azimutal angle \\cite{warren05}. The regions where $R_\\mathrm{c}/R_\\mathrm{s}> 0.99$ are shadowed. } \\label{w} \\end{figure}  Another interesting peculiarity of Tycho's shock structure, which we would like to discuss here, is the quite irregular behaviour of the radius of the forward shock around the edge of the visible SNR disk, that is clearly seen in Fig.\\ref{w}. Large shock distortions of this kind are not expected to result from the R-T instability. At first sight such a variation of the values of $R_\\mathrm{s}$ and $R_\\mathrm{c}$ can be easily attributed to fluctuations of the ambient ISM density and/or to an inhomogeneously distributed density and velocity field of the ejected matter. The local shock part, which encounters a lower ISM density or has a faster ejecta portion behind, propagates faster compared with the neighbouring shock pie\\-ces. This will lead to the formation of local forward displacements of the CD and the forward shock whose number and relative sizes are determined by the specific structure of the ISM and/or ejecta. Since a larger shock compression ratio is expected for higher shock speed (see Fig.1b), one should also expect a smaller difference between $R_\\mathrm{s}$ and $R_\\mathrm{c}$ on the top region of each such displacement. This is exactly what is observed.  Such arguments would give a good explanation for the observed picture only if CR injection/acceleration took place uniformly across the entire shock surface. However, in our view the actual situation is expected to be more complicated -- and physically more interesting. Efficient injection of suprathermal nuclear particles into the acceleration process takes place in those local shock regions, where the forward shock is quasi-parallel, and these regions are distributed over the shock surface according to the ambient ISM magnetic field structure and occupy in total about 20\\% of the shock \\cite{vbk03}. As a result one would expect that only about 20\\% of the local forward shock displacements efficiently accelerate nuclear CRs, and therefore display the extraordinary high ratio $R_\\mathrm{c}/R_\\mathrm{s}$. In reality the most extreme values $R_\\mathrm{c}/R_\\mathrm{s}\\geq 0.99$ are observed on randomly distributed local shock regions, and the positions of these regions roughly coincide with the positions of extended shock displacements. To explain such a correlation a strong physical connection between the shock speed and the efficiency of CR injection/acceleration should exist: the local parts of the shock with large speed effectively produce CRs and vice versa.  There are at least two physical processes which can resolve the above problem. The first one gives high injection on the leading part of the displacements, if they are formed due to ISM and/or ejecta inhomogeneities and initially the injection was suppressed here as a result of a highly oblique magnetic field.  Since after its formation the leading part of the displacement develops a large curvature, it will have a significant portion which becomes quasi-parallel, and therefore efficient CR injection/acceleration is expected over the whole leading part of such a displacement.  The second factor can presumably itself lead to the formation of displacements, even in the case of a uniform ISM and a spherically symmetric ejecta distribution. A significant fraction of the internal energy behind a quasi-parallel part of the shock front is contained in CRs. Since the shock is expected to be modified, the CR spectrum is very hard and therefore the pressure is mainly carried by the CRs with the highest energies. With their high mobility these CRs can diffuse laterally into neighbouring downstream volumes, which are located behind shock surfaces not producing CRs. This diffusive loss leads to a decrease of the internal pressure behind the quasi-parallel parts of the shock. Therefore the corresponding downstream ejecta undergo less deceleration. As a result these ejecta accelerate relative to their surroundings. This leads to the formation of a local displacement. One can expect that this also distorts the initially spherically symmetric shock. We therefore conclude that the effectively accelerating parts of the forward shock are among those regions, where the displacements occur, and they may actually delineate them. Predominantly fingers are situated on the effectively accelerating parts of the shock surface. If this is correct it gives a unique experimental identification of the areas at the outer SNR shock which efficiently produce nuclear CRs.   Since the most outwardly displaced parts of the CD and the shock with a high ratio $R_\\mathrm{c}/R_\\mathrm{s}$ have to be interpreted as the regions with efficient CR injection/acceleration, this opens the possibility to experimentally distinguish the shock areas with efficient CR production from those where CRs are not produced. A rough estimate shows that in Tycho's SNR the regions with extremely high ratios $R_\\mathrm{c}/R_\\mathrm{s}>0.99$ and the displacements occupy about 20\\% of the shock surface see Fig.4 of \\cite{warren05}. This corresponds to the theoretical expectation.   Our last point regards constraints which the recent re-evaluation of the mechanical output $E_\\mathrm{sn}=1.2 \\times 10^{51}$~erg together with the {\\it HEGRA} upper limit for the TeV \\gr flux \\cite{aha01} approximately impose on the distance and ambient density for Tycho's SNR. With this new $E_\\mathrm{sn}$-value we find a consistent fit for all existing data -- the SNR size, its expansion rate, overall synchrotron spectrum and the filament structure of the X-ray emission -- like it was done for the previously defined explosion energy $E_\\mathrm{sn}=0.27 \\times10^{51}$~erg (see \\cite{vbkr,vbk05} for details). In particular, the fit of $R_\\mathrm{s}$ and $V_\\mathrm{s}$, which is of the same quality as in Fig.1 for $E_\\mathrm{sn}=0.27 \\times10^{51}$~erg, gives for each assumed distance $d$ a rather definite value of the ISM number density $N_\\mathrm{H}(d)$. A rough explanation of this numerical result is the following: since for a given experimentally measured angular SNR size and its expansion rate the linear size $R_\\mathrm{s}$ and the speed $V_\\mathrm{s}$ scale proportionally to distance $d$, and since in the nearby (in time) Sedov phase $R_\\mathrm{s}\\propto (E_\\mathrm{sn}/N_\\mathrm{H})^{1/5}$, the density $N_\\mathrm{H}\\propto E_\\mathrm{sn}/d^5$ decreases with increasing distance $d$. The hadronic \\gr flux $F_{\\gamma}\\propto R_\\mathrm{s}^3V_\\mathrm{s}^2 N_\\mathrm{H}^2/d^2$ is then expected to scale as $F_{\\gamma}\\propto E_\\mathrm{sn}^2/d^7$.  The ejected mass is still assumed to be $M_{\\mathrm{ej}}=1.4 M_{\\odot}$. We also find the same nuclear injection rate $\\eta=3\\times 10^{-4}$ for all cases, and downstream magnetic field values $B_\\mathrm{d}\\approx 400$~$\\mu$G. At the same time, the linear size $L$ of an X-ray filament increases proportional to $d$. Therefore the magnetic field strength $B'_\\mathrm{d}\\propto L^{-2/3}$ \\cite{bv04b}, determined from the filament sizes (see Introduction), decreases with $d$. \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{ksenofontov2_fig4.eps} \\caption{ Spectral energy distribution of the \\gr emission from Tycho's SNR, as a function of \\gr energy $\\epsilon_{\\gamma}$, for a mechanical SN explosion energy of $E_\\mathrm{sn}= 1.2 \\times 10^{51}$~erg and four different distances $d$ and corresponding values of the ISM number densities $N_\\mathrm{H}$. All cases have a dominant hadronic compared to Inverse Compton \\gr flux. Experimental data are the upper limits of the {\\it HEGRA} (H-CT; \\cite{aha01}) and {\\it Whipple} (W; \\cite{buckley98}) Cherenkov telescopes and the 95\\% confidence {\\it HEGRA AIROBICC} (HA; \\cite{prahl97}) upper limit.} \\label{f3} \\end{figure}  In order to find the constraint on the distance $d$ and the ISM density $N_\\mathrm{H}$, we then compare in Fig.3 the resulting \\gr spectral energy distribution with the {\\it HEGRA} and Whipple upper limits at TeV energies. It is seen that all distances $d < 3.3$~kpc are inconsistent with the {\\it HEGRA} data. Distances of $d>4$~kpc are still consistent with the \\gr data, although there is a growing discrepancy between $B_\\mathrm{d}$ and $B'_\\mathrm{d}$: at $d=4.5$~kpc $B'_\\mathrm{d}\\approx 300$~$\\mu$G which is already considerably smaller than $B_\\mathrm{d}\\approx 400$~$\\mu$G. Therefore we believe that we can constrain the source distance also from above, $d<4$~kpc.  In Fig.2 we also show the values for $R_\\mathrm{c}/R_\\mathrm{s}$ and $R'_\\mathrm{c}/R_\\mathrm{s}$ for the case $E_\\mathrm{sn}=1.2 \\times 10^{51}$~erg and these increased distances. Within our approximate determination of $R'_\\mathrm{c}$ from $R_\\mathrm{c}$ they still agree with the {\\it Chandra} data, in particular because the CR production rates are comparable. Our calculations of the \\gr emission lead us to predict that the new Northern Hemisphere TeV detectors should detect this source at TeV-energies in, predominantly, hadronic $\\gamma$-rays: the expected $\\pi^0$-decay $\\gamma$-ray energy flux $(2-5)\\times 10^{-13}$~erg/(cm$^2$s) extends up to almost 100~TeV if the distance is indeed within the range 3.3-4 kpc. As a corollary the detection of a TeV signal is not only important by itself, but it is also crucial for the correct determination of all other key Supernova parameters."
    },
    "0512/gr-qc0512008_arXiv.txt": {
        "abstract": "In this work, numerical simulations were used to investigate the gravitational stochastic background produced by coalescences occurring up to $z \\sim 5$ of double neutron star systems. The cosmic coalescence rate was derived from Monte Carlo methods using the probability distributions for forming a massive binary and to occur a coalescence in a given redshift. A truly continuous background is produced by events located only beyond the critical redshift $z_* = 0.23$. Events occurring in the redshift interval $0.027<z<0.23$ give origin to a ``popcorn\" noise, while those arising closer than $z = 0.027$ produce a shot noise. The gravitational density parameter $\\Omega_{gw}$ for the continuous background reaches a maximum around 670 Hz with an amplitude of $1.1\\times 10^{-9}$, while the ``popcorn\" noise has an amplitude about one order of magnitude higher and the maximum occurs around a frequency of 1.2 kHz. The signal is below the sensitivity of the first generation of detectors but could be detectable by the future generation of ground based interferometers. Correlating two coincident advanced-LIGO detectors or two EGO interferometers, the expected S/N ratio are respectively 0.5 and 10. ",
        "introduction": "The merger of two neutron stars, two black holes or a black hole and a neutron star are among the most important sources of gravitational waves (GW), due to the huge energy released in the process. In particular, the coalescence of double neutron stars (DNS) may radiate about 10$^{53}$ erg in the last seconds of their inspiral trajectory, at frequencies up to 1.4-1.6 kHz, range covered by most of the ground-based laser interferometers like VIRGO \\citep{bra}, LIGO \\citep{abr}, GEO \\citep{hou} or TAMA \\citep{kur}. Besides the amount of energy involved in these events, the rate at which they occur in the local universe is another parameter characterizing if these mergings are or not potential interesting sources of GW. In spite of the large amount of work performed in the past years, uncertainties persist in estimates of the DNS coalescence rate. In a previous investigation, we have revisited this question \\citep{dfp,reg05}, taking into account the galactic star formation history derived directly from observations and including the contribution of elliptical galaxies when estimating the mean merging rate in the local universe. Based on these results, we have predicted a detection rate of one event every 125 and 148 years by initial LIGO and VIRGO respectively and up to 6 detections per year in their advanced configurations.  Besides the emission produced by the coalescence of the nearest DNS, the superposition of a large number of unresolved sources at high redshifts will produce a stochastic background of GW. In the past years, different astrophysical processes susceptible to generate a stochastic background have been investigated. On the one hand, distorted black holes \\citep{fer99,ara}, bar mode emission from young neutron stars \\citep{reg} are examples of sources able to generate a shot noise (time interval between events large in comparison with duration of a single event), while supernovas or hypernovas \\citep{bla,cow01,cow02,buo} are expected to produce an intermediate ``popcorn\" noise. On the other hand, the contribution of tri-axial rotating neutron stars \\citep{reg01}, including magnetars \\citep{rp05}, constitutes a truly continuous background.  Populations of compact binaries as, for instance, the cataclismic variables are responsible for the existence of a galactic background of GW in the mHz domain, which could represent an important source of confusion noise for space detectors as LISA \\citep{eva,hils,ben,pos,nel,tim}. These investigations have been extended recently to the extra-galactic contribution. \\citet{sch}, \\citet{kim04} and \\citet{coo} considered cosmological populations of double and mixed systems involving black holes, neutron stars and white dwarfs, while close binaries originated from low and intermediate mass stars were discussed by \\citet{far}.  In this work, using the DNS merging rate estimated in our precedent study, we have estimated the gravitational wave background spectrum produced by these coalescences. Numerical simulations based on Monte Carlo methods were performed in order to determine the critical redshift $z_*$ beyond which the duty cycle condition required to have a continuous background ($D > 1$) is satisfied. Unlike previous studies which focus their attention on the early low frequency inspiral phase covered by LISA \\citep{sch,far,coo}, here we are mainly interested in the few thousand seconds before the last stable orbit is reached, when more than 96\\% of the gravitational wave energy is released. The signal frequency is in the range 10-1500 Hz, covered by ground based interferometers. The paper is organized as follows. In \\S2, the simulations are described; in \\S3 the contribution of DNS coalescences to the stochastic background is calculated; in \\S4 the detection possibility with laser beam interferometers is discussed and, finally, in \\S5 the main conclusions are summarized. ",
        "conclusions": "In this work, we have performed numerical simulations using Monte Carlo techniques to estimate the occurrence of double neutron star coalescences and the gravitational stochastic background produced these events. Since the coalescence timescale obeys a well defined probability distribution ($P(\\tau) \\propto 1/\\tau$), derived from simulations of the evolution of massive binaries \\citep{dfp}, the cosmic coalescence rate does not follow the cosmic star formation rate and presents necessarily a time-lag. In the case where the sources are supernovae or black holes, the gravitational burst is produced in a quite short timescale after the the formation of the progenitors. Therefore, the time-lag is negligible and the comoving volume where the progenitors are formed is practically the same as that where the gravitational wave emission occurs, introducing a considerable simplification in the calculations. This is not the case when NS-NS coalescences are considered, since timescales comparable or even higher than the Hubble timescale have non negligible probabilities. The maximum probability to form a massive binary occurs at $z \\sim 1.7$, depending slightly on the adopted cosmic star formation rate, whereas the maximum probability to occur a coalescence is around $z \\sim 1.4$.  We have found that a truly continuous background is formed only when sources located beyond $z > 0.23$ ($z > 0.27$ for the SFR1 case), including 96\\% (94\\% for SFR1) of all events and the critical redshift corresponds to the condition $D > 1$. Sources in the redshift interval $0.027 < z < 0.23$  ($0.032 < z < 0.27$ for SFR1) produce a ``popcorn\" noise. Our computations indicate that the density parameter $\\Omega_{gw}$ has a maximum around 670 Hz (630 Hz for SFR1), attaining an amplitude of about of $1.1 \\times 10^{-9}$ ($8.3 \\times 10^{-10}$ for SFR1). The low frequency cutoff around 1.2 kHz corresponds essentially to the gravitational redshifted wave frequency associated to last stable orbit of sources located near the maximum of the coalescence rate.  The computed signal is below the sensitivity of the first and the second generation of detectors. However, using the planned sensitivity of third generation interferometers, we found that after one year of integration, the cross-correlation of two EGO like coincident antennas, gives an the optimized signal-to-noise of $S/R \\sim 10$.  The ``popcorn\" contribution is one order of magnitude higher with a maximum of $\\Omega_{gw} \\sim 1.3 \\times 10^{-8}$ ($8.8 \\times 10^{-9}$ for SFR1) at $\\sim 1.2$ kHz. This signal, which is characterized by the spatial and temporal evolution of the events as well as by its signature, can be distinguished from the instrumental noise background and adequate data analysis strategies for its detection are currently under investigation \\citep{all,cor,dra,cow05}.  {\\bf Acknowledgement}  The authors thanks the referee by his useful comments, which have improved the early version of this paper."
    },
    "0512/astro-ph0512279_arXiv.txt": {
        "abstract": "A new method of measuring cosmology with gamma-ray bursts(GRBs) has been proposed by Liang and Zhang recently. In this method, only observable quantities including the rest frame peak energy of the $\\nu F_{\\nu}$ spectrum $(E_{p}^{'})$, the isotropic energy of GRB $(E_{\\gamma,iso})$, and the break time of the optical afterglow light curves in the rest frame $(t_{b}^{'})$ are used. By considering this method we constrain the cosmological parameters and the redshift at which the universe expanded from the deceleration to acceleration phase. We add five recently-detected GRBs to the sample and derive $E_{\\gamma, {\\rm{iso}}}/10^{52} {\\rm ergs}=(0.93\\pm0.25)\\times (E^{'}_{\\rm {p}}/{\\rm 100 \\ keV})^{1.91\\pm0.32}\\times(t^{'}_{\\rm{b}}/1{\\rm day})^{-0.93\\pm0.38}$ for a flat universe with $\\Omega_M=0.28$ and $H_0=71.0$ km s$^{-1}$ Mpc$^{-1}$. This relation is independent of the medium density around bursts and the efficiency of conversion of the explosion energy to gamma-ray energy. We regard the $E_{\\gamma, {\\rm{iso}}}(E_{\\rm {p}}^{'}, t_{\\rm{b}}^{'}$) relationship as a standard candle and find $0.05<\\Omega_{\\rm{M}}<0.48$ and $\\Omega_{\\Lambda}<1.15$ (at the $1\\sigma$ confidence level). In a flat universe with the cosmological constant we obtain $0.25<\\Omega_{\\rm{M}}<0.46$ and $0.54<\\Omega_{\\Lambda}<0.78$ at the $1\\sigma$ confidence level. The transition redshift is $z_{T}=0.69_{-0.12}^{+0.11}$. Combining 20 GRBs with 157 type Ia supernovae, we find $\\Omega_{M}=0.29\\pm0.03$ for a flat universe and the transition redshift is $z_{T}=0.61_{-0.05}^{+0.06}$, which is slightly larger than the value found by considering SNe Ia alone. In particular, We also discuss several dark-energy models in which the equation of state $w(z)$ is parameterized, and investigate constraints on the cosmological parameters in detail. ",
        "introduction": "The property of dark energy and the physical cause of acceleration of the present universe are two of the most difficult problems in modern cosmology. In past several years, many authors used distant type Ia supernovae (SNe Ia) (Riess et al. 1998; Perlmutter et al. 1999; Riess et al. 2004), cosmic microwave background (CMB) fluctuations (Bennett et al. 2003; Spergel et al. 2003), and large-scale structure (LSS) (Tegmark et al. 2004) to explore cosmology. Very recently, there have been extensive discussions on using gamma-ray bursts (GRBs) to constrain cosmological constraints (Dai et al. 2004; Ghirlanda et al. 2004; Xu et al. 2005; Firmani et al. 2005; Friedman \\& Bloom 2005; Mortsell \\& Sollerman 2005; Di Girolamo et al. 2005; Liang \\& Zhang 2005; Lamb et al. 2005).  SNe Ia have been considered as astronomical standard candles and used to measure the geometry and dynamics of the universe. Phillips (1993) found the intrinsic relation in SNe Ia: $L_{p}=a\\times\\bigtriangleup{m}_{15}^{b}$, where $L_{p}$ is the peak luminosity and $\\bigtriangleup{m}_{15}$ is the decline rate in the optical band at day 15 after the peak. This relation and other similar relations can be used to explore cosmology. Riess et al (1998) considered 16 high-redshift supernovae and 34 nearby supernovae and found that our present universe has been accelerating. Perlmutter et al (1999) used 42 SNe Ia and drew the same conclusion. Riess et al (2004) selected 157 well-measured SNe Ia, which is called the``gold\" sample. Assuming a flat universe, they concluded that: (1) Using the strong prior of $\\Omega_{M}=0.27\\pm0.04$, fitting a static dark energy equation of state yields $-1.46<w<-0.78$ (95\\%C.L.). (2) Assuming a possible redshift dependence of $w(z)$ (using $w(z)=w_{0}+w_{1}z$), the data with the strong prior indicate that the region $w_{1}< 0$ and especially the quadrant ($w_{0}>-1$ and $w_{1}<0$) are the least favored. (3) Expand $q(z)$ into two terms: $q(z)=q_{0}+zdq/dz$. If the transition redshift is defined through $q(z_{T})=0$, they found $z_{T}=0.46\\pm0.13$. The cosmological use of SNe Ia has the following advantages: the SN Ia sample is very large and includes low-$z$ sources, so the parameters $a$ and $b$ can be calibrated by using low-$z$ SNe Ia. The Phillips relation and other similar relations are intrinsic and cosmology-independent so that they can be used to explore cosmology. But they also have disadvantages: the interstellar medium extinction may exist when optical photons propagate towards us. In addition, the maximum redshift of SNe Ia is only about 1.7 and thus the earlier universe may not be well-studied. Higher-redshift SNe Ia are necessary to eliminate parameter degeneracies in studying the evolution of dark energy (Weller \\& Albrencht 2002; Linder \\& Huterer 2003).  GRBs are the most intense electromagnetic explosions in the universe after the big bang. They have been well understood since the discovery of afterglows in 1997 (for review articles see Piran 1999, 2004; van Paradijs et al. 2000; M\\'esz\\'aros 2002; Zhang \\& M\\'esz\\'aros 2004). It has been widely believed that they should be detectable out to very high redshifts (Lamb \\& Reichart 2000; Ciardi \\& Loeb 2000; Bromm \\& Loeb 2002; Gou et al. 2004). Schaefer (2003) derived the luminosity distances of 9 GRBs with known redshifts by using two quantities (the spectral lag and the variability). He obtained the first GRB Hubble diagram with the mass density $\\Omega_M<0.35$ (at the $1\\sigma$ confidence level). Ghirlanda et al. (2004a) found the relation between isotropic-equivalent energy $E_{\\gamma,iso}$ and the local-observer peak energy $E_{p}^{'}$ (i.e., the Ghirlanda relation). Unfortunately, because of the absence of low-$z$ GRBs, the Ghirlanda relation has been obtained only from moderate-$z$ GRBs. So this relation is cosmology-dependent. Dai, Liang \\& Xu (2004) used for the first time the Ghirlanda relation with 12 bursts and found the mass density $\\Omega_M=0.35\\pm0.15$ (at the $1\\sigma$ confidence level) for a flat universe with the cosmological constant and the $w$ parameter of the static dark energy model $-1.27<w<-0.50$ ($1\\sigma$). Combining 14 GRBs with SNe Ia, Ghirlanda et al. (2004b) obtained $\\Omega_{\\rm M}=0.37\\pm0.10$ and $\\Omega_{\\Lambda}=0.87\\pm 0.23$. Assuming a flat universe, the cosmological parameters were constrained: $\\Omega_{\\rm M}=0.29\\pm0.04$ and $\\Omega_{\\Lambda}=0.71\\pm 0.05$. Firmani et al (2005) used the Bayesian method to solve the circularity problem. For a flat universe they found $\\Omega_{M}=0.28\\pm0.03$ and $z_{T}=0.73\\pm0.09$ for the combined GRB+SN Ia sample. In the dark energy model of $w_{z}=w_{0}$, they found $\\Omega_{M}=0.44$ and $w_{0}=-1.68$ with $z_{T}=0.40$ for the combined GRB+SN Ia sample. In the dark energy model of $w_{z}=w_{0}+w_{1}z$, they found the best values for GRB+SN Ia sample were $w_{0}=-1.19$ and $w_{1}=0.98$ with $z_{T}=0.55$. Xu, Dai \\& Liang (2005) obtained $\\Omega_{M}=0.15^{+0.45}_{-0.13}(1\\sigma)$ using 17 GRBs. Friedmann \\& Bloom (2005) discussed several possible sources of systematic errors in using GRBs as standard candles. Liang \\& Zhang (2005a) presented a multi-variable regression analysis to three observable quantities for a sample of 15 GRBs without assumption of any theoretical models. They obtained a relation among the isotropic gamma-ray energy ($E_{\\gamma,{\\rm{iso}}}$), the peak energy of the $\\nu F_{\\nu}$ spectrum in the rest frame ($E_{\\rm {p}}^{'}$), and the rest frame break time of the optical afterglow light curves ($t_{\\rm{b}}^{'}$). Using this relation, they found the $1\\sigma$ constraints are $0.13<\\Omega_{\\rm{M}}<0.49$ and $0.50<\\Omega_{\\Lambda}<0.85$ for a flat universe. They also obtained the transition redshift $0.78^{+0.32}_{-0.23}$ ($1\\sigma$). Ghirlanda et al (2005) used their relation in a homogeneous density profile and a wind density profile to explore cosmology. Liang \\& Zhang (2006) proposed an approach to calibrate the GRB luminosity indicators without introduction of a low-redshift GRB sample. The cosmological use of GRBs has advantages: First, gamma-ray photons suffer from no extinction when they propagate towards us. Second, GRBs are likely to occur at high redshifts. We can thus study the early universe. Recently GRB 050904 whose redshift is 6.29 was detected (Kawai et al. 2005; Haislip et al. 2005). But the low-$z$ GRB sample is so small that the intrinsic relation cannot be obtained. A cosmology-dependent relation is now used to constrain the cosmological parameters and transition redshift.  In this paper we investigate cosmological constraints and the transition redshift following the method of Liang \\& Zhang (2005). Our GRB sample contains 20 bursts. Combining this sample with 157 SNe Ia we constrain the cosmological parameters and the transition redshift in several dark energy models in detail. The structure of this paper is arranged as follows: In section 2, we list our sample and results from the regression analysis. In section 3, we explore the constraints on cosmological parameters and transition redshift using the GRB and SN Ia sample in different dark-energy models. In section 4, we present conclusions and a brief discussion. ",
        "conclusions": "Following Liang \\& Zhang (2005), we have derived an empirical relation among the rest frame peak energy of the $\\nu F_{\\nu}$ spectrum $(E_{p}^{'})$, the isotropic energy of GRB $(E_{\\gamma,iso})$, and the break time of the optical afterglow light curves in the rest frame $(t_{b}^{'})$ without any theoretical assumptions. The relation is $E_{\\gamma, {\\rm{iso}}}/10^{52} {\\rm ergs}=(0.93\\pm0.25)\\times (E^{'}_{\\rm {p}}/{\\rm 100 \\ keV})^{1.91\\pm0.32}\\times(t^{'}_{\\rm{b}}/1{\\rm day})^{-0.93\\pm0.38}$ in a flat universe with $\\Omega_M=0.28$ and $H_0=71.0$ km s$^{-1}$ Mpc$^{-1}$ for our sample with 20 GRBs. We find $0.05<\\Omega_{\\rm{M}}<0.48$ and $\\Omega_{\\Lambda}<1.15$ ($1\\sigma$). For a flat universe with the cosmological constant we obtain $0.25<\\Omega_{\\rm{M}}<0.46$ and $0.54<\\Omega_{\\Lambda}<0.78$ at the $1\\sigma$ confidence level. The transition redshift is $z_{T}=0.69_{-0.12}^{+0.11}$. Because of the small number of useful GRBs, we combine GRBs with SNe Ia. Using this joint sample, we find $\\Omega_{M}=0.29\\pm0.03$ for a flat universe. The transition redshift is $z_{T}=0.61_{-0.05}^{+0.06}$, which is slightly larger than the value found by considering SN Ia alone. The best values for $\\Omega_{M}$ and $w_{0}$ are $\\Omega_{M}=0.48$ and $w_{0}=-1.90$ with the GRB and SN Ia sample in the $w_{z}=w_{0}$ model. The transition redshift is $z_{T}=0.50_{-0.10}^{+0.09}$. For the SN Ia and GRB sample the best values are $w_{0}=-1.20_{-0.24}^{+0.18}$ and $w_{1}=0.85_{-0.25}^{+0.30}$ with $z_{T}=0.44_{-0.08}^{+0.09}$ at the $1\\sigma$ confidence level in the model of $w_{z}=w_{0}+w_{1}z$. The best values for the SN Ia and GRB sample are $w_{0}=-1.12_{-0.30}^{+0.32}$ and $w_{1}=1.90_{-1.75}^{+2.10}$ (at the $1\\sigma$ confidence level) with $z_{T}=0.49_{-0.09}^{+0.07}$ in the model of $w_{z}=w_{0}+w_{1}z/(1+z)$. The best values for the SN Ia and GRB sample are $w_{0}=-1.00_{-0.25}^{+0.22}$ and $w_{1}=0.54_{-1.20}^{+0.85}$ with $z_{T}=0.37_{-0.07}^{+0.10}$ at the $1\\sigma$ confidence level in the model of $w_{z}=w_{0}+w_{1}z/(1+z)^2$. The contours are tighter than those only by using the SN Ia data. From these figures a clear trend can be seen.  GRBs appear to be natural events of studying the universe at very high redshifts. They can bridge the gap between the nearby SNe Ia and the cosmic microwave background. GRBs establish a new insight on the cosmic acceleration. However the density of the circumburst medium and the efficiency of conversion of the explosion energy to gamma-rays should be assumed in the Ghirlanda relation. This is naturally overcome in the Liang \\& Zhang relation.  Since the empirical relation of Liang \\& Zhang (2005) is cosmology-dependent, we use a strategy through weighing this relation in all possible cosmological models. We find that the transition redshift varies from $0.37_{-0.07}^{+0.10}$ to $0.69_{-0.06}^{+0.05}$. Although our constraints from the GRB sample are weaker than those from SNe Ia, the constraints from the GRB and SN Ia sample are more stringent than those from the SN Ia sample. A low-$z$ GRB sample is needed to calibrate the relation in a cosmology-independent way. MIDEX-class mission, which would obtain $>$ 800 bursts in the redshift range $0.1\\leq z \\leq10$ during a 2-year mission (Lamb et al. 2005), will be dedicated to using GRBs to constrain the properties of dark energy. This burst sample would enable both $\\Omega_{M}$ and $w_{0}$ to be determined to $\\pm 0.07$ and $\\pm 0.06$ (68\\% C.L.), respectively, and $w_a$ to be significantly constrained. Probing the properties of dark energy by using GRBs is complementary (in the sense of parameter degeneracies) to other probes, such as CMB anisotropies and X-ray clusters (Lamb et al 2005). New constraints from GRBs detected in the future would improve the study of cosmology.   We thank Enwei Liang and Bing Zhang for helpful discussions, and the referee for valuable suggestions. This work was supported by the National Natural Science Foundation of China (grants 10233010 and 10221001)."
    },
    "0512/astro-ph0512423_arXiv.txt": {
        "abstract": "We describe several techniques developed by the High Resolution Fly's Eye experiment for measuring aerosol vertical optical depth, aerosol horizontal attenuation length, and aerosol phase function.  The techniques are based on measurements of side-scattered light generated by a steerable ultraviolet laser and collected by an optical detector designed to measure fluorescence light from cosmic-ray air showers. We also present a technique to cross-check the aerosol optical depth measurement using air showers observed in stereo.  These methods can be used by future air fluorescence experiments. ",
        "introduction": "Cosmic-ray air fluorescence experiments measure light produced by high-energy air showers.  The technique is calorimetric, with the atmosphere as the calorimeter.  The integrated scintillation light produced by an extensive air shower is proportional to the the primary particle energy and is almost independent of the primary particle composition. However, once this light is produced, the amount that reaches the observatory depends on how this light propagates from the shower through the molecular and aerosol components of the atmosphere to the light-collection optics of the detector.  The benefit of the atmosphere that makes the air fluorescence technique possible is balanced by the challenge to understand how the light propagates through the atmosphere.  Several aspects of the High Resolution Fly's Eye (HiRes) experiment address this challenge.  To reduce atmospheric effects, the experiment is located in a remote desert where the average humidity and cloud cover are relatively low.  The detector stations are located on hills approximately 100 m above the desert floor.  This places the optical detectors above many aerosols, including low-lying dust and ground fog.  To measure atmospheric effects, the experiment includes steerable uv lasers that probe the aperture of the detector while the experiment is running.  The light scattered out of a laser beam produces tracks in the same detectors that measure tracks of uv light from air showers.  Changes in atmospheric aerosols change the amount of light scattered and the amount of light that reaches the detector.  Modeling the detector, the laser, and the atmosphere generates a simulated set of detector measurements.  The simulated data are compared and fit to the actual detector measurements of the laser. This procedure yields a set of parameters that describe the aerosols. ",
        "conclusions": "We have developed several methods to measure the atmospheric aerosol properties of optical depth, extinction length, and phase function. These methods use side-scattered light from a 355 nm uv laser that produces tracks in the same optical fluorescence detectors that measure tracks from extensive air showers.  We have measured distributions of aerosol optical depth and horizontal extinction length at the location of the High Resolution Fly's Eye experiment. We found the distribution of aerosols is horizontally uniform over a distance of 20 km to less than 0.01 in $\\tau_{A}$ on more than 90\\% of the nights observed, and that most of these aerosols are in the lower 2 km of the atmosphere.  Between fall 1999 and spring 2001, the optical depth of the aerosol component is on average less than 25\\% of the optical depth of the molecular component.  As a cross check, we note that the average $\\tau_{A}$, as obtained by laser measurements, is consistent an average determined by comparing commonly viewed segments of extensive air showers observed by the two HiRes fluorescence detectors."
    },
    "0512/astro-ph0512109_arXiv.txt": {
        "abstract": "{We present a statistically decontaminated Color Magnitude Diagram of a $1\\degr\\times 1\\degr$ field in the core of the Sagittarius dSph galaxy. Coupling this CMD with the most recent metallicity distributions obtained from high resolution spectroscopy we derive robust constraints on the mean age of the stellar population that dominates the galaxy (Pop A). Using three different sets of theoretical isochrones in the metallicity range $-0.4\\le [M/H]\\le -0.7$ and taking into consideration distance moduli in the range $16.90\\le (m-M)_0\\le 17.20$ we find that the mean age of Pop A is larger than 5 Gyr, and the best-fit value is age$ = 8.0\\pm 1.5$ Gyr. Since Pop A provides the vast majority of the M giants that traces the tidal stream of Sgr dSph all over the sky, our estimate resolves the so called ``M giant conundrum'' (Majewski et al. 2003). The time needed by the M giants that currently populates the stream to diffuse within the main body of Sgr and to reach the extremes of the tidal tails once torn apart from the parent galaxy ($\\simeq 3-4$ Gyr) can be easily accommodated into the time lapsed since their birth ($\\simeq 5.5-9.5$ Gyr). ",
        "introduction": "The Sagittarius dwarf spheroidal galaxy (Sgr dSph; Ibata et al. \\cite{iba1}) currently provides the cleanest possibility to study in detail the tidal disruption and the accretion of a dwarf satellite into a large galaxy. The tidal tails of the disrupting galaxy have been observed in widely different positions with different tracers (see, for instance, Newberg et al. \\cite{sdss}, Ibata et al. \\cite{iba2}, and references therein). In particular, Majewski et al. (\\cite{maj03}, hereafter M03) showed that M giants traces the tidal tails of Sgr as a coherent and  dynamically cold filamentary structure (hereafter Sgr Stream) extending for tens of kpc from the parent galaxy and nicely aligned along the rosetta orbit of  Sgr (see also Law, Johnston \\& Majewski \\cite{law}).  The study of the physical properties (age and chemical abundances)  and the kinematics of stars into the main body of Sgr and into the Stream may provide for the first time the possibility to link the Star Formation History (SFH) of the galaxy with its ``dynamical/orbital'' history and evolution (see Bellazzini, Ferraro \\& Buonanno \\cite{BFBb}, hereafter BFBb) and/or to constrain the  one with the other. A first relevant case of this interlacing was pointed out by M03. These authors obtained rough estimates of the metallicity of M giants in the Sgr Stream from their J-K colors and using the Age-Metallicity Relation (AMR) for Sgr provided by Layden \\& Sarajedini (\\cite{ls00}, hereafter LS00), they concluded that Stream M giants should be younger than 5 Gyr and a significant fraction of them have an age of 2-3 Gyr. The N-body models that best reproduces the Sgr galaxy + Stream system (Law et al. \\cite{law}) show that a comparable amount of time is needed to produce tidal tails of the observed extension. According to these models, we are observing in the Stream M giants that were torn apart from their parent galaxy  up to 3.2 Gyr ago. Since it is reasonable to imagine that in a spheroidal system star formation episodes occur preferentially toward the center of the galaxy, the above time lapse should be (possibly) increased by the time needed to the newly born stars to diffuse out to the ``edges'' of the system (M03). Hence, the comparison  between the stellar evolution timescales (e.g., the age of M giants) and the dynamical timescales (e.g., the time needed to populate the whole extension of the Stream) give rise to a possible inconsistency, or at least a fine-tuning problem, since stars cannot be torn apart from a galaxy before their birth. M03 dubbed this apparent mismatch between the two timescales as \"the M giants conundrum''.  However, the method adopted by M03 to estimate the age of Sgr M giants is prone to uncertainties both in the previously established Sgr AMR as well as in the predicted colors of RGB tip stars from theoretical isochrones. An age estimate based on more reliable age-sensitive observables (such as the Main Sequence Turn Off - TO -  and/or the Sub Giant Branch, SGB) would greatly help in clarifying the issue. This kind of estimate can be performed only in the main body of the galaxy, where the stellar density is sufficient to allow the derivation of well populated Color Magnitude Diagrams (CMD) from the observation of reasonably small fields ($\\simgt 0.1$ deg$^2$). In this Letter we use a new large photometric dataset of Sgr stars to obtain a robust estimate of the mean age of the population that dominates the stellar content of Sgr, shedding a new light on the M giant {\\em conundrum}.  \\begin{figure} \\centering \\vskip 4truecm \\caption{ --- SEE fig1.jpg - POSTSCRIPT VERSION NOT INCLUDED ---- Statistically decontaminated CMD of the Sgr34 Field. The color of the stars is coded according to the local density of stars on the diagram. A few representative density contours are also overplotted. Some remarkable features of the CMD are labeled: the Red Clump (RC), the Blue Horizontal Branch (BHB), the RGB Bump, the Blue Plume (BP), the Main Sequence Turn Off point (TO) and the Sub Giant Branch (SGB). The spurious residuals of the statistic decontamination process (mainly due to edge effects) to the lower left and the upper right of the main locus of the Sgr populations in the CMD have been plotted as small points to provide a clearer view of the most significant parts of the diagram. } \\label{fig1} \\end{figure}    \\subsection{The dataset}  We have obtained B,V,I photometry of a  $1\\degr \\times 1\\degr$ field located at (l,b) $\\simeq(6.5\\degr,-16.5\\degr$), $\\sim 2\\degr$ to the East of the galaxy center, along the major axis (Sgr34 Field, see Bellazzini et al. \\cite{BFBa}, hereafter BFBa). The field was imaged with a mosaic of four pointings of the WFI camera mounted at the ESO/MPI 2.2m telescope at La Silla, Chile. The data reduction was performed as in Monaco et al. (\\cite{lbump}, hereafter Mo02) and the  absolute photometric calibration was achieved with repeated observations of Landolt's (\\cite{land}) standard fields. All the details of the data acquisition and reduction will be described in a future contribution (Bellazzini et al., in preparation). A $0.5\\degr\\times 0.5\\degr$ field sampling the Galactic population at similar angular distance from the Galactic Center (Gal Field, at (l,b)$\\simeq(-6.0\\degr,-14.5\\degr$)) was also observed with the same camera, to perform the statistical decontamination of the Sgr34 CMD from the foreground/background Galactic Stars  (see BFBb). The interstellar reddening was interpolated for each star from the Schlegel et al. (\\cite{cobe}) maps and corrected according to Bonifacio et al. (\\cite{bonir}). We found that the reddening variation over the considered fields are negligible (with standard deviations $<0.01$ mag) and we adopted the average reddening values $E(B-V)=0.116$ for Sgr34 and $E(B-V)=0.096$ for Gal Field.  The statistically decontaminated CMD obtained from the above described datasets is presented in Fig.~1. The statistical decontamination has been performed as in BFBb. The large samples available for both Sgr34 (more than 300000 stars) and for Gal Field ($\\simeq 57000$ stars) ensure a reliable and clean recovery of all the main features of the CMD, most of which have been labeled in Fig.~1 according to the nomenclature introduced in BFBa and Mo02. In the present context, the most  relevant features of Fig.~1 are the well identified TO point at $V_{TO}=21.4\\pm 0.15$ and the single, narrow and well defined SGB that is clearly visible at $V\\simeq 20.8$ and $B-V\\simeq 0.85$. In the following we will use these observables to constrain the mean age of Sgr.  \\begin{figure*} \\centering \\includegraphics[width=11.5cm]{fig2.ps} \\caption{Isochrone fitting of the TO/SGB region of the Sgr34 CMD, here represented with the isodensity contours shown in Fig.~1. The isochrones plotted as continuous lines provides the best fit to the TO luminosity, those plotted as dashed lines provides the best fit to the luminosity of the SGB. The left panels show the comparison with Z=0.004 isochrones, the right panel show the comparison with Z=0.007-0.008 isochrones. Isochrones from the P04 set (upper panels), G00 set (middle panels), and Y$^2$ set (lower panels) are considered. The age, metallicity and helium content (Y) of the adopted isochrones are reported in each panel.} \\label{isoc}% \\end{figure*} ",
        "conclusions": "The above described results robustly establishes that an age of 5 Gyr is not an upper limit for the bulk of Pop~A (and, consequently, of M giants in the main body of Sgr and in the Stream) as assumed by M03, but, in fact, a strong {\\em lower} limit to the age of most of these stars. Moreover, independently of the adopted theoretical models and of the assumed distance, the best fit age for this population lies in the range $5.5\\le$ age $\\le 9.5$ Gyr, and our preferred solution (Z=0.004 and $\\mu_0=17.10$, averaging over the results from the different isochrones sets) is $\\langle$age$\\rangle_{PopA} = 8.0\\pm 1.5$ Gyr. This {\\em evolutionary timescale} is now comfortably larger than the {\\em dynamical timescale} provided by realistic simulations of the formation of the Sgr Stream (Law et al. \\cite{law}), hence the {\\em M giant conundrum} appears to be solved."
    },
    "0512/astro-ph0512615_arXiv.txt": {
        "abstract": "Until recently, X-ray flares during the afterglow of gamma ray bursts (GRBs) were a rarely detected phenomenon, thus their nature is unclear.  During the afterglow of GRB 050502B, the largest X-ray flare ever recorded rose rapidly above the afterglow lightcurve detected by the {\\it{Swift}} X-ray Telescope.  The peak flux of the flare was $>500$ times that of the underlying afterglow, and it occurred at $>12$ minutes after the nominal prompt burst emission.  The fluence of this X-ray flare, (1.0 $\\pm$ 0.05) $\\times 10^{-6}$ erg cm$^{-2}$ in the 0.2--10.0 keV energy band, exceeded the fluence of the nominal prompt burst.  The spectra during the flare were significantly harder than those measured before and after the flare.  Later in time, there were additional flux increases detected above the underlying afterglow, as well as a break in the afterglow lightcurve.  All evidence presented below, including spectral and particularly timing information during and around the giant flare, suggests that this giant flare was the result of internal dissipation of energy due to late central engine activity, rather than an afterglow-related effect.  We also find that the data are consistent with a second central engine activity episode, in which the ejecta is moving slower than that of the initial episode, causing the giant flare and then proceeding to overtake and refresh the afterglow shock, thus causing additional activity at even later times in the lightcurve. ",
        "introduction": "Since its launch on 2004 November 20, {\\it{Swift}} \\citep{geh04} has provided detailed measurements of numerous GRBs and their afterglows with unprecedented reaction times.  Of the 57 bursts detected by the Burst Alert Telescope (BAT; Barthelmy et al. 2004) as of 2005 August 3, 43 were observed by the narrow field instruments in less than 200 ks (typical reaction time was much less, but occasionally BAT detected a burst that was observationally constrained).  The narrow field instruments include the X-ray telescope (XRT; Burrows et al. 2005) and the Ultraviolet-Optical Telescope (UVOT; Roming et al. 2005).  Of these 43 observations, 42 afterglows were detected by the XRT, and 30 of them received prompt ($< 300$ s) observations with the pointed instruments.  By detecting burst afterglows promptly, and with high sensitivity, the properties of the early afterglow and extended prompt emission can be studied in detail for the first time.  This also facilitates studies of the transition between the prompt emission and the afterglow.  While there are still many unknown factors related to the mechanisms that produce GRB emission, the most commonly accepted model is that of a relativistically expanding fireball with associated internal and external shocks \\citep{mes97}.  In this model, internal shocks produce the prompt GRB emission. Observationally, this emission typically has a timescale of $\\sim30$ s for long bursts and $\\sim$0.3 s for short bursts \\citep{mee96}.  The expanding fireball then shocks the ambient material to produce a broadband afterglow that decays quickly (typically as ${\\sim}t^{-\\alpha}$).  When the Doppler boosting angle of this decelerating fireball exceeds the opening angle of the jet into which it is expanding, then a steepening of the lightcurve (jet break) is also predicted \\citep{rho99}.  For a description of the theoretical models of GRB emission and associated observational properties, see \\citet{mes02}, \\citet{zha04}, \\citet{piran05}, and \\citet{van00}.  Several authors have suggested reasons to expect continued activity from the internal engine of the GRB after the classical \"prompt\" emission time frame.  \\citet{kat97} considered a model in which a magnetized disk around a central black hole could lead to continued energy release in the form of internal shocks.  The parameters of this energy release would depend on the complex configuration of the magnetic field and the magnetic reconnection dynamics, but time periods as long as days for the delayed emission were predicted.  \\citet{kin05} have speculated that episodic accretion processes could explain continued internal engine activity.  These authors expect that fragmentation and subsequent accretion during the collapse of a rapidly rotating stellar core could explain observations of extended prompt emission.  In general, the dominant model of an expanding fireball with internal/external shocks \\citep{mes97} allows for continued prompt emission, provided that the internal engine is capable of continuing the energy injection.  A few previous observations have included indications of flaring from GRBs after the nominal prompt emission phase.  \\citet{wat03} used XMM-Newton to detect line emission from GRB 030227 nearly 20 hours after the prompt burst.  They inferred continued energy injection at this late time, and concluded that a nearly simultaneous supernova and GRB event would require sporadic power output with a luminosity in excess of $\\sim5\\times10^{46}$ erg s$^{-1}$.  \\citet{pir05} used Beppo-SAX to observe two GRBs with relatively small X-ray flares.  The X-ray flare times for GRB 011121 and GRB 011211 were reported as t=240 s and t=600 s, respectively.  The spectral parameters of these two X-ray flares were consistent with afterglow parameters, and these flares were interpreted as the onset of the afterglow \\citep{pir05}.  Two other examples of flaring and/or late timescale emission can be found in \\citet{int03} and \\citet{gal05}.  Although not a detection of late flares from a particular GRB, the work of \\citet{con02}, in which an ensemble of GRBs was analyzed, should also be mentioned.  In this study, 400 long GRBs detected by the Burst and Transient Source Experiment (BATSE) were analyzed together in the form of a summed lightcurve above 20 keV.  Significant emission was found at late times (at least to 1000 s).  There are several possible explanations for this emission that do not require flares, but flares at various times are certainly one possible explanation.  More recently, \\citet{bur05b} provided the initial report that two bursts detected by {\\it{Swift}} showed strong X-ray flares.  The first of these, XRF 050406, was an X-ray flash with a short, and relatively weak, X-ray flare that peaked 213 s after the nominal prompt emission.  Due to the fast rise/decay, the most natural explanation for this flare is continued internal engine activity at late times (i.e. delayed prompt emission).  A detailed analysis of XRF 050406 is currently in preparation \\citep{romano05}.  GRB 050502B, the subject of this paper, was also reported on by \\citet{bur05b} since it had a dramatic X-ray flare that peaked 740 s after the nominal prompt emission.  This paper will now explore this event in more detail.  GRB 050502B was detected by the {\\it{Swift}}-BAT at 09:25:40 UT on 2005 May 02 \\citep{fal05}.  According to \\citet{cum05}, the T90 duration for the prompt emission detected by BAT was (17.5 $\\pm$ 0.2) s, and the burst had three individual peaks.  The main hard peak had a 6 s duration, was well fit by a power law with photon index 1.6 $\\pm$ 0.1, and had a 15--350 keV fluence of (8.0 $\\pm$ 1.0) $\\times10^{-7}$ erg cm$^{-2}$ \\citep{cum05}.  The spacecraft slewed promptly and observations with {\\it{Swift}}-XRT and {\\it{Swift}}-UVOT began 63 s after the BAT trigger time.  Since the flux was initially low, the XRT Image Mode data did not produce an initial onboard centroid position; however the first pass of data was analyzed on the ground, leading to an XRT position of RA $09^{h}30^{m}10.1^{s}$, Dec $+16^{\\circ}59^{m}44.3^{s}$ (J2000), with a 90\\% containment uncertainty of 5$\\arcsec$ \\citep{pag05}.  There was no counterpart found by UVOT, but ground-based optical observations reported by \\citet{cen05} revealed a fading afterglow at RA 09:30:10.02, Dec +16:59:48.07 (J2000), which is within 4$\\arcsec$ of the reported XRT position.  Following the initial low-flux detection by XRT, continued monitoring revealed increased flux that turned out to be the largest X-ray flare ever detected during a GRB afterglow.  This giant X-ray flare was not accompanied by any detected emission in the BAT energy band. ",
        "conclusions": "The complex lightcurve of GRB 050502B has many interesting features.  The most interesting of these is a giant flare with a fast rise/decay that began at 345 $\\pm$ 30 s, well after the nominal 17.5 $\\pm$ 0.2 s prompt GRB phase.  This is the largest X-ray flare ever detected after the apparent cessation of prompt emission.  After compiling all of the evidence, we come to the conclusion that the simplest explanation for this flare is continued activity in the internal engine of the GRB, not associated with the afterglow external shock.  This evidence includes: 1) The temporal decay index before and after the flare are identical, indicating that the afterglow had already begun before the flare, 2) the rise time and decay time of the flare are very fast, thus the flare is difficult (although not impossible) to explain with mechanisms associated with the external shock, 3) there is even faster time structure near the peak of the flare in the band above 1 keV, 4) the spectra during the giant flare are represented better by a Band function or cutoff power law model, rather than a simple power law, 5) the spectra before and after the flare are consistent with an afterglow that has already begun before the flare and continues with approximately the same spectral index after the flare, whereas the spectra during the flare is significantly harder, 6) based on recent data there appear to be many other similar flares, although not as intense as this one, with even faster time profiles (suggesting extended internal engine activity) in up to $\\sim$1/2 of the {\\it{Swift}} detected GRBs.  The two late bumps in the lightcurve (at $> 10^{4}$ s) are not as well constrained so conclusions are not as firm.  They could be due to either more internal engine activity, such as that which created the giant flare, or they could be due to another process associated with the afterglow shock.  If one or both of these bumps is due to re-energization of the afterglow shock due to the same internal engine ejection episode that created the giant flare at $\\sim$345 s, then this ejecta must have been emitted with a Lorentz factor $\\Gamma_2\\lesssim20$, and it must have been travelling slower than the primary ejecta that created the BAT detected prompt emission.  However, the limits on $\\Gamma_2$ are obviously dependent on the presumption that the bumps at $>10^4 s$ are due to re-energization of the afterglow shock by the same ejecta responsible for the giant flare.  Although this presumption is consistent with the data, it is not the only possible explanation, as stated earlier.  We can also conclude that there was a steepening of the lightcurve at some time during or after the late bumps (i.e. after $\\sim5\\times10^{4}$ s).  This steepening could have been a jet break, or it could have been the end of a phase of continuous energy injection into the afterglow shock front.  This GRB, and the associated X-ray flare, can be most easily explained within the framework of the standard GRB fireball model \\citep{mes97}, provided that there is some mechanism to feed the internal engine activity for extended time periods.  Although \"Type I\" collapsar models with prompt black hole formation cannot explain late time internal engine activity \\citep{mac01}, the fallback of material onto the central black hole after a stellar collapse could last for long time periods \\citep{woo93, mac01} and lead to late internal engine activity, albeit with significantly reduced luminosity.  Continued energy release due to dynamics of a magnetized disk around a black hole, as described by \\citet{kat97}, and/or continued and sporadic emission due to fragmentation and subsequent accretion during the collapse of a rapidly rotating stellar core, as described by \\citet{kin05}, could both explain observations of extended production of internal shocks.  In the time period since the analysis of the data from this GRB, {\\it{Swift}} has detected several more GRBs with X-ray flares \\citep{nou05}.  In the near future, it will be possible to use samples of many GRBs with X-ray flares to test models of long timescale prompt emission."
    },
    "0512/hep-ph0512234.txt": {
        "abstract": "\\noindent We analyze the constraints on neutrino mass spectra with extra sterile neutrinos as implied by the LSND experiment. The various mass related observables in neutrinoless double beta decay, tritium beta decay and cosmology are discussed. Both neutrino oscillation results as well as recent cosmological neutrino mass bounds are taken into account. We find that some of the allowed mass patterns are severely restricted by the current constraints, in particular by the cosmological constraints on the total sum of neutrino masses and by the non-maximality of the solar neutrino mixing angle. Furthermore, we estimate the form of the four neutrino mass matrices and also comment on the situation in scenarios with two additional sterile neutrinos. ",
        "introduction": "Introduction} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Scenarios with four neutrinos became popular on the wake of the LSND evidence of $\\bar{\\nu}_\\mu -\\bar{\\nu}_e$ transitions \\cite{lsnd}. Interpreted in terms of neutrino oscillations, the indicated mass scale for LSND is in the eV$^2$ range. Together with the evidence for neutrino oscillations from atmospheric (plus K2K) and solar (plus KamLAND) neutrino observations, requiring mass scales around $10^{-3}$ eV$^2$ and $10^{-4}$ eV$^2$, respectively, a fourth sterile neutrino has to be introduced in order to accommodate the presence of three distinct mass squared differences.  A priori, four neutrino scenarios allow for two possible mass patterns: \\begin{itemize} \\item[(i)] 2+2 scenarios, in which two pairs of neutrino states are separated from each other by the LSND mass scale. There are two possibilities for 2+2 scenarios; \\item[(ii)] 3+1 scenarios, in which one single neutrino state is separated by the LSND mass scale from the other three states. There are four possibilities for 3+1 scenarios; \\end{itemize} Oscillation analyzes in both schemes were performed by a number of authors \\cite{4mix1,4mix2,eitel,Giunti} and also the astrophysical and cosmological implications were investigated \\cite{4cosmo,4cosmo_new}. Historically, among the above two alternatives the 3+1 scenarios were at first relatively disfavored \\cite{4mix2} because of the non-observation of oscillations in short baseline experiments like KARMEN \\cite{karmen}, Bugey \\cite{bugey} and CDHS \\cite{cdhsw}. Therefore the 2+2 scenarios were found to be more compatible with the existing data. The sterile neutrino oscillation solution in 2+2 scenarios was viable for both solar and atmospheric neutrinos. However, SuperKamiokande data disfavored oscillation of the atmospheric $\\nu_\\mu$ to purely sterile neutrinos \\cite{SK_no_st}, and later on the SNO data started establishing the neutral current component in the solar $\\nu_e$ flux \\cite{sno}. For some time a mixed scenario, where the atmospheric neutrino anomaly is due to $\\nu_\\mu -\\nu_{s,\\tau}$ and the solar neutrino anomaly is due to $\\nu_e - \\nu_{s,\\tau}$, remained compatible with all data \\cite{sterilemix}. However, all recent analyzes show that 2+2 scenarios are ruled out at a high $\\sigma$ from the existing data \\cite{vallesterile,valle_rev}. Both atmospheric and solar neutrino data strongly disfavor oscillations to pure sterile species. This disfavored the 2+2 scenarios irrespective of whether LSND results are confirmed or not. The most updated analysis in the 3+1 scheme performed in \\cite{vallesterile,valle_rev} shows that non-evidence of neutrino oscillation in other short baseline (SBL) experiments combined with atmospheric neutrino data from SuperK and K2K is inconsistent with the LSND signal at 95\\% C.L.\\ and only marginal overlaps are found at 99\\% C.L. Thus, with increased precision of solar and atmospheric neutrino flux measurements the four neutrino explanation of the LSND anomaly suffered a setback. This led to many alternative explanations of the LSND anomaly including introduction of two sterile neutrinos -- the so-called 3+2 scenario \\cite{3plus2} --, CPT violation \\cite{cpt}, quantum decoherence effects violating CPT \\cite{decoh}, mass varying neutrinos \\cite{mv}, neutrino decay in four neutrino scenarios \\cite{4decay}, lepton number violating muon decay \\cite{babu}, decay of a heavy neutrino \\cite{sergio} or extra dimensional aspects \\cite{extra}.  Oscillation experiments can only measure the mass squared differences but not the absolute masses. The most direct and model independent way to measure the absolute masses is via kinematic measurements involving nuclear beta decay. The best bound at present is $m_{\\beta} < 2.3$ eV (95\\% C.L.) coming from the Mainz tritium beta decay experiment \\cite{mainz}. The KATRIN experiment is expected to increase the sensitivity down to $\\sim$ 0.2 eV \\cite{katrin}.   Information on absolute masses can also come from neutrinoless double beta decay (\\obb). Neutrinoless double beta decay experiments aim at observing the process \\[ (A,Z) \\rightarrow (A,Z + 2) + 2 \\, e^-~. \\] This is a lepton number violating process and its observation will establish the Majorana nature of neutrinos \\cite{scheval}. The decay width depends quadratically on the so-called effective mass. We assume here that only the light Majorana neutrinos implied by neutrino oscillation experiments are exchanged in the diagram of \\obb. In the basis in which the charged lepton mass matrix is real and diagonal, the effective mass is then nothing but the absolute value of the $ee$ element of the neutrino mass matrix. The best current limit on the effective mass is given by measurements of $^{76}$Ge established by the Heidelberg-Moscow collaboration \\cite{HM} (with similar results obtained by the IGEX experiment \\cite{IGEX}) \\be \\label{eq:current} \\meff \\le 0.35 \\, \\zeta~{\\rm eV}~, \\ee where $\\zeta={\\cal O}(1)$ indicates that there is an uncertainty stemming from the nuclear physics involved in calculating the decay width of \\obb.  The running projects NEMO3 \\cite{NEMO} and CUORICINO \\cite{CUORICINO} will be joined in the near future by next generation experiments such as CUORE \\cite{CUORE}, MAJORANA \\cite{MAJORANA}, GERDA \\cite{GERDA}, EXO \\cite{EXO}, MOON \\cite{MOON}, COBRA \\cite{COBRA}, XMASS, DCBA \\cite{DCBA}, CANDLES \\cite{CANDLES}, CAMEO \\cite{CAMEO} (for a review see \\cite{rev_ex}). One can safely expect that values of \\meff{} one order of magnitude below the limit from Eq.\\ (\\ref{eq:current}) will be probed within the next, say, 10 years\\footnote{Not to forget, those experiments aim also to put the controversial \\cite{contr} evidence of part of the Heidelberg-Moscow collaboration to the test.}. This means that scales of order $\\sqrt{\\lsnd}$ will be fully probed, and are even under investigation now. Since the effective mass measured in \\obb\\ also depends on the neutrino mixing angles, the neutrino mass scale and ordering, as well as the mass squared differences, it is possible to obtain additional constraints on sterile neutrino scenarios using neutrinoless double beta decay \\cite{4others,4vbb,4others1}.   Important constraints on sterile neutrinos can also come from cosmology. Inclusion of an extra neutrino, even if sterile, can be in conflict with cosmological observations. The problems are increased if the extra sterile neutrino is massive and has significant mixing with the active species. In particular the Big Bang Nucleosynthesis model of standard cosmology, which explains light element abundances of the Universe, puts constraints on the number of neutrino species. The latest bound found in \\cite{bbn1} for instance is $1.7 < N_{\\nu} < 3.0$ at 95\\% C.L.\\ and in \\cite{bbn2} it is quoted that $N_{\\nu}=3.14^{+0.70}_{-0.65}$. The differences in the results are due to different inputs regarding the uncertainties in the primordial He abundance. Observations of the Cosmic Microwave Background and of large scale structures can also constrain the number of neutrino species. A summary of these bounds obtained by various groups including different data sets can be found in \\cite{han05}. The upper limit on the number of neutrinos in these analysis can vary from 6 to 8. A recent bound as quoted in \\cite{han05} is $N_{\\nu} = 4.2 ^{+1.7}_{-1.2}$ at 95\\% C.L. Another important constraint from cosmology comes on the sum of total masses of all the neutrinos, $\\Sigma \\equiv \\sum m_i$. For four light neutrinos with degenerate masses the bound is $\\Sigma < 1.7$ eV (95\\% C.L.) from WMAP and 2dF data \\cite{han05}. For four (five) neutrinos, with one (two) of them carrying a mass, the bound is $\\Sigma < 1.05~(1.64)$ eV (95\\% C.L.) \\cite{hr0}. Improvement of these numbers within one order of magnitude is expected \\cite{han05}. Note that these bounds depend on the priors and data sets used, for slightly more stringent bounds see, e.g., \\cite{Dodelson:2005tp}. The above constraints can however be evaded if the abundances of sterile neutrinos in the early Universe can be suppressed. This requires going beyond the framework of standard cosmology and introducing mechanisms such as primordial lepton asymmetries \\cite{foot}, low re-heating temperature \\cite{gelmini}, additional neutrino interactions \\cite{bell}\\footnote{See however \\cite{hr}.} etc.  Turning back to oscillations, the MiniBooNE experiment \\cite{miniboone} is expected to confirm or refute the LSND signal and is expected to publish results within the next 6 months or so. If MiniBooNE does not confirm the LSND signal, then with the data collected with $10^{21}$ protons on target they can rule out the entire 90$\\%$ area allowed by LSND with 4 to 5$\\sigma$ \\cite{miniboone}. If however they confirm the LSND signal then this will give rise to an intriguing situation in what regards the explanation of global oscillation data from accelerator, reactor, atmospheric and solar neutrino experiments. If confirming the LSND result, MiniBooNE can not distinguish the allowed four (or five) neutrino mass spectra. To understand the implied mass and mixing scheme, other observables are therefore crucial. This concerns in particular observables depending on the neutrino masses and ordering. Inasmuch one can use these future measurements to identify the neutrino spectrum is one of the motivations of this work. We stress here that we assume only the neutrino oscillation explanation of the LSND result is correct, i.e., the new physics alternatives (not necessarily predicting a signal for MiniBooNE) put forward in Refs.~\\cite{cpt,decoh,mv,4decay,babu,sergio,extra} are not required. \\\\ %valid. \\\\    In this paper we examine what constraints from current and future data can be obtained on possible neutrino mass spectra in scenarios with one or more sterile neutrinos. For the four allowed 3+1 scenarios we give the neutrino masses, their sum as testable in cosmology, the kinematic neutrino mass for tritium experiments and the effective mass in \\onbb$\\!\\!$. We include the most recent values of mass-squared differences and mixing angles from latest global analyzes of oscillation data. We furthermore reconstruct the possible mass matrices in four neutrino scenarios that are consistent with the current data. Finally, we also comment on 3+2 scenarios.  The paper is build up as follows: In Section \\ref{4nu} we discuss our parametrization of the four neutrino mixing matrix and summarize the relevant formulae for the neutrino masses, their sum, the kinematic neutrino mass measured in beta decay experiments and the effective mass that can be observed in neutrinoless double beta decay. In Section \\ref{sec:0vbb3+1} we apply this framework to 3+1 scenarios. Approximate forms of the neutrino mass matrices in 3+1 schemes that are consistent with the current data are given in Section \\ref{sec:matrices}. In Section \\ref{sec:3+2} we comment on the above quantities in the 3+2 scheme, before presenting our summary and conclusions in Section \\ref{sec:concl}. The oscillation probabilities for the relevant short baseline oscillation experiments are delegated to the Appendix. Although 2+2 scenarios are highly disfavored we also add for the sake of completeness an Appendix on the implications of such scenarios for neutrino masses from cosmology, beta decay and neutrinoless double beta decay. We also discuss the form of mass matrices in the 2+2 scenarios.  %\\newpage %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0512/astro-ph0512459_arXiv.txt": {
        "abstract": "{Having conducted a search for the $\\lambda$$\\sim$1.3\\,cm (22\\,GHz) water vapor line towards galaxies with nuclear activity, large nuclear column densities or high infrared luminosities, we present H$_2$O spectra for NGC\\,2273, UGC\\,5101, and NGC\\,3393 with isotropic luminosities of 7, 1500, and 400\\,L$_{\\odot}$. The H$_2$O maser in UGC\\,5101 is by far the most luminous yet found in an ultraluminous infrared galaxy. NGC\\,3393 reveals the classic spectrum of a `disk maser', represented by three distinct groups of Doppler components. As in all other known cases except NGC\\,4258, the rotation velocity of the putative masing disk is well below 1000\\,km\\,s$^{-1}$. Based on the literature and archive data, X-ray absorbing column densities are compiled for the 64 galaxies with reported maser sources beyond the Magellanic Clouds. For NGC\\,2782 and NGC\\,5728, we present {\\it Chandra} archive data that indicate the presence of an active galactic nucleus in both galaxies. Modeling the hard nuclear X-ray emission, NGC\\,2782 is best fit by a high energy reflection spectrum with $N_{\\rm H}$$\\ga$10$^{24}$\\,cm$^{-2}$. For NGC\\,5728, partial absorption with a power law spectrum indicates $N_{\\rm H}$$\\sim$8$\\times$10$^{23}$\\,cm$^{-2}$. The correlation between absorbing column and H$_2$O emission is analyzed. There is a striking difference between kilo- and megamasers with megamasers being associated with higher column densities. All kilomasers ($L_{\\rm H_2O}$$<$10\\,L$_{\\odot}$) except NGC\\,2273 and NGC\\,5194 are Compton-thin, i.e. their absorbing columns are $<$10$^{24}$cm$^{-2}$. Among the H$_{2}$O megamasers, 50\\% arise from Compton-thick and 85\\% from heavily obscured ($>$10$^{23}$cm$^{-2}$) active galactic nuclei. These values are not larger but consistent with those from samples of Seyfert 2 galaxies not selected on the basis of maser emission. The similarity in column densities can be explained by small deviations in position between maser spots and nuclear X-ray source and a high degree of clumpiness in the circumnuclear interstellar medium. ",
        "introduction": "The $J_{\\rm K_aK_c} = 6_{16}-5_{23}$ line of ortho-H$_2$O at 22\\,GHz ($\\lambda$$\\sim$1.3\\,cm) is one of the strongest and most remarkable spectral features of the electromagnetic spectrum. The transition, connecting levels located approximately 645\\,K above the ground state, traces warm ($T_{\\rm kin}$$\\ga$400\\,K) and dense ($n$(H$_2)$$\\ga$10$^{7}$\\,cm$^{-3}$) molecular gas (e.g. Kylafis \\& Norman 1987, 1991). Observed as a maser, the line can reach isotropic luminosities in excess of 10$^4$\\,L$_{\\odot}$ (Barvainis \\& Antonucci 2005). Being traditionally known to trace oxygen-rich red giant stars and to pinpoint sites of (massive) star formation, it has more recently become a major tool to determine geometric distances and three dimensional motions of nearby galaxies and to elucidate the nuclear environment of active galaxies, allowing us to map accretion disks and to determine masses of nuclear engines (for recent reviews, see Greenhill 2002, 2004; Maloney 2002; Henkel \\& Braatz 2003; Morganti et al. 2004; Henkel et al. 2005a; Lo 2005; for 3-D motions, see Brunthaler et al. 2005).  Extragalactic H$_2$O masers have been observed for a quarter of a century (Churchwell et al. 1977) and emission has been reported from more than 60 galaxies (e.g., Henkel et al. 2005b, Kondratko et al. 2006). The masers can be classified as related (1) to star formation, (2) to nuclear accretion disks (`disk-masers'), (3) to interactions of nuclear jet(s) with ambient molecular clouds or to amplification of the jet's seed photons by suitably located foreground clouds (`jet-masers') and (4) to nuclear outflows. H$_2$O masers with apparent (isotropic) luminosities $L_{\\rm H_2O}$ $<$ 10\\,L$_{\\odot}$, often associated with sites of massive star formation, are referred to as `kilomasers', while the stronger sources, known to arise from the innermost parsecs of their parent galaxy, are referred to as `megamasers'. The statistical properties of the sample of extragalactic H$_2$O masers are not well explored. The first and last such comprehensive study, that of Braatz et al. (1997), dates back to when little more than a dozen sources were known.  In order to improve our knowledge of the statistical properties of the greatly enlarged maser sample, we analyze the correlation between maser emission and X-ray absorbing column density toward the nuclear engine. In Sects.\\,2--3 we summarize recent observations and show spectra of sources detected in this work. In Sect.\\,4, we present X-ray column densities for the entire extragalactic maser sample, including a detailed analysis of {\\it Chandra} archive data of the megamaser galaxies NGC\\,2782 and NGC\\,5728. Specific statistical properties of the sample involving X-ray absorbing columns and maser luminosities are discussed in Sect.\\,5. Sect.\\,6 summarizes the main results. ",
        "conclusions": "In this paper, maser emission and column density are discussed for 64 extragalactic H$_2$O maser sources. The main results are:  (1) In searching for the $\\lambda$$\\sim$1.3\\,cm (22\\,GHz) water vapor line in galaxies with high nuclear column densities, active nuclei or high infrared luminosities, we detected H$_2$O emission towards NGC\\,2273, UGC\\,5101, and NGC\\,3393 with isotropic luminosities of 7, 1500, and 400\\,L$_{\\odot}$, respectively. UGC\\,5101, one of the most luminous H$_2$O megamasers ever found, is the first luminous ($L_{\\rm H_2O}$$>$100\\,L$_{\\odot}$) H$_2$O maser detected in an ULIRG. The maser is likely nuclear. NGC\\,2273 hosts a maser that is either nuclear or associated with massive star formation. Since its lineshape is wide, suggesting a long lifetime, interferometric measurements may soon reveal its nature. NGC\\,3393, first detected by Kondratko et al. (2006), shows a classical NGC\\,4258-like profile. If the analogy holds, the rotation velocity of the disk is well below the 1000\\,km\\,s$^{-1}$ encountered in NGC\\,4258. This and similarly low rotation velocities in other `disk-maser' galaxies suggest that the accretion disk in NGC\\,4258, surrounding a nuclear engine of relatively high mass, is particularly compact and well ordered.  (2) {\\it Chandra} X-ray images and spectra are presented and analyzed for NGC\\,2782 and NGC\\,5728. A hidden AGN was detected in NGC\\,2782. The X-ray spectrum is best fit by a high energy reflection model with $N_{\\rm H}$$\\ga$10$^{24}$cm$^{-2}$; its 6.4\\,keV Fe line appears to be buried in a spatially more extended, relatively bright continuum. The H$_2$O maser, so far believed to be associated with star formation, may be related to the nuclear engine. The X-ray spatial analysis of NGC\\,5728 confirms the biconical axisymmetric structure of its ionization region, which was previously revealed by optical HST emission line studies. The best fitting model, with a partially absorbed power law spectrum, yields $N_{\\rm H}$$\\sim$8$\\times$10$^{23}$\\,cm$^{-2}$.  (3) A statistical analysis of the correlation between X-ray absorbing column density and H$_2$O emission was carried out, encompassing 64 galaxies beyond the Magellanic Clouds. At a confidence level of at least 99.4\\%, the column density distributions of kilomaser and megamaser sources are different. More than 80\\% (22/26) of the megamaser sources are heavily obscured ($N_{\\rm H}$$>$10$^{23}$cm$^{-2}$) and half are Compton-thick ($N_{\\rm H}$ $\\geq$10$^{24}$cm$^{-2}$). Among the 13 known kilomaser sources ($L_{\\rm H_2O}$ $<$ 10\\,L$_{\\odot}$), however, only NGC\\,2273 and NGC\\,5194 are Compton-thick. The small subsamples of nuclear kilomasers and disk megamasers show average column densities that are undistinguishable from those of the entire megamaser sample and from samples of Seyfert 2 galaxies that were not selected with respect to maser emission. Presumably a clumpy cloud structure in the cirumnuclear environment, diverging positions between maser and nuclear sources and an occasional amplification of a background radio continuum source are sufficiently decoupling X-ray column densities and H$_2$O maser properties not to show a clear correlation.  (4) The identification of an AGN of type 2 is sufficient to justify a search for H$_2$O maser emission. The determination of X-ray column densities toward the nuclear engine does not enhance maser detection probabilities."
    },
    "0512/astro-ph0512173_arXiv.txt": {
        "abstract": "In this paper, we give the spectral energy distributions of 42 M81 globular clusters in 13 intermediate-band filters from 4000 to 10000{\\AA}, using the CCD images of M81 observed as part of the BATC multicolor survey of the Sky. The BATC multicolor filter system is specifically designed to exclude most of the bright and variable night-sky emission lines including the OH forest. Hence, it can present accurate SEDs of the observed objects. These spectral energy distributions are low-resolution spectra, and can reflect the stellar populations of the globular clusters. This paper confirms the conclusions of \\citet{sbkhp02} that, M81 contains clusters as young as a few Gyrs, which also were observed in both M31 and M33. ",
        "introduction": "The study of globular clusters (GCs) plays an important role in our understanding of the evolution and history of galaxies. They are bright and easily identifiable star clusters typically with homogeneous abundances and ages. The Galactic GCs, the stars of which are thought to be among the oldest stars in the universe provide important information regarding the minimum age of the universe and the early formation history of our Galaxy. However, we also find that the GC system of our neighboring galaxy, M31, contains at least 20 young GCs ranging in age from 100 Myr to $\\sim 5$ Gyr \\citep{bur04,b05}.  Except for the Local Group galaxies, M81 is one of the nearest large spirals outside the Local Group. As such, its globular cluster system has come under recent detailed scrutiny. \\citet{pr95} first attempted to identify GCs in M81 from ground-based images, sifting through over 3700 objects in a $50\\arcmin$ diameter field centered on M81. They found $~70$ GC candidates within 11 kpc galactocentric radius. \\citet{pbh95} then confirmed 25 as M81 GCs on basis of spectroscopy of 82 bright GC candidates in the M81 field. \\citet{sbkhp02} obtained moderate-resolution spectroscopy for 16 of the Perelmuter \\& Racine M81 GC candidates, and found that all of these are GCs.  Recently, \\citet{cft01} discovered 114 compact star clusters in M81 from $B$-, $V$-, and $I$-band {\\it {Hubble Space Telescope}} (HST) Wide Field Planetary Camera 2 images in eight fields, covering a total area of 40 arcmin, 54 of which are new GCs. Using these 95 M81 GCs, \\citet{ma05} presented that the intrinsic $B$ and $V$ colors and metallicities of these GCs are bimodal, with metallicity peaks at $\\rm {[Fe/H]}\\approx -1.45$ and $-0.53$, similar to what we find for the Milky Way and M31 GCs. In this paper we present new spectral energy distributions (SEDs) for 42 of these GCs, using M81 images observed as part by galaxy calibration program of the Beijing-Arizona-Taiwan-Connecticut (BATC) multicolor sky survey \\citep[e.g.,][]{fan96,zheng99}. The BATC filters are custom-designed set of 15 intermediate-band filters to do spectrophotometry for preselected 1 deg$^{2}$ regions of the northern sky.  Details of our observations and data reduction are given in \\S~2. \\S~3 gives our summary.  \\section {Observations and Data Reduction}  \\subsection {The Sample of GCs}  \\citet{ma05} studied the intrinsic $B$ and $V$ colors and metallicities of 95 M81 GCs. In order to study the stellar populations of these GCs, we extracted 311 images of M81 field as part of the BATC multicolor survey of the sky, taken in 13 intermediate-band filters with a total exposure time of $\\sim 100$ hours from February 5, 1995 to April 30, 2002. Multiple images of the same filter were combined to improve the signal-to-noise ratio. While the SEDs of the sample GCs brighter than $V\\sim 20$ mag can be obtained in the BATC multicolor system, we are constrained in obtaining full SEDs for these GCs by the limited field of view around M81, as well as by some of these GCs being in the high background of M81 itself. As such, we have obtained SEDs in 13 BATC filters for 42 of the 95 GCs previously presented.  \\subsection{Observations and data reduction}  The BATC multicolor survey uses a Ford Aerospace $2048\\times 2048$ CCD camera with 15 $\\mu$m pixel size on the 0.6/0.9m f/3 Schmidt telescope of the Xinglong Station of the National Astronomical Observatories, giving a CCD field of view of $58^{\\prime}$ $\\times $ $58^{\\prime}$ with a pixel size of $1\\arcsec{\\mbox{}\\hspace{-0.15cm}.} 7$. The typical seeing of the Xinglong station is $2\\arcsec$. The BATC multicolor filter system, which was specifically designed to avoid contamination from the brightest and most variable night sky emission lines, includes 15 intermediate-band filters covering 3300{\\AA} to 1$\\mu$. Calibrations of these images are made using observations of four $F$ sub-dwarfs, HD~19445, HD~84937, BD~${+26^{\\circ}2606}$, and BD~${+17^{\\circ}4708}$, all taken from \\citet{ok83}.  Hence, our magnitudes are defined in a way similar to the spectrophotometric AB magnitude system that is the Oke \\& Gunn $\\tilde{f_{\\nu}}$ monochromatic system. BATC magnitudes are defined on the AB magnitude system as  \\begin{equation} m_{\\rm batc}=-2.5{\\rm log}\\tilde{F_{\\nu}}-48.60, \\end{equation}  \\noindent where $\\tilde{F_{\\nu}}$ is the appropriately averaged monochromatic flux in unit of erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ at the effective wavelength of the specific passband. In the BATC system \\citep{yan00}, $\\tilde{F_{\\nu}}$ is defined as  \\begin{equation} \\tilde{F_{\\nu}}=\\frac{\\int{d} ({\\rm log}\\nu)f_{\\nu}r_{\\nu}} {\\int{d} ({\\rm log}\\nu)r_{\\nu}}, \\end{equation}  \\noindent which links the magnitude to the number of photons detected by the CCD rather than to the input flux \\citep{fuku96}. In equation (2), $r_{\\nu}$ is the system's response, $f_{\\nu}$ is the SEDs of the source.  Of the 15 BATC filters, we did not use the two bluest filters. Data reduction of the CCD data proceeds with removal of bias subtraction and flat-fielding with dome flats. These steps were performed with our custom-made, automatic data reduction software, PIPELINE I, developed for the BATC multicolor sky survey \\citep{fan96,zheng99}. The dome flat-field images were taken by using a diffuser plate in front of the correcting plate of the Schmidt telescope, a flatfielding technique which has been verified with the photometry we have done on other galaxies and fields of view \\citep[e.g.,][]{fan96,zheng99,wu02,yan00,zhou01,zhou04}. Spectrophotometric calibration of the M81 images using the Oke-Gunn standard stars is done during photometric nights \\citep[see details from][]{yan00, zhou01}.  Using the images of the standard stars observed on photometric nights, we derived iteratively the atmospheric extinction curves and the variation of these extinction coefficients with time \\citep[cf.][]{yan00,zhou01}. The extinction coefficients at any given time in a night $[K+ \\Delta K (UT)]$ and the zero points of the instrumental magnitudes ($C$) are obtained by  \\begin{equation} m_{\\rm batc}=m_{\\rm inst}+[K+\\Delta K(UT)]X+C, \\end{equation}  \\noindent where $X$ is the air mass. The instrumental magnitudes ($m_{\\rm inst}$) of the selected bright, isolated and unsaturated stars on the M81 field images of the same photometric nights can be readily transformed to the BATC AB magnitude system ($m_{\\rm batc}$). The calibrated magnitudes of these stars are obtained on the photometric nights, which are then used as secondary standards to uniformly combine images from calibrated nights to those taken during non-photometric weather. Table~1 lists the parameters of the BATC multicolor filter system and the statistics of observations. Column 6 of Table~1 gives the scatter, in magnitudes, for the photometric observations of the four primary standard stars in each filter.  \\subsection{Integrated photometry}  For each M81 GC, the PHOT routine in DAOPHOT \\citep{stet87} is used to obtain magnitudes. To avoid contamination from nearby objects, we adopt a small aperture of $6\\arcsec{\\mbox{}\\hspace{-0.15cm}.}8$ corresponding to a diameter of 4 pixels in the Ford CCD. Aperture corrections are determined as follows, using the isolated M81 GC, Is40165: determining the magnitude differences between photometric diameters 4 and 10 pixels in each of the 13 BATC filters. Inner and outer radii of the sky apertures are from 4 to 7 pixels for a diameter of 4 pixels, and from 6 to 10 pixels for a diameter of 10 pixels. The SEDs obtained in this manner are given in Table~2. Column 1 is GC's name taken from \\citet{pbh95}, \\citet{sbkhp02} and \\citet{cft01}. Column 2 to column 14 give the magnitudes of the 13 BATC passbands observed. The second line for each GC gives the $1-\\sigma$ errors in magnitudes for the corresponding passband. The errors for each filter are given by DAOPHOT. Magnitudes in the BATC03 filter could not be obtained for Is50286, Id50357, Id70319, SBKHP16 and  CFT41 owing to low signal-to-noise ratios. Because of low angular resolution, given the Schmidt pixel size of $1\\arcsec{\\mbox{}\\hspace{-0.15cm}.} 7$, the different sizes of different clusters are not evident in our CCD images.  \\subsection{Comparison with previous photometry}  \\citet{zhou03} presented the relationships between the BATC intermediate-band system and $UBVRI$ broadband system using the standard stars catalogs of \\citet{land83,land92} and \\citet{gtj00}. The coefficients of two relationships are showed by equations (4) and (5) below:  \\begin{equation} m_B=m_{04}+0.2201(m_{03}-m_{05})+0.1278\\pm0.076, \\end{equation}  \\begin{equation} m_V=m_{07}+0.3292(m_{06}-m_{08})+0.0476\\pm0.027. \\end{equation}  Using equations (4) and (5), we transformed the magnitudes of 42 GCs in BATC03, BATC04 and BATC05 bands to ones in the $B$ band, and in BATC06, BATC07 and BATC08 bands to ones in $V$ band. Figure 1 plots the comparison of $V$ (BATC) and ($B$$-$$V$) (BATC) photometry with previously published measurements of \\citet{pr95} and \\citet{cft01}. In this figure, our  magnitudes/colors are on the x-axis, the difference between our and \\citet{pr95} and \\citet{cft01} magnitudes/colors are on the y-axis. Table 3 lists this comparison. The mean $V$ magnitude and color differences (in the sense of this paper $-$ \\citet{pr95} and \\citet{cft01}) are $<\\Delta V> =-0.116\\pm0.028$ and $<\\Delta (B-V)>=-0.017\\pm 0.027$, respectively. The uncertainties in $<\\Delta V>$ and $<\\Delta (B-V)>$ are calculated by  \\begin{equation} \\sqrt{\\frac{\\Sigma(<\\Delta V>-\\overline{<\\Delta V>})^2}{N(N-1)}}, \\end{equation} and \\begin{equation} \\sqrt{\\frac{\\Sigma(<\\Delta (B-V)>-\\overline{<\\Delta (B-V)>})^2}{N(N-1)}}. \\end{equation}  Uncertainties in $B$ (BATC) and $V$ (BATC) have been added linearly, i.e. $\\sigma_B=\\sigma_{04}+0.2201(\\sigma_{03}+\\sigma_{05})$, and $\\sigma_V=\\sigma_{07}+0.3292(\\sigma_{06}+\\sigma_{08})$, to reflect the errors in the three filter measurements. For the colors, we add the errors in quadrature, i.e. $\\sigma_{(B-V)}={(\\sigma_B2+\\sigma_V2)}^{1/2}$. From Figure 1 and Table 3, it can be seen that there is good agreement in the photometric measurements.  \\begin{figure*} \\begin{center} \\centerline{\\includegraphics[angle=-90,width=120mm]{f1.eps}} \\caption{Comparison of cluster photometry with previous measurements by \\citet{pr95}, shown as open circles, and \\citet{cft01}, shown as crosses.} \\label{fig1} \\end{center} \\end{figure*}  \\subsection{Reddening}  In order to obtain intrinsic SEDs for the sample GCs, the photometric data are corrected for the reddening from the foreground extinction contribution of the Milky Way and for the internal reddening due to varying optical paths through the disk of M81. The total reddening determination for the M81 field (the foreground plus M81 contribution) has been measured by a number of authors \\citep[e.g.,][]{fwm94,kong00}. We only mention here that, \\citet{kong00} obtained the reddening maps of M81 field based on the images observed by the BATC multicolor sky survey in the 13 intermediate-band filters from 3800 to 10000 \\AA. To determine the metallicity, age, and reddening distributions for the M81 field, \\citet{kong00} found the best match between the observed colors and the predictions from single stellar population models of \\citet{bc96}. A map of the interstellar reddening in a substantial portion of M81 was obtained. For a few clusters that fall near the edges of the images, \\citet{kong00} did not obtain reddenings. For these clusters we adopt a mean reddening value of 0.13 as \\citet{cft01} did.  The local reddening values for these GCs are listed in column (4) of Table 1 in \\citet{ma05}. Figure 2 plots the intrinsic SEDs of 42 GCs (relative to the flux of filter BATC08) in the 13 BATC intermediate-band filters.  \\begin{figure*} \\begin{center} \\centerline{\\includegraphics[angle=-90,width=170mm]{f2.eps}} \\caption{Intrinsic spectrophotometric energy distributions for 42 GCs in M81.} \\label{fig2} \\end{center} \\end{figure*}   \\subsection{Analysis of the SEDs of the GCs}  \\citet{sbkhp02} observed moderate-resolution spectroscopy of 16 M81 GCs using the Low Rosolution Imaging Spectrograph on the Keck I telescope. By comparing between the observed age-sensitive index $\\rm H\\beta$ against Mg2 and isochrones from the evolutionary synthesis models from \\citet{worthey94} and from \\citet{fab95}, \\citet{sbkhp02} find that SBKHP15\\footnote{In fact, SBKHP15 should be SBKHP16. We referred the Table 6 of \\citet{pr95} and found that the R.A. (J2000) and Decl. (J2000) of ID 50867 are 09:55:40.194 and 69:07:30.82, and the R.A. (J2000) and Decl. (J2000) of ID 50889 are 09:55:51.995 and 69:07:39.32, i.e. the R.A. (J2000) and Decl. (J2000) of SBKHP15 and 16 should be exchanged in \\citet{sbkhp02}.} is younger than the other GCs. As we know that, lower metallicity and younger age can make the fluxes in longer filter bands to be lower. The metallicity of this cluster is nearly the same as one of SBKHP8. So, from the SEDs of Figure 2, we can conclude that SBKHP16 is younger than SBKHP8, as its SED is lower than those of SBKHP8 in longer filter bands. In particular, SBKHP16 has very low fluxes in the BATC14 and BATC15 filter bands. We compare its SEDs with ones of SBKHP8, and find that, the intrinsic flux (relative to the flux of BATC08 filter band) is 1.607 versus 0.641 in BATC14 filter band, and 1.505 versus 0.649 in BATC15 filter band, nearly 2.5 times.  \\subsection{Ages}  A single GC is a stellar population having a single age and chemical abundance. Globular clusters are ideal systems to be characterized by simple stellar populations (SSPs) models.  BC96 models \\citep{bc96} are given for simple stellar populations (SSPs) of metallicities $Z=0.0004, 0.004, 0.008, 0.02, 0.05$, and $0.1$. These models are based on the Padova group evolutionary tracks \\citep{bres93,fbbc94,gbcbn96}, which use the radiative opacities of \\citet{irw92} together with a helium abundance $Y=2.5Z+0.23$ (The reference solar metallicity is $Z_\\odot=0.02$). BC96 models further use the \\citet{lcb97} standard star library. The ages in the BC96 models range from 0 to 20 Gyr. A \\citet{salp55} IMF of $\\Phi(M)=A \\times M^{-\\alpha}$ with $\\alpha=2.35$ is used with a normalization constant $A=1$, a lower cutoff mass $M_{\\rm l}=0.1M_{\\odot}$ and an upper cutoff mass $M_{\\rm u}=125M_{\\odot}$.  To proceed with the comparisons, we first convolve the SEDs of BC96 models with the BATC filter profiles to obtain the optical and near-infrared integrated luminosities. The integrated luminosities $L_{\\lambda_i}(t,Z)$ of the $i$th BATC filter can be calculated as  \\begin{equation} L_{\\lambda_i}(t,Z) =\\frac{\\int F_{\\lambda}(t,Z)\\varphi_i(\\lambda)d\\lambda} {\\int \\varphi_i(\\lambda)d\\lambda}, \\end{equation}  \\noindent where $F_{\\lambda}(t,Z)$ is the SED at age $t$ in metallicity $Z$ model, $\\varphi_i(\\lambda)$ is the response functions of the $i$th filter of the BATC filter system ($i=3, 4, \\cdot\\cdot\\cdot, 15$), respectively. All integrated colors of BC96 models are calculated relative to the BATC08 filter band ($\\lambda=6075${\\AA}):  \\begin{equation} \\label{color} C_{\\lambda_i}(t,Z)={L_{\\lambda_i}(t,Z)}/{L_{6075}(t,Z)}. \\end{equation}  \\noindent From this equation, we can obtain model intermediate-band colors for SSP models of different metallicities.  In order to study the stellar populations of these GCs, we use 5 ages of 1, 2, 3, 8 and 16 Gyrs of BC96 SSP models of two metallicities: a metal-poor model of $Z=0.0004$, and a more metal-rich model of $Z=0.004$.  The best SED-fit between a globular cluster and SSP models is found by minimizing the color differences between intrinsic integrated color of a cluster and integrated color of models:  \\begin{equation} R^2(n,t,Z)=\\frac{\\sum_{i=3}^{15}{[C_{\\lambda_i}^{\\rm intr}(n)-C_{\\lambda_i}^ {\\rm ssp}(t, Z)]^2}/\\sigma_{i}^{2}} {{\\sum_{i=3}^{15}}{1/\\sigma_{i}^{2}}}, \\end{equation}  \\noindent where $C_{\\lambda_i}^{\\rm ssp}(t, Z)$ represents the integrated color in the $i$th filter of a SSP at age $t$ in a metallicity $Z$ model. $C_{\\lambda_i}^{\\rm intr}(n)$ is the intrinsic integrated color for a cluster. The differences are weighted by $1/{\\sigma_i}^{2}$, where the ${\\sigma_i}$'s are observational uncertainties of the passbands. Figure 3 shows the results of SED-fits, in which filled circle represents the intrinsic integrated color of a cluster, and the thick line represents the best fit of the integrated color of a SSP model.  \\begin{figure*} \\begin{center} \\centerline{\\includegraphics[angle=-90,width=170mm]{f3.eps}} \\caption{Map of the fit of the integrated color of a SSP model with intrinsic integrated color for sample GCs. Filled circle represents the intrinsic integrated color of a GC, and the thick line represents the best fit of the integrated color of a SSP model.} \\label{fig3} \\end{center} \\end{figure*}  From Figure 3, we can see that, of these 42 M81 GCs, there are 11 for which our estimates give ages younger than 8 Gyrs. The results tell us that, M81 includes a population of intermediate-age GCs with ages of a few Gyr. Similar clusters have been observed in both M31 and M33 \\citep{brodie90,brodie91,jiang03,bb04,bur04,ppb05, sarajedini98,sarajedini00,ma02}. For SBKHP16, \\citet{sbkhp02} derived its age is between 1.5 and 3 Gyrs, and our result is consistent with this estimte, given an age of between 1 and 2 Gyrs. Our results also shows that the age of SBKHP13 is between 2 and 3 Gyrs, which was not presented by \\citet{sbkhp02}, because the value of index $\\rm H\\beta$ was not derived in \\citet{sbkhp02}. The ages of the other GCs of \\citet{sbkhp02} obtained in this paper are also fully consistent with the results of \\citet{sbkhp02}. ",
        "conclusions": "We have obtained SEDs of 42 M81 GCs in 13 intermediate-band filters with the BATC 0.6/0.9m Schmidt telescope. The BATC filter system is specifically designed to exclude most of the bright and variable night-sky emission lines including the OH forest, and it can present the accurate SEDs of the observed objects. This paper confirms the conclusions of \\citet{sbkhp02} that, M81 contains clusters as young as a few Gyrs. Such young GCs have also been observed in both M31 and M33 \\citep{brodie90,brodie91,jiang03,bb04,bur04,ppb05, sarajedini98,sarajedini00,ma02}."
    },
    "0512/astro-ph0512490_arXiv.txt": {
        "abstract": "We present a series of numerical studies of the interaction of colliding radiative, hydrodynamic young stellar outflows. We study the effect of the collision impact parameter on the acceleration of ambient material and the degree to which the flow is isotropized by the collision as a mechanism for driving turbulence in the parent molecular cloud. Our results indicate that the high degrees of compression of outflow material, achieved through radiative shocks near the vertex of the interaction, prevents the redirected outflow from spraying over a large spatial region.  Furthermore, the collision reduces the redirected outflow's ability to entrain and impart momentum into the ambient cloud.  Consideration of the probabilities of outflow collisions leads us to conclude that individual low velocity fossil outflows are the principle coupling between outflows and the cloud. ",
        "introduction": "\\label{s1} Molecular Clouds have long been a subject of interest in astrophysics since they are the exclusive environments in which stars form in galaxies.  The expected lifetimes for molecular clouds has become a topic of considerable debate as numerical simulations have shown that MHD turbulence, the nominal means of support for the clouds against self-gravity, will decay on a crossing timescale \\citep{maclow,stone,vazquez-semadeni}.  In light of this result it is difficult to understand why molecular clouds do not fully collapse in an efficient burst of star formation on timescales no longer than a few crossing times. Thus it appears that either molecular clouds are transient features or they are resupplied with turbulent energy through some other means.  Jets and molecular outflows are recognized as a ubiquitous phenomena associated with star formation.  It is expected that most if not all low mass stars produce a collimated outflow during their formation from a parent molecular cloud core (massive stars may also produce collimated outflows though this point remains somewhat speculative \\citep{Shepherd}).  Stars do not, however, form in isolation.  Rich star forming regions such as Orion can contain as many as 1000 stars per $\\textrm{pc}^3$ \\citep{testi}.  Low mass star forming regions such as Taurus or Perseus will contain hundreds of stars in a similar volume.  The ubiquity and high density of outflows from young stars make them an intriguing candidate for the source of turbulent energy in molecular clouds.  The idea that feedback from TT winds could lead to a self-regulating state of star formation dates back as far as \\cite{Norman Silk}. Consideration of combined energy budget for the outflows in some clouds compared with the energy in the cloud's turbulent motions support notions of feedback showing an approximate balance between outflow input and turbulent support \\cite{Bally,Bally Reip,Knee,Matzner 2002,Warin}. Thus the combined action of many outflows could, in principle, provide the required deposition of turbulent energy to support a cloud against collapse. More recent observational studies have explored multiple (though apparently) non-interacting outflow structures in individual clouds and come to similar conclusions. For example direct observational evidence showing that the giant stellar outflows associated with HH 300 and HH 315 have disrupted their cloud's density and velocity distributions at parsec scale distances from their source has been provided by \\citet{Arce 2003}. The actual global disruptive effect these individual flows have depends on the ability that these outflows have to impart their momentum into their parent clouds by entraining and accelerating molecular gas \\citep{Arce 2003, Arce 2001}.  While invoking jets and outflows to drive turbulent motions appears attractive for molecular cloud studies, there is a potential problem with such a scenario. The principle means of energy transfer from jet to cloud appears to come via shock waves, the so-called ``prompt entrainment'' mechanism \\cite{Chernin}. This is to be compared with ``turbulent entrainment'' mechanism which occurs via a turbulent boundary at the edge of a jet \\cite{Canto}.  Thus the effect of a single supersonic outflow is bounded by the shock wave which defines it.  Only those regions of a cloud which have been swept over by the outflow will gain any energy. Given such a localization of energy and momenta deposition, the action of multiple, isotropically oriented outflows is required to drive the random motions associated with isotropic turbulence.  Somehow the energy and momenta in the localized region engulfed by a jet or outflow must be randomized and distributed over many scales.  The results presented here consider the facility of the collision of two protostellar outflows toward this result.  We consider the interaction of two heavy jet flows oriented at $90\\,^{\\circ}$ to each other with different impact parameters.  Our goals in this study are to examine the resultant flows and attempt to distinguish between those with low and high impact parameters in terms of how ambient gas is accelerated. ",
        "conclusions": "The radiative energy losses when unshocked protostellar outflow streams directly collide reduces the kinetic energy available to deposit into the molecular cloud as turbulent energy.  The high degree of compression of outflow gas induced by cooling from such a collision prevents the redirected outflow from spraying over a large spatial region.  Furthermore, the collision reduces the redirected outflow's ability to entrain and impart momentum into the ambient cloud.  The cooling of the interaction region produces a reduced bow shock surface area over which outflow momentum can be exchanged with ambient gas. ``Grazing'' collisions, where only the cocoon of shock-decelerated gas or a small fraction of the high speed outflow gas collide, however, have little effect on radiative energy loss or the rate of entrainment of ambient material into the flow.  Because the direct collision of protostellar outflows is rare, we conclude that such collisions have little effect on the turbulent energy budget of molecular clouds.  Based on the results of this study we conclude that if turbulence is energized by outflows it does not occur through collisions of active outflows. Instead the mechanical energy of an outflow is most likely supplied to the turbulent motions of the cloud through the action of {\\it fossil} cavities that remain after the driving source of the outflow has expired.  The role of individual or possibly overlapping fossil outflows was explored before \\citep{Quillen et al 2005}. This study focused on NGC 1333 and explored the interaction of the cloud with slowly moving shells which remain after the outflow source has either shut down or become signifigantly weakened by the decrease in $\\dot{M}_j$. The fossil cavities were shown to carry significant momentum and can provide the coupling mechanism between outflow and turbulent motions in the cloud. Using the bow shock radius and outflow length in \\S\\ref{s2}, without accounting for collisons, the volume fill ratio of outflow cavities exceeds unity at stellar density $> 32~\\textrm{pc}^{-3}$. We speculate that as the density of protostars approaches this value, the parent cloud will become subsumed in motion driven by randomly oriented fossil cavities.  Future work should focus on the interaction of fossil outflows that were launched at different times within a turbulent cloud and their overlap to unravel the exact mechanisms by which this provides a route to sustained turbulence."
    },
    "0512/astro-ph0512345_arXiv.txt": {
        "abstract": "We study the phenomenology of cosmic-rays (CRs) at the galactic/extragalactic transition, focusing on two opposite models for the composition of the extragalactic (EG) component. Model A assumes a mixed source composition, with nuclear abundances similar to that of the low-energy CRs, while model B assumes that EG sources accelerate only protons. We study the limits within which both scenarios can reproduce the observed high-energy CR spectrum and composition. The ankle in model A is interpreted as the GCR/EGCR transition, while in model B it is the pair-production dip. Model A has a source spectrum $\\propto E^{-x}$ with $x \\sim 2.2 - 2.3$, while model B requires $x \\sim 2.6 - 2.7$. We compare the predictions of both models with the available data on the energy evolution of the high-energy CR composition using the two main composition-related observables: $X_{\\max}$ and $\\langle\\ln A\\rangle$. We conclude that model~A is currently favoured. Uncertainties are discussed and distinctive features of the two models are identified, which should allow one to distinguish between the models in the near future when more precise measurements are available with higher-statistics experiments. ",
        "introduction": "\\label{sec:introduction}  Despite its remarkable regularity over more than 12 orders of magnitude in energy, the cosmic-ray (CR) spectrum exhibits a few features that are subject to intense observational and theoretical studies, as they may provide valuable information to eventually understand the origin of both galactic (GCRs) and extragalactic (EGCRs) cosmic-rays. Along with the so-called knee around $5\\,10^{15}$~eV, whose precise energy and composition structure is not firmly established yet \\cite{Kascade}, and the expected, but observationally uncertain, suppression at high energies due to CR interactions with the cosmological microwave background (CMB), the energy range between $10^{18}$ and $10^{19}$~eV is of particular interest since it is believed to host the transition from GCRs to EGCRs. Such a transition must indeed occur as it is established (notably from observations of Galactic gamma-rays due to $\\pi^{0}$ CR interactions) that the low-energy CRs have a Galactic origin, while simple considerations about the confinement of particles in the Galaxy and the Galactic halo strongly suggest that most of the highest-energy CRs must have an extragalactic origin (unless their charge is unexpectedly large, which is not favoured by the observations).  The most natural shape that a transition between two steady (featureless) components may take \\emph{a priori} is that of an ``ankle',' i.e., a hardening of the spectrum, where the harder, initially subdominant component at low energy simply takes over from the softer component at \\emph{some} transition energy, $E_{\\mathrm{t}}$. Conversely, a ``knee-like'' transition corresponding to a steepening of the spectrum with no associated discontinuity is \\emph{a priori} unlikely, since it can be smooth only if the second component starts (roughly) at the energy where the first one stops \\emph{and} if at that energy they have (roughly) the same flux. Interestingly, an ankle is indeed observed in the CR spectrum, around $3\\,10^{18}$~eV, precisely in the energy range where the confining effect of the Galactic magnetic fields is expected to lose its efficiency (because the gyroradius of the particles of charge $Z$ in a microGauss field, $r_{\\mathrm{g}} \\simeq 1\\,\\mathrm{kpc}\\ Z^{-1}B_{\\mu\\mathrm{G}}^{-1}$, becomes comparable to the thickness of the Galaxy) and the GCR component is thus expected to die out. It is thus tempting to interpret the observed ankle as a natural signature of the GCR/EGCR transition (e.g., \\cite{Nagano92}).  In the $10^{18}$ to $10^{19}$~eV energy range, another feature in the EGCR spectrum is expected, if one assumes that the EGCR sources only accelerate protons. In this case, the energy losses associated with the production of e$^{+}$e$^{-}$ pairs during the proton transport through the CMB lead to a so-called ``pair production dip'' in the propagated spectrum \\cite{Berezinsky+88}, the shape of which is reminiscent of the CR ankle. From this point of view, the ankle can be interpreted as a signature of a pure proton EGCR component  \\cite{Berezinsky+02,Berezinsky+04}, in which case the GCR/EGCR transition must occur at a lower energy, possibly around the so-called ``second knee'' that may be present in the data but with an uncertain energy scale. Some support for such an interpretation can be found in the recent results of the HiRes experiment \\cite{HiRescomp} reporting a rapid variation of the depth of CR-induced extensive air showers (EAS), interpreted as a change of composition from heavy (typically iron) to light nuclei (protons), between $10^{17}$ and $10^{18}$ eV, i.e., below the ankle. A key assumption in the interpretation of the ankle as a pair production dip is the effective  absence of nuclei heavier than hydrogen among EGCRs, not only by the time they reach the Earth but also at their source \\cite{Berezinsky+04,Allard2005}.  Another important parameter of the EGCR source scenario is the logarithmic slope of the assumed power-law injection spectrum, $Q(E) \\propto E^{-x}$. In a recent study \\cite{Allard2005}, we have drawn attention to the influence of the EGCR source composition on both the interpretation of the ankle and the acceleration scenario. Our results can be summarized as follows: i) if the sources accelerate only protons, then the observed high-energy CR spectrum can be well reproduced with a steep injection spectrum, $x\\simeq 2.6$--2.7 \\cite{Berezinsky+02,De marco+03}, and the ankle can be interpreted as a pair-production dip, with a GCR/EGCR transition  $\\la 10^{18}$~eV; ii) if the EGCR source composition is roughly similar to that of GCRs, i.e., a substantial faction of heavier nuclei from the ambient medium around the source are accelerated together with protons, then the observed CR spectrum can be reproduced equally well with a less steep injection spectrum, $x\\simeq 2.2$--2.3, and the ankle can be interpreted as the GCR/EGCR transition (around $3\\,10^{18}$~eV).  While an ankle-like transition is \\emph{a priori} more natural than a knee-like transition, the rough coincidence between the end of the GCR component and the beginning of the EGCR component may arise naturally around $10^{18}$~eV, as a result of the shortage of high-energy GCRs (due to either acceleration or propagation effects) and the suppression of low-energy EGCRs because of expansion losses and/or magnetic field filtering \\cite{Aloisio2005,Lemoine+05}. Also, while a very steep spectrum and a discrimination of heavy nuclei by the acceleration process is \\emph{a priori} less natural, it has been claimed that strong extended magnetic fields around EGCR sources might substantially increase the path of highly charged nuclei compared to protons and thus tend to suppress heavy nuclei, turning an originally mixed composition into an effective pure proton source, with a subsequent phenomenology similar to that of the genuinely pure-proton scenario \\cite{Gunter}.  However that may be, we do not focus here on the theoretical plausibility of the different models, which remains somewhat subjective at this stage, but rather investigate the two main scenarios  from the point of view of composition. In particular, we make  definite predictions about observable features in the GCR/EGCR transition region for both models and compare our results with the available data.  In Sect.~2, we summarize the different composition hypotheses and briefly recall our standard scheme for the photo-disintegration of ultra-high-energy nuclei. In Sect.~3, we compare the propagated spectra obtained within the two main scenarios investigated, their implications for the remaining galactic component and the phenomenological limits of both models. We characterize the transition in terms of composition observables in Sect.~4 and compare our results with the available data. Finally, we discuss our results and conclude in Sect.~5. ",
        "conclusions": "In this paper, we have studied the phenomenology of the GCR/EGCR transition within two types of models for the extragalactic component: model A assumes a mixed source composition, with nuclear abundances similar to that of the GCRs, while model B assumes that the EG sources accelerate only protons. From the point of view of the CR energy spectrum, both models appear to be able to reproduce the high-energy CR data equally well, although with a different source spectrum and a correspondingly different interpretation of the ankle. Within model A, the source spectrum is typically in $E^{-2.3}$, and the ankle marks the transition from GCRs to EGCRs. Within model B, the source spectrum is steeper, in $E^{-2.6}$ or $E^{-2.7}$, and the ankle is fitted all the way down to $10^{18}$~eV by the EGCR component alone, exhibiting a characteristic dip due to the pair production interactions of ultra-high-energy protons with the cosmological microwave background.  Although we have assumed that the intensity of the sources did not evolve in the last few Gyr, none of the results presented here would be significantly modified if the sources actually evolved with redshift, either in number or in intensity. The exact slope of the required source spectrum could be slightly modified, but model A would also require a significantly harder spectrum, i.e. less luminous sources (from the low-energy extrapolation). We also studied the phenomenological limits of the models, and showed that the main features of model A were quite robust to a reasonable modification of the source composition, while the pair-production dip interpretation of the ankle provided by model B would be destroyed if a relatively small fraction of He or heavier nuclei were also accelerated in (and escaped from) the sources.  Since the available data on the CR energy spectrum are not accurate enough yet to discriminate between models A and B, we focused on the constraints that may be derived from composition estimates. We considered two complementary composition-related observables, $X_{\\max}$ and $\\langle \\ln A\\rangle$, accessible respectively to fluorescence detectors (e.g., HiRes, Auger) or ground arrays (e.g., AGASA, Auger). We obtained very different predictions for the evolution of these observables with energy, mostly because of the different scenarios implied for the transition from a heavy GCR component to EGCRs  : model B requires that the transition towards a light (pure proton) EGCR component be completed at an energy as low as $10^{18}$ eV, whereas the heavy GCR component extends up to $3-6\\,10^{18}$ eV in model A. The early GCR/EGCR transition of model B also implies an even steeper final drop of the Galactic component if a break of the source spectrum occurs at an energy larger than $10^{17}$~eV (remind that such a break is required for energetics reasons and possibly also not to violate the proton fraction measured at $10^{17}$~eV).  In any case, model B predicts a dramatic change of composition between $10^{17}$ and $10^{18}$~eV, which we have shown in Sect.~\\ref{sec:compoConstraints} to be disfavored by composition analysis of  most experiments expect for the very high HiRes-Mia elongation rate between $10^{17.5}$ and $10^{18}$ eV (when no low energy cut is applied). By contrast, the smooth transition implied by the mixed EGCR source composition of model A appears more in conformity with most currently available data.  It should be acknowledged, however, that the composition measurement remain difficult at high energy, and in particular the absolute normalisation of the $X_{\\max}$ observable is hadronic-model dependent. While the predictions of model A are in excellent quantitative agreement with the current data, an important result of our study is that the \\emph{shape} of the energy evolution of $X_{\\max}$ is also very different between both models. A distinctive prediction of model A is the presence of an ``s'' feature in the $X_{\\max}$ curve, corresponding to ``delay in lightening'' of the high-energy CRs, directly related to the heavy nuclei accelerated at the sources. This feature may be tested by future, higher-statistics experiments such as Auger, (to some extent) independently of the assumed hadronic interaction model.  The understanding of the origin of the ankle is very important for the phenomenology of not only the ultra-high-energy CRs, but also the Galactic ones. We wish to stress here that it cannot be obtained by the study of the sole spectrum, but benefits a lot from composition studies. It was the purpose of this paper to derive general predictions within the framework of two important and recently discussed models, independently of any consideration about their plausibility or the actual sources and acceleration mechanism involved. We argued elsewhere that model A is from our point of view more natural than model B, because i) a mixed-composition source is expected for processes accelerating particles out of the interstellar medium, ii) a source spectrum in $E^{-2.3}$ seems easier to interpret (notably in terms of relativistic shock acceleration) than a steeper spectrum in $E^{-2.6}$, which would also lead to an energetics problem if not cut at low-energy, and iii) an ankle-like feature is \\emph{a priori} natural for a transition between two CR components. However, we should not underestimate the argument that the ankle is very well reproduced by a single component of extragalactic protons over its whole energy range, with a very limited number of free parameters \\cite{Berezinsky+02,Berezinsky+04}. We nevertheless note here that the pure proton model meets some difficulties to account for the available composition measurements. It is a fact, however, that the experimental data are still quite uncertain and do not lead to a completely coherent picture above $10^{18}$~eV, even though they all favour a significant fraction of non-proton nuclei at the ankle and a smooth GCR.EGCR transition. Future experiments with larger statistics and higher resolution will be needed to fully understand the origin of the ankle: Kascade-Grande \\cite{Grande} should be able to measure the spectrum and constrain the composition above $10^{17}$~eV, where the transition is expected to begin for the two models; and the Pierre Auger Observatory will be able to measure the elongation rate above $5\\,10^{17}$~eV using the hybrid detection technique with a very high statistics above $10^{19}$~eV. The joint use of the fluorescence and ground array techniques will also be crucial to constrain the hadronic models, which are key to detailed understanding of the high-energy CR composition. Furthermore, the very high expected statistics coupled with the full sky coverage provided by the two sites of the Auger (North and South) should reveal the sources of UHECRs and thus provide the last ingredient necessary to the understanding of their origin.  Finally, we briefly comment on the possible influence of extragalactic magnetic fields (EGMF), which have been neglected here. Definite predictions cannot be proposed, mostly because we still lack strong constraints from both observational and theoretical studies about the strength, the coherence length and the structure of the EGMF (see for instance the large differences obtained for similar approches of magnetic fields in large scale structures: \\cite{Sigl05,Dolag05}). Qualitatively, however, we noted above that a magnetic horizon effect is expected at low energy (\\cite{Lemoine+05,Parizot04}), which could affect the extragalactic spectrum and make the GCR/EGCR transition sharper. This would thus amplify the specific features of model B described above, and possibly make it even more difficult for it to account for the composition-related data. Likewise, in the case of model A, the magnetic horizon would be rigidity-dependent and thus show at higher energy for heavier nuclei. The mixed composition would thus get heavier with increasing energy (with the heaviest nuclei entering the horizon at higher energy), which could partly (depending on the intensity of the magnetic field) counterbalance the steepening of the $X_{\\max}$ evolution curve discussed above in the case of a pure proton composition, and also make the characteristic ``s'' shape of model A more pronounced if heavy nuclei start becoming abundant in the mixed composition regime (corresponding to the flattening of the $X_{\\max}$ evolution; see above). In any case, whatever the effect of the hypothetic magnetic horizon, it would not lead to a confusion between the predictions of models A and B since the energy scale of the transition would still remain different. Furthermore the expected $X_{\\max}$ features (steep-flatter-steep) should remain present unless nuclei enter the horizon at energies above their photo-disintegration threshold, which would require quite a large EGMF (larger than used in \\cite{Lemoine+05}). A large EGMF could also increase the path length of the UHE nuclei at high energy, as studied in \\cite{Gunter} and thus shift the photo-disintegration cut-off of the different types of nuclei to a lower energy. However, as we have shown above, the predicted spectrum is not very sensitive to the presence of nuclei at the highest energies. In the case of our generic composition, the proton component already represents $70\\%$ of the total flux at $1.5\\,10^{19}$~eV, and a lighter composition could thus provide a similarly good fit of the data. In conclusion, although strong extragalactic magnetic fields would certainly have some impact on the phenomenology of the high-energy CRs, our general results should not be strongly affected.  \\subsubsection*{Acknoledgements:} We thank A. G. Mariazzi, T. Dova, A. Watson, D. Heck, S. Sciutto, R. Engel, M. Roth, M. Ave, J. Cronin, T. Yamamoto and G. Pelletier for valuable discussions. This work was supported in part by NSF through grant PHY-04057069 and the KICP under NSF PHY-0114422."
    },
    "0512/astro-ph0512290.txt": {
        "abstract": "{The debate on the oxygen abundances of metal-poor stars has its origin in contradictory results obtained using different abundance indicators. To achieve a better understanding of the problem we have acquired high quality spectra with the Ultraviolet and Visual Echelle Spectrograph at VLT, with a signal-to-noise of the order of 100 in the near ultraviolet and 500 in the optical and near infrared wavelength range. Three different oxygen abundance indicators, OH ultraviolet lines around 310.0\\,nm, the \\oifor\\ line at 630.03\\,nm and the \\oi\\ lines at 777.1-5\\,nm were observed in the spectra of 13 metal-poor subgiants with $-3.0\\le\\mathrm{[Fe/H]}\\le-1.5$. Oxygen abundances were obtained from the analysis of these indicators which was carried out assuming local thermodynamic equilibrium and plane-parallel model atmospheres. Abundances derived from \\oi\\ were corrected for departures from local thermodynamic equilibrium. Stellar parameters were computed using \\teff-vs-color calibrations based on the infrared flux method and Balmer line profiles, Hipparcos parallaxes and \\ion{Fe}{ii} lines. [O/Fe] values derived from the forbidden line at 630.03\\,nm are consistent with an oxygen/iron ratio that varies linearly with [Fe/H] as $\\mathrm{[O/Fe]}=-0.09(\\pm0.08)\\mathrm{[Fe/H]}+0.36(\\pm0.15)$. Values based on the \\oi\\ triplet are on average $0.19\\pm0.22\\dex$(s.d.) higher than the values based on the forbidden line while the agreement between OH ultraviolet lines and the forbidden line is much better with a mean difference of the order of $-0.09\\pm0.25\\dex$(s.d.). In general, our results follow the same trend as previously published results with the exception of the ones based on OH ultraviolet lines. In that case our results lie below the values which gave rise to the oxygen abundance debate for metal-poor stars. ",
        "introduction": "Oxygen plays a major role in astrophysics through its high cosmic abundance. It is predominantly produced and ejected back into the interstellar medium in connection with core-collapse supernovae (SNe II). By comparing oxygen abundances in stars of different ages with those of iron, which is produced both in SNe II and in thermo-nuclear supernovae of white dwarfs (SNe Ia), one can probe the star formation rate history and initial mass function of a galaxy as well as the physics of supernovae \\citep[e.g.][]{Tinsley80,Wheeler89}. Due to the longer time-scale of Fe production in SNe Ia , one can expect an oxygen over-abundance relative to iron at low metallicity: \\ofe\\,$> 0$\\footnote{The abundance ratios are defined by the customary [X/Fe]\\,$=\\log{(N_{\\rm X} / N_{\\rm Fe})_*} - {\\log{(N_{\\rm X} / N_{\\rm Fe})_\\odot}} $.}.  While such an oxygen over-abundance in metal-poor stars has long been known to exist \\citep{Conti67}, the exact amount is still contested with no unanimous agreement apparently in sight. Broadly speaking the derived \\ofe\\ values depend on which oxygen diagnostic is employed in the analysis and what type of stars is observed. There are four different types of spectral transitions which have been utilised for this purpose: the forbidden \\oifor\\ 630.0 and 636.3\\,nm\\ lines, high excitation \\oi\\ lines, in particular the \\oi\\ 777.1-5\\,nm\\ triplet, OH A-X electronic lines in the ultraviolet (UV) and OH vibration-rotation lines in the infrared (IR). The \\oi\\ and OH lines have mainly been used in metal-poor turn-off stars (\\teff\\,$\\ga 6000$\\,K) while the O abundances in halo giants (\\teff\\,$\\la 5000$\\,K) are typically based on the \\oifor\\ lines. The \\oifor\\ lines suggest a quasi-plateau at \\ofe\\,$\\sim +0.5$ for \\feh\\,$\\le -1$ \\citep[e.g.][]{Barbuy88,Sneden91,Nissen02}. This has traditionally been advocated as the correct metallicity trend, in particular since the first studies based on the OH lines in the UV gave consistent results \\citep{Bessell84,Bessell91,Nissen94}. More recently, the OH lines in the IR have also been studied in a few stars with a similar result, at least when restricting to weak, higher excitation OH lines \\citep{Balachandran01,Melendez01}. Systematically higher \\ofe\\ values are normally derived when using the \\oi\\ triplet \\citep[e.g.][]{Abia89,Israelian98,Israelian01,Boesgaard99,Carretta00,Nissen02,Fulbright03}, sometimes interpreted as a steady linear increase in \\ofe\\ towards lower {\\feh}. The \\oi\\ triplet results have often been blamed on departures from local thermodynamic equilibrium (LTE), effects of inhomogeneities and/or an erroneous \\teff -scale in order to reconcile them with the \\oifor\\ results. The issue of oxygen abundances in halo stars received much renewed attention following the surprising results of \\citet{Israelian98,Israelian01} and \\citet{Boesgaard99} who found a near linear trend in \\ofe\\ with a significant slope of about $-0.4$ towards lower metallicity from OH lines in the UV in turn-off stars and a couple of subgiants.  All of the above-mentioned oxygen diagnostics have their advantages and disadvantages. The \\oi\\ triplet is susceptible to non-LTE effects (NLTE) \\citep[] [and references therein]{Kiselman01} and is, due to the high excitation potential, quite sensitive to the adopted {\\teff}. The \\oifor\\ line is immune to departures from LTE but is essentially undetectable in turn-off stars with \\feh\\,$\\la -2$ due to its weakness. The OH lines are very sensitive to the temperature structure in the stellar atmosphere. Given the much lower temperatures encountered in the optically thin layers in 3D hydrodynamical model atmospheres compared with classical 1D hydrostatic models for metal-poor stars \\citep{Asplund99}, large downward abundance corrections have been flagged for the OH lines in turn-off stars. The \\oifor\\ lines are also affected by such temperature inhomogeneities but to a lesser degree than the OH lines \\citep{Nissen02}, while the NLTE line formation of the \\oi\\ lines can be suspected to be dependent on the 3D atmospheric structure as well \\citep{Kiselman95, Asplund01,Asplund04}. Additional complications and confusion stem from the differences and uncertainties in the adopted \\teff\\ and \\logg , the use of \\fei\\ or \\feii\\ lines to derive \\feh , possible missing continuous opacities in the UV \\citep{Balachandran98,Bell01} and which solar oxygen abundance the stellar values are referenced to. Given the recent large downward revision of the solar O abundance to \\logo\\,$\\sim 8.7$ \\citep{Allende01,Asplund04}, this difference of about 0.2\\,dex compared with the previously advocated value \\citep{Anders89} will directly translate to a corresponding {\\em increase} in \\ofe .  Considering all these remaining uncertainties it is perhaps not surprising that no consensus has as yet been achieved in terms of the behaviour of \\ofe\\ with \\feh . It is noteworthy that significant differences between the different oxygen indicators only reveal themselves at low metallicities (\\feh\\,$\\la -2$), which is observationally very challenging due to the weakness of some of the spectral features. The fact that the different oxygen diagnostics are normally used in different types of stars constitutes a major problem in this regard. The only metal-poor stars for which it is possible to employ molecular as well as forbidden and permitted atomic lines are subgiants but also then exceptionally high-quality spectra are required. Until now, only two halo subgiants (\\object{HD\\,140283} and \\object{BD$+23\\degr3130$}, both with \\feh\\,$\\sim -2.4$) have been analysed with \\oifor , \\oi\\ and OH UV lines \\citep{Israelian98,Israelian01,Fulbright99,Cayrel01,Balachandran01,Nissen02}, with somewhat ambiguous results. Our best hope to finally resolve this important outstanding problem is offered by halo subgiants. In the present article we present such a study based on very high-quality UV and optical spectra for a sample of 13 subgiants with $-3.0 \\le {\\rm [Fe/H]} \\le -1.5$. ",
        "conclusions": "Until 3D hydrodynamic model atmospheres for subgiants are available which will make 3D-NLTE line formation computations possible, 1D-LTE analyses based on the forbidden 630.03\\,nm line of oxygen will probably provide the best estimates of stellar [O/Fe] when observable. These values are by far the most accurate, with typical errors of the order of $0.10$\\,dex. Note although that they may not be free of significant systematic errors associated with inhomogeneities in the stellar atmospheres. Our best [O/Fe] estimates suggest a value at a level of $0.55\\pm0.13$ for Pop II subgiants, relatively independent on metallicity; the best linear fit to these estimates is $\\mathrm{[O/Fe]}=-0.09(\\pm0.08)\\mathrm{[Fe/H]}+0.36(\\pm0.15)$.  Both [O/Fe] values based on the IR triplet lines and on the OH UV lines depart from the the {\\oifor}-based estimates. The departures from the {\\oifor} results are more significant for the abundances derived from the triplet lines, on average $0.19\\pm0.22\\dex$ (s.d.), than for the values from the molecular lines, $-0.09\\pm0.25\\dex$ (s.d.). In the case of the triplet lines, the departures decrease with increasing effective temperatures while in the case of the OH lines they increase, turning from being negative to positive. According to line computations for 3D model atmospheres of main-sequence stars, 3D abundance corrections are negative and decrease with decreasing effective temperature which is not in contradiction with the present observations. In general, 3D effects on abundances and departures from LTE depend on stellar parameters and this can be the reason for the fact that differences in [O/Fe] between the indicators depend on the fundamental stellar parameters.  In order to increase the accuracies of [O/Fe] the $S/N$ in the observations should be further improved and the effective temperatures should be determined more accurately. Accurate gravities are as well necessary to improve the values based on OH UV lines whose typical errors are of the order of $0.3\\dex$.  Our results for subgiants follow the trends found by many others for main-sequence stars and giants in the last few years with the exception of our OH-based results which are significantly lower."
    },
    "0512/astro-ph0512035_arXiv.txt": {
        "abstract": "{IRAS 08339+6517 is a luminous infrared and Ly$\\alpha$-emitting starburst galaxy that possesses a dwarf % companion object at a projected distance of 56 kpc. An \\ion{H}{i} tidal tail has recently been detected between both % galaxies, suggesting that about 70\\% of the neutral gas has been ejected from them. We present deep broad-band % optical images, together with narrow band H$\\alpha$ CCD images, and optical intermediate-resolution spectroscopy of % both galaxies. The images reveal interaction features between both systems and strong H$\\alpha$ emission in the inner % part of IRAS 08339+6517. The chemical composition of the ionized gas of the galaxies is rather similar. The analysis % of their kinematics also indicates interaction features and reveals an object that could be a candidate tidal dwarf % galaxy or a remnant of an earlier merger. Our data suggest that the \\ion{H}{i} tail has been mainly formed from % material stripped from the main galaxy. We find weak spectral features that could be attributed to the presence of % Wolf--Rayet stars in this starburst galaxy and estimate an age of the most recent burst of around 4 -- 6 Myr. A more % evolved underlying stellar population, with a minimal age between 100 -- 200 Myr, is also detected and fits an % exponential intensity profile. A model which combines 85\\% young and 15\\% old populations can explain both the % spectral energy distribution and the \\ion{H}{i} Balmer and \\ion{He}{i} absorption lines presented in our spectrum. % The star formation rate of the galaxy is consistently derived using several calibrations, giving a value of $\\sim$9.5 % \\Mo\\ yr$^{-1}$. IRAS 08339+6517 does satisfy the criteria of a luminous compact blue galaxy, rare objects in the % local universe but common at high redshifts, being a very interesting target for detailed studies of galaxy evolution % and formation.}   \\titlerunning{SF and stellar populations in the LCBG IRAS 08339+6517} \\authorrunning{L\\'opez-S\\'anchez, Esteban \\& Garc\\'{\\i}a-Rojas} ",
        "introduction": "Determining the star formation history in galaxies is fundamental for the understanding of their evolution. In nearby % galaxies, such as those belonging to the Local Group, it is performed studying the stellar content via % color-magnitude diagrams (CMD) and stars as age tracers [see \\citet{Grebel99} and references there in]. However, in % distant galaxies it is only interpreted by means of spectral synthesis techniques. The analysis of the interstellar % medium (ISM) in and around galaxies complements the study of their star formation histories. Star formation activity is stronger in starburst galaxies, where the dominant and very young stellar population could % even hide the more evolved stellar population presented in it. The more extreme cases are the blue compact dwarf % (BCDs) galaxies, that have very low metal abundances and exhibit a global starburst activity. Although the majority % of these systems possesses an old underlying population with ages of several Gyrs \\citep{LT86}, a few of them do not % present it, suggesting that they are really young galaxies \\citep{IT99}. The best example of this kind of BCDs is I % Zw 18, a bona fide young galaxy in the local universe \\citep{IT04}. At intermediate and high redshifts a % heterogeneous class of vigorous starburst systems with luminosities around $L^*$ ($L^{\\star} = 1.0 \\times 10^{10} % L_{\\odot}$) are observed \\citep{GOK03}. They are designated as luminous compact blue galaxies (LCBGs) and their % evolution and nature are still open questions, being fundamental to get a sample of them in the local universe to % investigate the origin of their activity \\citep{WJS04}.  An important subset of starbursts are the so-called Wolf--Rayet (WR) galaxies, whose integrated spectra have a broad % emission feature around 4650 \\AA\\ that has been attributed to WR stars. This feature consists of a blend of % \\ion{N}{iii} $\\lambda$4640, \\ion{C}{iii}/\\ion{C}{iv} $\\lambda$4650 and \\ion{He}{ii} $\\lambda$4686 emission lines % \\citep{C91}, the last one being the most prominent line. The WR feature indicates the presence of a substantial % population of these kinds of massive stars whose ages are less than 6 Myr and offers the opportunity to study very % young starbursts \\citep{SV98}. Furthermore, they constitute the best direct measure of the upper end of the initial % mass function (IMF), a fundamental ingredient for studying unresolved stellar populations \\citep{SGIT00,PSGD02}. % Studying a sample of WR galaxies, \\citet{ME00} suggested that interactions with or between dwarf objects could be the % main star formation triggering mechanism in dwarf galaxies and noted that the interacting and/or merging nature of WR % galaxies can be detected only when deep and high-resolution images and spectra are available. The compilation of WR % galaxies performed by \\citet{SCP99} lists 139 members, but since then this number has increased % \\citep{PH00,BO02,CTS02,PSGD02,LTD03,Tran03,FCCG04,IPG04,PKP04,Kniazev04}, these galaxies have even been detected at % high $z$ \\citep{VMCGD04}. In this paper, we add a new member to the list of probable WR galaxies: the luminous % infrared galaxy IRAS 08339+6517, which, furthermore, could be also classified as a luminous compact blue galaxy % (LCBG).       IRAS 08339+6517 is at 80 Mpc (at that distance, 1$\\arcsec$ = 388 pc, assuming H$_0$=75 km s$^{-1}$ Mpc$^{-1}$) and was firstly reported in the IRAS Point Source Catalog (1986). \\citet{MAM88} performed a % multi-wavelength study and described the object as an exceptionally bright and compact starburst nucleus. % \\citet{GDL98} presented a detailed UV analysis, suggesting that O stars must be present in the ionizing cluster(s) of % the galaxy in order to explain the blue wing of the \\ion{C}{iv} $\\lambda$1550 and \\ion{Si}{iv} $\\lambda$1400 stellar % absorption lines, although these stars cannot be more massive than about 40 \\Mo. \\citet{KMH98} noted a P Cygni % profile in the Ly$\\alpha$ emission, which also showed two components indicating the chaotic structure of the % interstellar medium in this object. \\citet{MHK03} included IRAS 08339+6517 in their $HST$ UV and optical study of % Ly$\\alpha$ starbursts.  In a very recent paper, \\citet{CSK04} presented  VLA \\ion{H}{i} imaging of this starburst galaxy and found an % extended tidal structure in neutral hydrogen indicating that it is interacting with a nearby companion, 2MASX % J08380769+6508579 (which we designate as the companion galaxy throughout this paper). This feature provides an % evidence for tidally induced starburst episodes. \\citet{CSK04} estimated neutral hydrogen masses of $(1.1\\pm % 0.2)\\times 10^9$ \\Mo\\ and $(7.0\\pm 0.9)\\times 10^8$ \\Mo\\ for IRAS 08339+6517 and its companion galaxy, respectively, % and a mass of $(3.8\\pm 0.5)\\times 10^9$ \\Mo\\ in the tidal material between them. Consequently, $\\sim$ 70\\% of the % neutral gas has been removed from one or both galaxies.  We are developing a detailed analysis of Wolf-Rayet galaxies using both deep images and spectra in optical % wavelenghts. Our data show increasing evidences of the connection between the detection of WR features and the % finding of interaction signatures in this kind of starbursts \\citep{ME00,LSE03,LSER04a,LSER04b,LS06}. The % characteristics previously observed in IRAS 08339+6517 (strong \\Ha\\ emission, starburst nature and the detection of % an extended \\ion{H}{i} tidal tail in the direction of its companion dwarf galaxy) make it an ideal target to look for % the existence of WR stars. In this way, we have performed deep broad-band optical and narrow-band H$\\alpha$ CCD % images together optical intermediate-resolution spectroscopy of IRAS 08339+6517 and its dwarf companion galaxy. Our % main aim is to study their morphology and stellar populations, as well as the distribution, che\\-mi\\-cal composition % and kinematics of the ionized gas, and to check if WR stars are presented in their youngest bursts.    \\begin{figure*}[ht] \\includegraphics[angle=90,width=1\\linewidth]{IRAS08339figure1new.eps} \\caption{\\small{(a) Deep optical image of IRAS 08339+6517 and its companion galaxy, 2MASX J08380769+6508579, in $R$. % It has been saturated to reveal the weakest features. (b) The companion galaxy and (c) IRAS 08339+6517 in non-satured % $R$ images. The grayscale is in logarithmic scale in all three cases.}} \\label{figR} \\end{figure*} ",
        "conclusions": "We have used deep optical and \\Ha\\ imagery together deep optical intermediate-resolution spectroscopy to analyze the % morphology, colors, ages, stellar populations, phy\\-si\\-cal conditions, kinematics and chemical abundances of the % galaxy IRAS 08339+6517 and its dwarf companion galaxy, 2MASX J08380769+6508579. Our data reinforce the results of the % \\ion{H}{i} observations performed by \\citet{CSK04} that both objects are in interaction and that the \\ion{H}{i} tidal % tail seems to be formed mainly by material stripped from IRAS 08339+6517.  We have obtained the oxygen abundances of both galaxies using empirical calibrations. The O/H and N/O ratios of the % two galaxies are rather similar, suggesting a similar degree of chemical evolution of the system. We have analyzed % the chemistry of the ionized gas using different aperture sizes and found an apparent metallicity gradient between % the nucleus and the external zones of IRAS 08339+6517. The reddening coefficient also shows differences between the % central and external zones.  IRAS 08339+6517 shows important \\Ha\\ emission in its inner regions. IRAS 08339+6517 is not powered by an AGN but it % is a nuclear starbursting galaxy. We have detected at least two different popu\\-lations in the galaxy, the age of the % youngest one being around 4 -- 6 Myr. But the more evolved stellar population, with age older than 100 -- 200 Myr, % fits an exponential intensity profile. A model which combines 85\\% of a young (6 Myr) population with 15\\% of an old % (140 -- 200 Myr) population can explain both the spectral energy distribution and the H Balmer and \\ion{He}{i} % absorption lines observed in our spectrum. Furthermore, the NIR colors suggest that even an older stellar population, % with age upper than 1 -- 2 Gyr, is also present in IRAS 08339+6517. Its companion galaxy is practically dominated by % an old population with age $>$ 250 Myr, although it also hosts a recent $\\sim$ 6 Myr starburst in its external areas. % The Keplerian, ionized and cold dust masses derived for the main galaxy seem to be the usual for young starbursts.  Our spectra seem to show weak Wolf-Rayet features in IRAS 08339+6517, but they are located only in the central knot % of the galaxy. However, deeper spectra are needed to confirm the WR nature of this galaxy. We discuss that aperture % effects and localization of the bursts with WR stars could play a fundamental role in the detection of this sort of % massive stars in starburst galaxies.    We have derived the SFR of IRAS 08339+6517 using multi-wavelength correlations. The \\Ha, FIR and radio luminosities % give similar values, around 9.5 \\Mo\\ yr$^{-1}$, suggesting that the contribution of reddening by dust has been % properly estimated. The high SFR derived using X-ray luminosities implies massive stars formation in the last few % Myr, as seems to be supported by the high rate of supernova explosions and the probable detection of a large number % WR stars.  Knot B, located in the NW arm of the main galaxy, has the higher metallicity and possesses peculiar kinematics. It % could be a TDG candidate because its metallicity is higher than that expected for a dwarf galaxy with its % $B$-luminosity and because the kinematics of the ionized gas are decoupled from the general kinematic pattern. % However, it could also be a remnant of a previous early merger because of the peculiar morphology of the galaxy and % the existence of the \\ion{H}{i} tidal tail. Besides it, at the present time the interaction with the external % companion galaxy is much more prominent and probably the origin of the \\ion{H}{i} tidal tail.  Finally, IRAS 08339+6517 could be classified as a luminous compact blue galaxy (LCBG) because of its color, absolute % magnitude and surface brightness. There are very few local LCBGs nowadays detected but nearly half of them have % optical companions, present disturbed morphologies and/or are clearly interacting \\citep{GPW04}. An example of local % LCBG is Mkn~1087 \\citep{LSER04b}, that is in interaction with two nearby galaxies and shows bridges and tails that % connect dwarf surrounding objects with the main body. If interactions were the responsible of the activity in LCBGs, % it would indicate that they were perhaps more common at high redshifts, as the hierarchical galaxies formation models % predict [i.e., \\citet{KW93,Springel05}]. This fact would also support the idea that interaction with dwarf companion % objects could be an important trigger mechanism of the star formation activity in local starbursts."
    },
    "0512/astro-ph0512203_arXiv.txt": {
        "abstract": "We stack {\\it Spitzer} 24\\,$\\mu$m images for $\\sim$ 7000 galaxies with $0.1 \\le z < 1$ in the Chandra Deep Field South to probe the thermal % dust emission in low-luminosity galaxies over this redshift range. Through stacking, we can detect mean 24\\,$\\mu$m fluxes that are more than an order of magnitude below the individual detection limit.  We find that the correlations for low and moderate luminosity galaxies between the average $L_{\\rm IR}/L_{\\rm UV}$ and rest-frame $B$-band luminosity, and between the star formation rate (SFR) and $L_{\\rm IR}/L_{\\rm UV}$, are similar to those in the local Universe. This verifies that oft-used assumption in deep UV/optical surveys that the dust obscuration--SFR relation for galaxies with SFR $\\leq$ 20\\,M$_\\odot$\\,yr$^{-1}$ varies little with epoch. We have used this relation to derive the cosmic IR luminosity density from $z = 1$ to $z = 0.1$. The results also demonstrate directly that little of the bolometric luminosity of the galaxy population arises from the faint end of the luminosity function, indicating a relatively flat faint-end slope of the IR luminosity function with a power law index of $1.2\\pm0.3$. ",
        "introduction": "The rapid decay of the cosmic mean star formation rate (SFR) density from $z = 1$ to the present epoch has been convincingly established over the last decade (see, e.g., Hopkins~\\citeyear{Hopkins04} and references therein). The focus has now turned to characterizing the types of galaxies responsible for this decay. Rest-frame ultraviolet (UV) and optical emission-line studies indicate the decay since z $\\sim$ 1 is strongly influenced by the behavior of the relatively low-mass galaxies (e.g., Brinchmann \\& Ellis et al. ~\\citeyear{Brinchmann00}; Juneau et al. \\citeyear{Juneau05}; Bundy et al. \\citeyear{Bundy05}; Bauer et al. \\citeyear{Bauer05}; Wolf et al. \\citeyear{Wolf05}). On the other hand, surveys in the thermal infrared (IR) have found many galaxies of intermediate and high mass with intense, deeply obscured star formation (Franceschini et al. \\citeyear{Franceschini03}; Zheng et al. \\citeyear{Zheng04}; Bell et al. \\citeyear{Bell05}; P\\'erez-Gonz\\'alez et al. \\citeyear{Perez05}; LeFloc'h et al. \\citeyear{LeFloc'h05}). However, the role of obscured star formation in lower-mass galaxies remains unknown (although see Heavens et al. \\citeyear{Heavens04} for a powerful and complementary approach).   Among the many observational SFR estimators (see Kennicutt \\citeyear{Kennicutt98a} for a review), at high redshifts UV radiation is the most easily measured proxy for the SFR. It provides the basis for much of our understanding of the evolution of the cosmic SFR density (e.g., Madau et al. \\citeyear{Madau96}; Steidel et al. \\citeyear{Steidel99}; Sullivan et al. \\citeyear{Sullivan00}; Hopkins et al. \\citeyear{Hopkins01}; Schiminovich et al. \\citeyear{Schiminovich05}). Yet, young stars are usually born in dust-rich environments; the dust absorbs the vast majority of the UV light and re-radiates this energy in the thermal IR. Therefore, to obtain a complete census of the bolometric luminosity from young stars, observations in both the UV and IR are required (see, e.g., Gordon et al. \\citeyear{Gordon00}). However, the available infrared facilities have lacked the sensitivity and imaging resolution to explore obscured star formation in low-luminosity, relatively low-mass galaxies, except for nearby examples.  Thus, indirect estimates have had to be used for the IR outputs of such galaxies even at moderate redshift, using tools such as the empirical calibration of UV color to extinction (e.g., Meurer et al. \\citeyear{Meurer99}; Calzetti \\& Heckman \\citeyear{Calzetti99}), or the locally-observed trend between $L_{\\rm IR}/L_{\\rm UV}$ and luminosity (showing that more luminous galaxies tend to have higher extinction; e.g., Wang \\& Heckman \\citeyear{Wang96}; Bell \\citeyear{Bell03}; Buat et al. \\citeyear{Buat05}). These approaches have been applied to observations of high-redshift galaxies to derive the SFR per unit comoving volume (e.g., Adelberger \\& Steidel \\citeyear{Adelberger00}; Hopkins \\citeyear{Hopkins01}).   \\begin{figure*} \\centering \\includegraphics[width=0.35\\textwidth,clip]{f1a.eps}\\hskip 6mm \\includegraphics[width=0.35\\textwidth,clip]{f1b.eps} \\caption{Illustration of Spitzer/MIPS 24\\,$\\mu$m images before and after PSF subtraction of individually-detected sources.  {\\it Left}: an image section (10$\\arcmin \\times 10\\arcmin$) in the CDFS; {\\it Right}: the corresponding residual image after removing all detected sources brighter than 83\\,$\\mu$Jy (at the 80\\% completeness).  Two images are shown in the same greylevel scale.  Sources seen in the residual images are fainter than 83\\,$\\mu$Jy.  } \\label{exampleimage} \\end{figure*}    The goal of this paper is to measure directly the average IR outputs of low-luminosity galaxies at intermediate redshift. We use a deep 24\\,$\\mu$m map of the Chandra Deep Field South (CDFS) from the Multiband Imaging Photometer on {\\it Spitzer} (MIPS: Rieke et al. \\citeyear{Rieke04}). The 24\\,$\\mu$m map is limited by both photon and confusion noise (Dole et al. \\citeyear{Dole04}). We find that stacking on the position of known intermediate redshift galaxies substantially reduces both sources of noise, allowing secure detection of average flux densities substantially below the conventional confusion limit. We have stacked 24\\,$\\mu$m images of several thousand intermediate- and low-luminosity galaxies at $0.1 \\le z < 1$. We combine the resulting detections with the COMBO-17 data set to investigate the effects of dust obscuration in galaxies with $m_R < 24$ mag out to $z \\sim$ 1. We find no evidence for evolution in the relationship between obscuration and the SFR over this range of redshift. Our results suggest that the IR luminosity function has a relatively flat slope (power law index of $1.2\\pm 0.3$) toward low luminosities, at least out to $z \\sim$ 1.  The paper is organized as follows. Section 2 describes the data and the selection of the galaxy sample. Section 3 addresses the methods for stacking the images and describes tests of their validity. Section 4 shows the results. Section 5 uses these results to refine estimates of the cosmic SFR density from z $\\sim$ 1 to the present epoch. Our work is summarized in Section 6. Throughout the paper we adopt a cosmology with $H_0$\\,=\\,70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$\\,=\\,0.3 and $\\Omega_{\\Lambda}$\\,=\\,0.7. ",
        "conclusions": "Estimating the importance of dust-obscured star formation for intermediate-redshift normal, or even low-luminosity galaxies on a galaxy-by-galaxy basis is impossible through direct thermal IR observations with current technology.  Using 24\\,$\\mu$m images from Spitzer, we show that one can determine mean thermal IR fluxes and hence 10-fold fainter average obscured star formation rates by stacking 24\\,$\\mu$m images centered on the optical positions of known intermediate-redshift galaxies.  We use a sample of galaxies from the COMBO-17 photometric redshift survey of the Extended Chandra Deep Field South, which provides astrometry, photometric redshift and rest-frame 2800\\AA, $U$, $B$, and $V$-band absolute magnitudes for thousands of $z \\la 1$ galaxies.  We stack MIPS 24\\,$\\mu$m images for subsamples of galaxies in redshift slices and rest-frame $B$-band luminosity bins, allowing detection of average 24\\,$\\mu$m fluxes down to $<$\\,10\\,$\\mu$Jy, an order of magnitude deeper than those accessible on a galaxy-by-galaxy basis.  The mean total IR luminosity of these galaxy subsets is estimated from the observed 24\\,$\\mu$m luminosity, taking into account the uncertainty of IR SED shapes. Analogous mean UV luminosities are derived, and average SFRs are estimated from the UV and IR luminosities of each subsample. We use these data to examine the correlations among optical luminosity, dust obscuration, SFR and redshift.  We find that the correlation between dust obscuration, i.e. the ratio of $L_{\\rm IR}/L_{\\rm UV}$, and rest-frame $B$-band luminosity seen in local star-forming galaxies holds over all redshifts $z \\la 1$, with brighter galaxies showing a higher $L_{\\rm IR}/L_{\\rm UV}$ ratio. Our averaged 24\\,$\\mu$m detections show directly that even in low luminosity galaxies (to 0.05\\,$L^\\ast$) the majority of the bolometric luminosity from young stars is re-radiated in the thermal IR. Nonetheless, the decrease of $L_{\\rm IR}/L_{\\rm UV}$ ratio with decreasing $L_{\\rm B}$ implies that globally star formation in faint objects is lower than the estimate one would derive the level of dust obscuration typical of normal galaxies.  We explore the correlation between average $L_{\\rm IR}/L_{\\rm UV}$ ratios and SFRs for the different optically-selected subsamples. Different subsamples populate different parts of the $L_{\\rm IR}/L_{\\rm UV}$--SFR plane; however, our data indicate that this correlation does not evolve much between $z = 1$ and the present day.  In closing, we briefly consider some of the factors determining the degree of dust obscuration indicated by $L_{\\rm IR}/L_{\\rm UV}$ and roughly parameterized as an optical depth, $\\tau$. The optical depth of a galaxy $\\tau \\propto \\Sigma_{\\rm gas}\\,Z\\,\\alpha$, where $\\Sigma_{\\rm gas}$ is gas density, $Z$ is metallicity and $\\alpha$ is a geometric term to account for the gas and dust distribution relative to massive stars.  Star formation intensity is strongly correlated with gas density, e.g., the Schmidt law (Kennicutt~\\citeyear{Kennicutt98b}). Since our results are drawn from a large number of galaxies, geometric effects may cancel out, to first order. Then the correlation of $L_{\\rm IR}/L_{\\rm UV}$, or  $\\tau$ with SFR (i.e., more intense star-forming environments show larger dust obscuration) may suggest that gas density drives $L_{\\rm IR}/L_{\\rm UV}$ correlation to a much greater extent than $Z$, at least over the magnitude and redshift ranges we consider in this work.  Independent from the infrared luminosity function, we estimate the cosmic infrared luminosity density from the rest-frame $B$-band luminosity function. An increase by a factor of $9\\pm$3 is found for the comoving infrared luminosity density from $z$\\,=\\,0.1 to $z$\\,=\\,1,  consistent with Le Floc'h et al. (\\citeyear{LeFloc'h05}). Based on our estimate of the comic infrared luminosity density, it is suggested that the infrared luminosity function of Le Floc'h et al. (\\citeyear{LeFloc'h05}) is well determined at intermediate redshifts, supporting their claim that faint end slope of the luminosity function is relatively flat. Our result suggests a power law index of $1.2\\pm 0.3$."
    },
    "0512/astro-ph0512258_arXiv.txt": {
        "abstract": "{We report the initial results of a VLT/NACO high spatial resolution imaging survey for multiple systems among 58 M-type members of the nearby Upper Scorpius OB association. Nine pairs with separations below 1\\arcsec\\, have been resolved. Their small angular separations and the similarity in the brightness of the components ($\\Delta$Mag$_{K}<$1 for all of them), indicate there is a reasonable likelihood several of them are true binaries rather than chance projections. Follow-up imaging observations with WHT/LIRIS of the two widest binaries confirm that their near-infrared colours are consistent with physical very low-mass binaries. For one of these two binaries, WHT/LIRIS spectra of each component were obtained. We find that the two components have similar M6-M7 spectral types and signatures of low-gravity, as expected for a young brown dwarf binary in this association. Our preliminary results indicate a possible population of very low-mass binaries with semimajor axis in the range 100~AU--150~AU, which has not been seen in the Pleiades open cluster. If these candidates are confirmed (one is confirmed by this work), these results would indicate that the binary properties of very low-mass stars and brown dwarfs may depend on the environment where they form. ",
        "introduction": "Multiple systems and their physical properties are important testimonies of the formation and early stages of evolution of a class of astrophysical objects. For that reason, the study of multiplicity among ultracool dwarfs has been an intense field of research over the last few years \\citep[see e.g][]{1999AJ....118.2460B,1999Sci...283.1718M,2000ApJ...529L..37M,2003ApJ...594..525M,2003AJ....125.3302G,2004A&A...424..213B,2003AJ....126.1526B,2003ApJ...586..512B,2005Natur.433..286C,2003ApJ...587..407C,2002ApJ...567L..53C,2003IAUS..211..233J,2005astro.ph..9134J,2005ApJ...633..452K,2005ApJ...621.1023S}. While all these studies have reported a cut-off in the distribution of separation at about 30~AU, consistent with a dynamical formation mechanism involving gravitational interactions, a still very small but increasing number of wide binary candidates are reported in the field \\citep{1998ApJ...509L.113M, 2005A&A...439L..19P,2004A&A...427L...1F,2004AJ....128.1733G,2005A&A...440L..55B} and in young associations \\citep{2004ApJ...614..398L,2005A&A...438L..25C,2005A&A...435L..13N}. Together with the relatively high frequency of multiple systems reported by the above mentioned authors, these wide multiple systems challenge the ejection models and show that this scenario, although certainly at work, can probably not explain the formation of the majority of the very low mass objects.  A recent study performed by \\citet{2005ApJ...633..452K} in Upper Scorpius has revealed a significant fraction of visual binaries among late type members of the association. The three candidates they resolve among their twelve targets have small separations (less than 18~AU) and flux ratios close to unity, thus similar to the overall field and Pleiades binary populations. More recently \\citet{2005ApJ...633L..41L} reported a wide binary (at 130~AU) discovered serendipitously during a spectroscopic survey. The existence of such a fragile wide pair at the age of Upper Scorpius \\citep[5~Myr][hereafter Usco]{2002AJ....124..404P}, suggests that some very low-mass objects do not form via ejection. A detailed knowledge of the properties of wide mutliple systems in different environments would provide key constraints regarding the contribution of ejection to the formation of very low mass objects. In this paper, we report the initial results of a systematic search for multiple systems among the very low mass stars and brown dwarfs members of the young Upper Scorpius OB association with the NACO adaptive optics system on the VLT. In section \\ref{observations}, we present the observations and processing of the data. In section \\ref{dataanalysis} we describe the analysis of the data and the discovery of multiple system candidates. In section \\ref{comp} we discuss the companionship of the candidates, and in section \\ref{discussion} we discuss the results in the context of the current models of formation and in comparison with previous studies. ",
        "conclusions": "From our high angular resolution survey among a sample of 58 ultracool dwarfs in Upper Scorpius, we report 9 new binary candidates, among which 2 might be triples. Complementary near-infrared imaging and spectroscopic observations indicate that one of the wide multiple systems (DENIS161833AB) is a physical pair of M7--M7.5 brown dwarfs at a high level of confidence. Independent observations by \\citet{2005ApJ...633L..41L} lead to the same conclusion. Second epoch imaging and spectroscopic observations are required to confirm the multiplicity of the eight other candidates, but the near-infrared colours of one of them (USCO-160028.5AB) are consistent with a companion of similar spectral type. All binary candidates have differences of magnitude corresponding to mass ratio close to unity, confirming a preference for equal mass systems similar to what was observed up-to-now for field and Pleiades objects. The distribution of separations is extending up to separations of 100~AU, contrasting with the cut-off observed at 30~AU for field and Pleiades ultracool objects. With one third (3/9) of the candidates having separations greater than 100~AU, two third (6/9) in the range of separation between 0--30~AU, and two triple systems, the properties of the new candidates contrast with the properties of multiple ultracool dwarfs from the field or from the Pleiades, but the multiplicity of 8 of them must be confirmed before any firm conclusion can be drawn."
    },
    "0512/physics0512115_arXiv.txt": {
        "abstract": "We apply the Smaller ALignment Index (SALI) method to a 4--dimensional mapping of accelerator dynamics in order to distinguish rapidly, reliably and accurately between ordered and chaotic motion. The main advantage of this index is that it tends {\\it exponentially} to zero in the case of chaotic orbits, while it fluctuates around non--zero values in the case of quasiperiodic trajectories. Thus, it avoids the notorious ambiguities concerning the eventual convergence of (maximum) Lyapunov exponents to (positive) non-zero values. Exploiting the different behavior of SALI in these two cases we produce phase space `charts' where regions of chaos and order are clearly identified. Evaluating the percentage of chaotic and escaping orbits as a function of the distance from the origin we are able to estimate rapidly and accurately the boundaries of the {\\it dynamical aperture} of a proton beam, passing repeatedly through an array of magnetic focusing elements. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512544_arXiv.txt": {
        "abstract": "{% } {% Our objective is to prove that integrated optics (IO) is not only a good concept for astronomical interferometry but also a working technique with high performance. } {% We used the commissioning data obtained with the dedicated K-band integrated optics two-telescope beam combiner which now replaces the fiber coupler MONA in the VLTI/VINCI instrument. We characterize the behaviour of this IO device and compare its properties to other single mode beam combiner like the previously used MONA fiber coupler. } {% The IO combiner provides a high optical throughput, a contrast of $89$\\% with a night-to-night stability of a few percent. Even if a dispersive phase is present, we show that it does not bias the measured Fourier visibility estimate. An upper limit of $5\\times10^{-3}$ for the cross-talk between linear polarization states has been measured. We take advantage of the intrinsic contrast stability to test a new astronomical prodecure for calibrating diameters of simple stars by simultaneously fitting the instrumental contrast and the apparent stellar diameters. This method reaches an accuracy with diameter errors of the order of previous ones but without the need of an already known calibrator. } {% These results are an important step of integrated optics, since they prove its maturity in an astronomical band where the technology has been specially developed for astronomical conveniences. It paves the road to incoming imaging interferometer projects. } ",
        "introduction": "Installed at the heart of the Very Large Telescope Interferometer \\citep[VLTI,][]{Glindemann-2003}, the VINCI instrument combines coherently the light coming from two telescopes in the infrared K band. Among the most impressive astrophysical results is the measurements of diameters of very low mass stars \\citep{segransan-2003}, the oblateness of the fast rotating star Achernar \\citep{Domiciano-2003}, the calibration of the brightness-color relation of Cepheids \\citep{Kervella-2004b} and its complementarity with asteroseismology to constrain the stellar structure \\citep{Pijpers-2003}. All of them benefit from high accuracy interferometric measurements, achieved in the near-infrared by modal filtering of the corrugated wavefront and real time monitoring of the stellar flux injection \\citep{Foresto-1997,Tatulli-2004}.  \\citet[paper I:][]{Malbet-1999} have suggested to combine beams with planar integrated optics (IO) components to take benefit of strong spatial filtering, stability and compactness. Afterwards, this technique has been successfully tested in laboratory \\citep[papers II and III:][]{Haguenauer-2000b,Berger-1999} and on the sky \\citep[paper IV:][]{Berger-2001}.  In the framework of the IONIC activities, a collaboration was initiated in 2003 between LAOG and the ESO interferometry group to implement an integrated optics beam combiner operating in the K band on the VLTI, sharing the VINCI optical interface. This proposal followed the study of the IO techniques towards the K band \\citep[paper V:][]{Laurent-2002} and a previous collaboration between the two institutes concerning an IO H band combiner \\citep{LeBouquin-2004}.  \\begin{figure*} % \\centering \\includegraphics[width=0.9\\textwidth]{Fig1.eps} \\caption{ {\\bf a:} Layout of the VINCI instrument where the MONA fiber coupler is replaced by the \\dtk{} Integrated Optics coupler. The LISA detector array has been expanded to show the four pixels: Pa and Pb monitor the photometry and the fringes are recorded on I1 and I2. {\\bf b:} Picture of the \\dtk{} coupler in front of the VINCI off-axis parabola. {\\bf c:} Sketch of the component with the two inputs (bottom) and the four outputs (top). {\\bf d:} Raw interferometric output I1 obtained in the K band on 88~Aqr with the siderostats. Note the large fringe contrast even though the photometry is not calibrated.} \\label{fig:implantation_2TK} \\end{figure*} %  In this paper, we report the validation of the new VINCI setup, equipped with a two-telescope IO coupler for the K band. In Sect.~\\ref{sec:commissioning}, we summarize the instrumental context, the observations and the data reduction techniques. Results are presented and discussed in Sect.~\\ref{sec:results}. The intrinsic instrumental stability allows us to validate a new interferometric calibration technique, based on the simultaneous fit of the instrumental contrast and stellar apparent diameter as detailed in Sect.~\\ref{sec:new_contrast}. Finally, perspectives of Integrated Optics in the framework of optical interferometry are discussed in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have characterized the new setup of the VINCI instrument, equipped with an Integrated Optics component for the K band. First, we show that VINCI keeps the same limiting magnitude. Coupled with a good stability and a high instrumental contrast, it makes VINCI an interesting instrument to perform observations on faint objects. Then, we have explored its chromatic response by Fourier Transform Spectrometry. A dispersive phase due to the fibers is present but does not bias the measured square visibility by the classical Fourier estimation. Nevertheless, careful attention has to be paid to the selection criteria of the Wavelet algorithm developed by \\citet{Kervella-2004} because the spectro-temporal fringe shape is slightly spread over more frequencies and times. Finally, we find an upper limit of $5\\times10^{-3}$ for the cross-talk between linear polarization states.  The intrinsic stability of the whole instrumental chain VLTI + VINCI + \\dtk{} allows us to try a new calibration technique, based on the simultaneous fit of the atmospheric contrast and stellar apparent diameter. We validate it with two well known stars. The recovered fit parameters are all consistent. We emphasize that this technique thus proves the interferometric quality and stability of VINCI.  These results are an important step for the development of Integrated Optics for astronomical interferometry. Previous works were all performed to validate the technologies or to test the performances in shared risk programs. For the first time, a component has been designed, manufactured and commissioned to answer an astrophysical institute request. Because of the lack of telecommunication or metrology applications at $2\\mu$m, the K band component used has been especially developed for astronomical convenience, proving the maturity of the technique.  In the future, the goal is to combine the entire VLTI array and to disperse the light in order to have spectro-imaging capabilities. Integrated optics is certainly a promising solution \\citep{Kern-2003}. The compactness of the planar optical component allows one to combine many beams in the same chip, which drastically reduces the instability and the required alignments. The observational strategies (number of baselines, wavelength, combination scheme...) can be adapted to the object thanks to the ``plug and play'' ability of IO combiners. Output beams of the planar component can act as the input slit of a spectrograph, avoiding complex anamorphic optics.  In this context, we develop a IO chip which combines four beams with a very photon efficient concept: it allows to measure the six complex coherencies with only 24 pixels and without external OPD modulation \\citep{LeBouquin-2004b}. This component is already under tests at the LAOG optical bench and could be a key part for incoming imaging interferometer projects such as VITRUV \\citep{Malbet-2004}."
    },
    "0512/astro-ph0512634_arXiv.txt": {
        "abstract": "Precise radial velocity measurements have led to the discovery of $\\sim170$ extrasolar planetary systems.  Understanding the uncertainties in the orbital solutions will become increasingly important as the discovery space for extrasolar planets shifts to planets with smaller masses and longer orbital periods.  The method of Markov chain Monte Carlo (MCMC) provides a rigorous method for quantifying the uncertainties in orbital parameters in a Bayesian framework (Ford 2005a).  The main practical challenge for the general application of MCMC is the need to construct Markov chains which quickly converge.  The rate of convergence is very sensitive to the choice of the candidate transition probability distribution function (CTPDF).  Here we explain one simple method for generating alternative CTPDFs which can significantly speed convergence by one to three orders of magnitude.  We have numerically tested dozens of CTPDFs with simulated radial velocity data sets to identify those which perform well for different types of orbits and suggest a set of CTPDFs for general application.  Additionally, we introduce other refinements to the MCMC algorithm for radial velocity planets, including an improved treatment of the uncertainties in the radial velocity observations, an algorithm for automatically choosing step sizes, an algorithm for automatically determining reasonable stopping times, and the use of importance sampling for including the dynamical evolution of multiple planet systems.  Together, these improvements make it practical to apply MCMC to multiple planet systems.  We demonstrate the improvements in efficiency by analyzing a variety of extrasolar planetary systems. ",
        "introduction": "Recent detections of planets around other stars have spurred a wide range of research on planet formation and planetary system evolution. The future of radial velocity planet searches promises to be exciting. Ongoing large surveys including a broad array of nearby main-sequence stars will continue to increase the number of known extrasolar planets.  The challenges of accurately determining orbital parameters will become even more important for three reasons.  First, continued monitoring with the radial velocity technique will permit the detection of planets with smaller masses.  Since the lowest mass planet will always be near the threshold of detection, they will typically have relatively low signal-to-noise ratios and hence relatively large uncertainties in model parameters.  Second, the increasing time span of precision observations will permit the discovery of planets with larger orbital periods.  Unfortunately, there can be large degeneracies in the orbital parameters for planets with orbital periods comparable to the duration of observations. Finally, the increasing precision and number of radial velocity observations is likely to reveal more multiple planet systems.  The large number of model parameters needed to model multiple planet systems can lead to degeneracies and some orbital parameters being poorly constrained.  All of these trends imply that it will become increasingly important to understand the uncertainties in orbital elements and other parameters derived from such observations. Thus, we must use the best possible statistical tools to analyze radial velocity data.   \\subsection{Introduction to Bayesian Inference}  To quantitatively analyze the available observational constraints, we employ the techniques of Bayesian inference.  The essential equations of Bayesian inference can be easily derived from the basic axioms of probability theory.  We start with a joint probability distribution, $p(x,y)$, ($x$ and $y$ may be scalars or vectors of several variables).  We then construct a marginalized probability distribution for $x$ by integrating over $y$, $p(x) = \\int p(x,y)\\, dy$.  We can then write the joint probability distribution as a product of the marginalized probability distribution and a conditional probability distribution, $p(x,y) = p(x) p(y|x)$.  Then, Bayes' theorem could be written as \\be p(y|x) = \\frac{p(x,y)}{p(x)} = \\frac{p(y) p(x|y)}{\\int p(y) p(x|y) \\, dy }. \\ee The real insight in Bayes' theorem is to identify $x$ with a set of observational data ($\\vec{d}$) and $y$ with a set of model parameters ($\\vec{\\theta}$) that are not directly observed.  By treating both the observation and the model parameters as random variables, Bayesian inference is able to address statistical questions in a mathematically rigorous fashion.  The joint probability, $p(\\vec{d}, \\vec{\\theta})$, can be expressed as the product of the likelihood ($p(\\vec{d} | \\vec{\\theta})$, the probability of the observables given the model parameters), and a prior probability distribution function ($p(\\vec{\\theta})$) which is based on previous knowledge of the model parameters.  Bayes's theorem allows one to compute a posterior probability density function, $p(\\vec{\\theta} | \\vec{d})$, which incorporates the knowledge gained by the observations $\\vec{d}$.  That is \\bea p(\\vec{\\theta}| \\vec{d}, \\mathcal{M} )  & = &\\frac{ p( \\vec{d}, \\vec{\\theta} | \\mathcal{M} ) }{ p(\\vec{d | \\mathcal{M} })} = \\frac{ p( \\vec{d}, \\vec{\\theta} | \\mathcal{M} ) }{ \\int p( \\vec{d}, \\vec{\\theta} | \\mathcal{M} ) p( \\vec{\\theta} | \\mathcal{M} ) \\,d\\vec{\\theta} } \\nonumber \\\\ & = & \\frac{ p( \\vec{\\theta} | \\mathcal{M} ) p(\\vec{d} | \\vec{\\theta}, \\mathcal{M} ) }{ \\int p( \\vec{\\theta} | \\mathcal{M} ) p( \\vec{d}| \\vec{\\theta}, \\mathcal{M} ) \\,d\\vec{\\theta} }, \\label{BayesEqn} \\eea where we have now explicitly added the fact that each of the probability distributions is conditioned on the assumption of a certain model, $\\mathcal{M}$, that includes the meaning of the model parameters, $\\vec{\\theta}$, and their relationship to the observational data, $\\vec{d}$.  Table \\ref{Tab1} provides a summary of the symbols which appear in multiple sections of this paper.  \\subsubsection{Advantages of Bayesian Inference}  Bayesian inference has several advantages over classical statistical techniques.  First, the Bayesian framework provides a rigorous mathematical foundation for making inferences about the model parameters.  Since all Bayesian inferences are based on the posterior probability distributions, there is a rigorous basis for quantifying uncertainties in model parameters, unlike frequentist methods which typically result in point estimates.  Frequentist methods are often combined with resampling techniques such as bootstrap to estimate uncertainties in model parameters.  However, these techniques rely on fictitious observations that the experimenter believes could have been made, while the Bayesian posterior probability distribution depends only on the observations that were actually made (and the prior knowledge of the model parameters).  While the posterior probability distribution is often abbreviated as simply the posterior, it is important to remember that it is a true probability distribution for the model parameters, unlike the results of resampling techniques such as bootstrap.  Because of these theoretical advantages, Bayesian inference typically produces more accurate estimates of the uncertainty in model parameters, as shown using actual radial velocity observations of extrasolar planet in Ford (2005a, Paper I).  For further discussion of alternative techniques see Paper I, \\S 3 \\& 4.  The Bayesian framework also has several practical advantages.  Because inference is based on probability distributions, it is straightforward to incorporate a variety of different types of information and observations.  Further, hierarchical Bayesian models can naturally accommodate uncertainties in models as well as in observations. Finally, the Bayesian framework provides a natural basis for making predictions about future observations.  These predictions can even be used to improve the efficiency of future experiments or observations (Loredo 2003; Ford 2005b).  The most commonly criticized aspect of Bayesian inference is the necessity of specifying prior probability distributions for the model parameters.  In many cases, the observational data provides such a strong constraint on the model parameters that the effect of the prior is negligible.  However, in cases where the observations provide only very limited constraints, the posterior distribution can be significantly influenced by the choice of prior.  A complete Bayesian analysis should include checking the sensitivity of any conclusions to the choice of prior.  It should be noted that this criticism is not unique to Bayesian inference.  Similar assumptions are being made implicitly in frequentist analyses, as the choice of model parameters can also significantly influence the results.  For a striking example of practical differences for orbit determination of binary stars, see Pourbaix (2002).  Indeed, the necessity to explicitly state what priors are being chosen can be viewed as a strength of the Bayesian method.   Given all the advantages of Bayesian inference, one might wonder why it is not the standard practice for data analysis in the physical sciences.  Unfortunately, the lower integral in Eqn.\\ \\ref{BayesEqn} can be extremely difficult to compute, particularly when $\\vec{\\theta}$ has a large number of dimensions.  The computational burden has long limited Bayesian methods to a small number of relatively simple problems.  Modern computers permit the application of Bayesian inference to an growing number of increasingly complex problem.  Even today, the brute force evaluation of the integrals is impractical for many real world problems.  Therefore, it is necessary to develop efficient algorithms for evaluating the necessary integrals.  The method of Markov chain Monte Carlo is quite general and it has become an increasingly common tool for performing the necessary Bayesian integrals in recent years (e.g., Paper I and reference therein).   \\subsection{Introduction to Markov chain Monte Carlo}  Paper I introduced Bayesian analysis for constraining the orbital parameters of extrasolar planets with radial velocity observations and presented an algorithm to perform the necessary integrations based on Markov chain Monte Carlo (MCMC) simulation.  This algorithm can accurately characterize the posterior probability distribution function for orbital parameters based on radial velocity observations.  Paper I described how to construct a Markov chain (i.e. sequence) of states (i.e. sets of parameter values, $\\vec{\\theta}_i$) which are sampled from the posterior probability function.  Such a chain can be calculated by specifying an initial set of parameter values, $\\vec{\\theta}_0$, and a transition probability, $p(\\vec{\\theta}_{i+1} | \\vec{\\theta}_i, \\mathcal{M})$.  To guarantee that the Markov chain will converge to the posterior probability distribution, the Markov chain must be aperiodic, irreducible (i.e., it must be possible for the chain to reach every state with non-zero probability from any other state with non-zero probability), and reversible, that is, \\begin{equation} p(\\vec{\\theta}| \\vec{d}, \\mathcal{M} ) p(\\vec{\\theta}|\\vec{\\theta}', \\mathcal{M}) = p(\\vec{\\theta}'| \\vec{d}, \\mathcal{M}) p(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}). \\end{equation} It is possible to construct a reversible transition probability, $p(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M} )$, from a non-reversible candidate transition probability distribution function (CTPDF), $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$, using the Metropolis-Hastings (MH) algorithm.  The MH algorithm involves the generation of a trial state ($\\vec{\\theta}'$) according to the CTPDF, $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$, and randomly accepting the trial as the next state or rejecting the trial state in favor of the current state.  The MH algorithm specifies an acceptance probability \\bea \\alpha(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) & = & \\min \\left\\{ \\frac{  p(\\vec{\\theta}'| \\vec{d}, \\mathcal{M}) q(\\vec{\\theta}|\\vec{\\theta}', \\mathcal{M})}{p(\\vec{\\theta}| \\vec{d}, \\mathcal{M}) q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) }, 1 \\right\\} \\\\ & = & \\min \\left\\{ \\frac{  p(\\vec{d} | \\vec{\\theta}', \\mathcal{M}) q(\\vec{\\theta}|\\vec{\\theta}', \\mathcal{M})}{p(\\vec{d}|\\vec{\\theta}, \\mathcal{M}) q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) }, 1 \\right\\}. \\label{AlphaEqn} \\eea When using this acceptance probability, the transition probability \\begin{equation} p(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) = q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) \\alpha(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M}) \\end{equation} is guaranteed to be reversible and irreducible, provided only that $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$ allows transitions to all $\\vec{\\theta}'$ for which $p(\\vec{\\theta}' | \\vec{d}, \\mathcal{M})$ is non-zero.  Note that the MH algorithm does not require that the normalization of $p(\\vec{\\theta}| \\vec{d}, \\mathcal{M})$ be known.  While the above algorithm guarantees that the Markov chain will converge to $p(\\vec{\\theta}| \\vec{d}, \\mathcal{M})$, it does not specify when the chain will achieve convergence.  The choice of $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$ can have a dramatic effect on the rate of convergence of the Markov chain.  Poor choices can lead to extremely inefficient sampling and hence slow convergence. The most efficient choice for $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$ would be $p(\\vec{\\theta}' | \\vec{d}, \\mathcal{M} )$, the posterior probability distribution itself.  However, this is rarely possible, since the whole purpose of the Markov chain is to calculate the posterior distribution.  Paper I presented a practical algorithm based on the Metroplis-Hastings algorithm, varying a subset of parameters at each step, and and a Gaussian CTPDF, $q(\\vec{\\theta}'|\\vec{\\theta}, \\mathcal{M})$ centered on $\\vec{\\theta}$ with a covariance matrix $I\\vec{\\beta}$, where $I$ is the identity matrix and $\\vec{\\beta}$ is a vector of scale parameters, $\\beta_\\nu$.  Throughout this paper, we use the index $\\nu$ to distinguish the model parameters.  The Gibbs sampler specifies that only a subset of $\\vec{\\theta}$ is altered at each step of the Markov chain.  While there are several variations, we choose which parameters are to be altered in the next trial state according to randomly generated permutations of the model parameters. At each step one function of the model parameters ($u_\\mu(\\vec{\\theta})$) is chosen to be updated using, \\begin{equation} q(u_\\mu(\\vec{\\theta}') | u_\\mu(\\vec{\\theta}), \\mathcal{M} ) = \\frac{1}{\\sqrt{2\\pi \\beta_\\mu^2}} \\exp \\left[ -\\frac{\\left[u_\\mu(\\vec{\\theta}')-u_\\mu(\\vec{\\theta})\\right]^2}{2\\beta_\\mu^2} \\right] \\label{eqnCandTransProb} \\end{equation} for valid $\\vec{\\theta}'$ (i.e., if the model dictates that $\\vec{\\theta}'_\\nu$ be positive definite, then trial states with negative $\\vec{\\theta}'_{\\nu}$ are rejected).  We use the index $\\mu$ to distinguish the different types of steps, and the index $i$ to indicate the $n$-th step of the Markov chain.  Each $\\beta_\\mu$ is a parameter which controls the scale for the steps based on the quantity indicated by $\\mu$.  In Paper I, we used $\\vec{u}(\\vec{\\theta}) = \\vec{\\theta}$, i.e., we only took steps in the model parameters.  In \\S 4 of this paper, we will present several alternative CTPDFs that can significantly improve the computational efficiency of the MCMC algorithm by reducing the number of steps required before a Markov chain can be used for inference.  \\subsection{Summary of Previous Results}  The techniques of Bayesian inference and Markov chain Monte Carlo have previously been applied to analyzing the radial velocity observations of several extrasolar planetary systems.  Originally, Paper I presented a complete algorithm for calculating the posterior probability distributions for orbital parameters based on radial velocity observations.  Paper I demonstrated several shortcomings of the conventional estimates of parameter uncertainties based on bootstrap-type resampling.  It found that Bayesian analyses, were particularly important for planets where the orbital period is comparable to the duration of observations.  Driscol (2006) is performing a systematic study to compare uncertainty estimates made with bootstrap-style resampling and uncertainty estimates made with MCMC. Gregory (2005a) used Bayesian inference and MCMC to reanalyze the observations of one particular system (HD 73526), and found two alternative orbital solutions, one of which had a larger posterior probability than the originally published orbital solution.  This demonstrated another shortcoming of frequentist methods that are based on only the maximum likelihood solution, rather than the posterior probability distribution.  Ford, Lystad \\& Rasio (2005) reanalyzed the observations of the three planets orbiting $\\upsilon$ Andromedae to derive improved constraints on the orbital parameters and perform a dynamical analysis (Ford et al.\\ 2005).  Gregory (2005b) used an improved algorithm to analyze another system (HD 208487) and claim that the present data show a $\\sim95\\%$ probability for a second planet.  \\subsection{Motivation for this work}  Each of the above studies using MCMC has demonstrated advantages of Bayesian inference for analyzing radial velocity observations of extrasolar planets.  However, significant obstacles remain.  For several of the systems analyzed in Paper I, the model parameters were highly correlated and caused the Markov chains to converge slowly. Convergence was particularly problematic for multiple planet systems, so Paper I included an analysis of only one multiple planet system (GJ 876).  While we had attempted to include other multiple planet systems, slow convergence rates and computational limitations prevented us from being sufficiently confident that the other Markov chains had converged to include them in the final paper.  Gregory (2005a) took the more risky approach of basing inferences on Markov chains which showed obvious signs of non-convergence, but argued that the main conclusions were robust.  These studies illustrate the importance of improving the computational efficiency of MCMC algorithms for analyzing radial velocity data.  Since Paper I, we have dramatically improved the computational efficiency of our MCMC algorithm by introducing alternative CTPDFs. Our improved CTPDFs allow for the very rapid analysis of the typical single planet systems, as well as the practical application of Bayesian inference and MCMC to multiple planet systems.  For example, incorporating just a few of these optimizations permitted the dynamical analysis of the $\\upsilon$ Andromedae system (Ford et al.\\ 2005).   \\subsection{Outline}  In this paper we describe several refinements to the algorithms presented in Paper I. In \\S 2 we present our physical model of the planetary system and observations.  We include modifications to the model that allow the MCMC technique to be applied to systems with a significant amount of stellar ``jitter'' and/or additional unknown planetary companions.  We also present a method for testing the sensitivity of posterior distribution functions to the assumption that the observational uncertainties are normally distributed. In \\S 3 we provide precise descriptions of a few important technical issues, including choosing priors, choosing step sizes, testing for non-convergence, and deciding when to stop calculating Markov chains. In \\S4 we describe modifications to the original MCMC algorithm presented in Paper I.  In particular, we present several alternative choices for the CTPDF, $q(\\vec{\\theta}'|\\vec{\\theta})$, and demonstrate the improved efficiency by applying the new sampling algorithms to several different types of simulated planetary systems. Of particular interest, we make practical suggestions for the choices of CTPDFs in \\S\\ref{SRec}. In \\S 5, we discuss the application to multiple planet systems. In \\S 6, we demonstrate our improved algorithms on a few examples of actual planetary systems. In \\S7, we summarize our conclusions and discuss areas for future research. ",
        "conclusions": "This and previous papers have demonstrated numerous advantages using Bayesian inference for analyzing the orbital parameters of extrasolar planets.  Unfortunately, Bayesian analyses typically require calculating multidimensional integrals which can be computationally challenging.  In Paper I, we introduced the method of MCMC, the Gibbs sampler, and the MH algorithm for performing these integrals.  This made it practical to use MCMC to perform Bayesian analyses of the orbital parameters of several extrasolar planets.  In this paper, we developed and tested several alternative CTPDFs to further improve the efficiency of MCMC.  In particular, we recommend two sets of CTPDFs which together can accelerate the convergence of Markov chains by multiple orders of magnitude (see \\S\\ref{SRec}).  These improvements make it practical to analyze all the known extrasolar planets and even multiple planet systems.  We demonstrate our optimized MCMC algorithms by analyzing several actual extrasolar planet systems.  We anticipate applying these algorithms to a large number of extrasolar planet systems to obtain more accurate orbital parameters and improve dynamical investigations of multiple planet systems.  Several significant challenges still prevent fully Bayesian analyses from being routinely applied to extrasolar planet observations.  In particular, it can be quite computationally demanding to sample from the posterior distribution when the posterior distribution has several well separated modes.  This commonly occurs for low-mass planets that are not yet clearly detected.  It would be desirable to be able to analyze such planets, since the marginally detected planets can influence the orbital parameters derived for other planets that have already been clearly detected around the same star.  Another important challenge is developing computationally efficient means of performing Bayesian model selection, i.e., simultaneously considering models with zero, one, two, or more planets.  This would allow planet detections to be based on Bayesian rather than maximum likelihood analyses. Gregory (2005a) has suggested parallel tempering as one possible approach, however even more efficient algorithms would be desirable."
    },
    "0512/astro-ph0512128_arXiv.txt": {
        "abstract": "``The Duck'' is a complicated non-thermal radio system, consisting of the energetic radio pulsar \\psr, its surrounding pulsar wind nebula \\pwn\\ and the adjacent supernova remnant (SNR) \\snr.  PSR~\\psr\\ was originally claimed to be a young ($\\approx15\\,000$~yr) and extreme velocity  ($\\ga1500$~\\kms) pulsar which had penetrated and emerged from the shell of the associated SNR~\\snr, but recent upper limits on the pulsar's motion have raised serious difficulties with this interpretation.  We here present 8.5~GHz interferometric observations of the nebula \\pwn\\ over a 12-year baseline, doubling the time-span of previous measurements. These data correspondingly allow us to halve the previous upper limit on the nebula's westward motion to 14~milliarcseconds yr$^{-1}$ (5-$\\sigma$), allowing a substantive reevaluation of this puzzling object.  We rule out the possibility that the pulsar and SNR were formed from a common supernova explosion $\\approx15\\,000$~yrs ago as implied by the pulsar's characteristic age, but conclude that an old ($\\ga70\\,000$~yr) pulsar / SNR association, or a situation in which the pulsar and SNR are physically unrelated, are both still viable explanations. ",
        "introduction": "The formation and subsequent evolution of pulsars are not yet fully understood. A powerful constraint on these processes is provided by an independent age estimate. In cases where both a pulsar and its associated supernova remnant (SNR) can be identified, an age estimate which is independent of both distance and inclination effects is simply $t_p = \\Theta / \\mu$, where $\\Theta$ is the angular separation between the pulsar and its inferred birth site, and $\\mu$ is the pulsar's proper motion \\citep{mgb+02,klh+03}. This can then be compared to the age of the system as expected from spin-down \\citep[e.g.,][]{lk05}: \\begin{equation} t_p=\\frac{P}{(n-1)\\dot{P}}\\left[1-\\left(\\frac{P_0}{P}\\right)^{n-1}\\right] , \\label{eqn_age} \\end{equation} where $P$ is the current spin period, $\\dot{P}$ is the time-derivative of $P$, $P_0$ is the period at birth, and $n$ is the ``braking index'' (see further discussion in \\S\\ref{sec_old}). Comparison of these two independent age estimates  provides information on $P_0$ and $n$, i.e., on the processes which impart and then dissipate the considerable angular momentum of neutron stars (e.g., Gaensler \\& Frail 2000, hereafter GF00; Kaspi \\etal\\ 2001b\\nocite{gf00,krv+01}). If one assumes that $P_0\\ll P$ and $n=3$, Equation~(\\ref{eqn_age}) reduces to the expression for the ``characteristic age'' of a pulsar, $\\tau_c \\equiv P/2\\dot{P}$.  PSR~\\psr\\ is an isolated 125-ms pulsar, surrounded by the cometary radio and X-ray pulsar wind nebula (PWN) \\pwn, which in turn is located just outside the SNR~\\snr\\ \\citep{ckk+87,mkj+91,fk91,fkw94,kggl01}. As shown in Figure~\\ref{fig_snr}, the combined system has a distinctive morphology which has led to it being termed ``the Duck''. Because of the proximity of the pulsar to the SNR, and because the PWN morphology suggests that the pulsar is moving away from the SNR interior, it has been widely assumed that the pulsar and the SNR are physically associated.  The SNR and PWN also have very similar \\HI\\ absorption spectra, both suggesting a distance of $\\sim5$~kpc \\citep{fkw94}, and consistent with the distance of 5.1~kpc implied by the pulsar's dispersion measure and the Galactic electron density model of \\cite{cl02}.  We consequently adopt a common distance of 5~kpc for pulsar, PWN and SNR in further discussion.  PSR~\\psr\\ has a characteristic age $\\tau_c = 15$~kyr \\citep{mhth05}. If we assume $t_p = \\tau_c$ and that the pulsar has traveled $\\Theta = 16\\farcm1 -20\\farcm6$ from the SNR's center in its lifetime \\citep{fkw94},\\footnote{The range of values quoted for $\\Theta$ reflects the uncertainty in locating the pulsar's presumed birthplace. This results from the fact that the pulsar's inferred trajectory does not pass through the SNR's geometric center \\citep[Fig.~\\ref{fig_snr};][]{fkw94}. See \\S\\ref{sec_old} for further discussion.} this implies a westward proper motion for the pulsar of magnitude $\\mu=\\Theta/\\tau_c = 63-80$~mas~yr$^{-1}$.  However, radio interferometric observations of the western tip of the PWN taken at the Very Large Array (VLA) over 6.7~years yielded a surprising 5-$\\sigma$ upper limit $\\mu< 25$~mas~yr$^{-1}$ (GF00\\nocite{gf00}), implying a pulsar age $t_p > 39-50$~kyr~$\\gg\\tau_c$, if the pulsar was born near the SNR's geometric center.  GF00 used the stand-off distance of the PWN, the radius of the SNR and the separation of the pulsar from the SNR's center to derive a solution for the system's evolution which predicted $t_p \\sim 90- 170$~kyr and $\\mu \\sim 10$~mas~yr$^{-1}$.  Subsequently, Thorsett \\etal\\ (2002, hereafter TBG02\\nocite{tbg02}) observed the pulsar itself with the VLA over a 3.9-year baseline, and placed an independent upper limit on westward proper motion\\footnote{TBG02 state that their limit is $\\mu < 16$~mas~yr$^{-1}$ at 95\\% confidence (i.e., $2\\sigma$).  Here we use 5-$\\sigma$ limits throughout, and have adjusted their limit accordingly.} of $\\mu < 37$~mas~yr$^{-1}$.  TBG02 argue that this is most easily explained if PSR~\\psr\\ and SNR~\\snr\\ are unrelated, and if the pulsar is instead moving away from the center of the PWN~\\pwn. Assuming $t_p \\approx \\tau_c$, one can then predict a proper motion $\\mu \\sim 5$~mas~yr$^{-1}$.  As a further alternative, \\cite{gva04b} has argued that this system results from a massive high-velocity progenitor star which went supernova inside its wind-blown bubble. The pulsar began its life substantially offset from the cavity's center, but the resulting SNR expands to take on the shape of the cavity.  \\cite{gva04b} subsequently develops a model in which $t_p \\approx 54$~kyr, $\\Theta \\approx 6'$ and $\\mu \\approx 7$~mas~yr$^{-1}$.  As an attempt to resolve this puzzling situation, we have conducted a new observation of the western cometary tip of PWN~\\pwn\\ (see inset to Fig.~\\ref{fig_snr}), which doubles the time baseline considered by GF00 to 12 years.  In \\S\\ref{sec_obs} we present our new observations and corresponding measurement of proper motion, while in \\S\\ref{sec_disc} we we interpret these new results in the context of various possibilities proposed for the origin and evolution of this system. ",
        "conclusions": "The upper limit on proper motion we find for the PWN~\\pwn\\ is the most restrictive one obtained to date for this system.  We are able to reject the original expectation that PSR~\\psr\\ was born $\\sim15$~kyr ago and is associated with SNR~\\snr, regardless of whether the corresponding supernova occurred at the geometric center of the SNR, or at a site substantially offset from this.  Two possibilities remain. First, the pulsar and SNR are associated and share an age $\\ga70$~kyr. In this case, the pulsar has a typical projected space velocity of $\\la330$~\\kms, which allows it to overtake its SNR at this stage and drive a bow shock through ambient gas. The sudden drop in pressure as the pulsar crosses the SNR's radiative shell might also explain the bulbous component of the PWN~\\pwn\\ seen between the pulsar and the SNR.  The pulsar's passage has re-energized the radio emission from the SNR through rapid diffusion of pulsar wind particles along magnetic field lines; without this new injection of particles, the SNR would be much dimmer, would have a steeper spectrum and generally would be more difficult to detect.  The slight offset between the pulsar's inferred trajectory and the SNR's geometric center possibly results from a density gradient into which the SNR is expanding, or from a slightly offset explosion within a pre-existing cavity. The only difficulty with this picture is that the system's age must then be many times larger than the pulsar's characteristic age, which is not seen in other pulsar / SNR associations. The proposed explanation is that the surface magnetic field of the pulsar is at the present epoch growing on a time scale of $\\sim15$~kyr. This effect is consistent with the properties of other pulsar / SNR associations, and is possibly also being seen in the long-term timing signatures of several pulsars \\citep{lpgc96,lyn04}.  The alternative is that the pulsar is $\\sim15$~kyr old as indicated by its characteristic age, and is unassociated with SNR~\\snr.  This explanation eliminates the need to invoke off-center cavity explosions, re-energized shells or growing magnetic fields to explain the observations. However, this interpretation offers no easy explanation for the morphology of \\pwn\\ which, contrary to the proposal of TBG02, cannot be easily explained as a relic nebula left behind at the pulsar's birth site. This model also requires that the brightened emission and flat spectrum of the SNR near the pulsar be a coincidence.  Frustratingly, a full understanding of the Duck remains elusive. However, future VLA or VLBA observations of PSR~\\psr\\ should eventually detect motion. Furthermore, an X-ray detection of \\snr\\ and of the eastern parts of \\pwn\\ should provide additional constraints on the properties  of the system \\citep[see discussion by][]{kggl01}, while continued timing of the pulsar may be able to indicate the nature of the braking torque and magnetic field evolution in this source \\citep[e.g.,][]{lpgc96}.  Finally, other similar systems might provide additional clues.  PSR~B1951+32 is clearly in the process of puncturing the shell of the SNR~CTB~80 \\citep{hk88,fss88}, while PSR~J1016--5857 and SNR G284.3--1.8 have been proposed as another such interacting system \\citep{cbm+01}. Over the  ensemble of Galactic pulsars and SNRs, \\cite{sfs89} predict approximately half a dozen systems in which the pulsar has recently penetrated the SNR shell.  Searches for further interacting pairs may thus prove fruitful."
    },
    "0512/astro-ph0512402_arXiv.txt": {
        "abstract": "% We study axisymmetric models of layered protoplanetary discs taking radiative transfer effects into account, and allowing for a residual viscosity in the dead zone. We also explore the effect of different viscosity prescriptions. In addition to the ring instability reported in the first paper of the series we find an oscillatory instability of the dead zone, accompanied by variations of the accretion rate onto the central star. We provide a simplified analytical description explaining the mechanism of the oscillations. Finally, we find that the residual viscosity enables stationary accretion in large regions of layered discs. Based on results obtained with the help of a simple 1-D hydrocode we identify these regions, and discuss conditions in which layered discs can give rise to FU~Orionis phenomena. ",
        "introduction": "%  The transport of angular momentum in protoplanetary discs is most probably driven by the magnetorotational instability, described by \\citet{bh91a,bh91b}, and hereafter referred to as MRI. Since the protoplanetary discs are ionized rather weakly, in some regions of the disc the ionization degree $\\xi$ can decrease below a critical level, $\\xi_c$, at which the gas decouples from the magnetic field and the MRI decays. Such region is referred to as a {\\em dead zone}, contrary to the so called {\\em active zone} where the MRI operates.  The most important processes responsible for the ionization of gas in protoplanetary discs are 1) particle collisions due to thermal motions, and 2) irradiation by cosmic rays and high-energy photons originating from the central star. The collisional ionization dominates in the innermost part of the disc, where the temperature is high enough to maintain $\\xi>\\xi_c$ at all distances from the disc plane. Further away from the star $\\xi_c$ is exceeded only in two surface layers (on both sides of the disc), each ionized by cosmic rays and stellar X-ray quanta, and having approximately constant column density $\\Sigma_\\mathrm{a}$. Both processes result in the following distribution of disc activity: at $r < r_1$ (where the thermal ionization degree exceeds $ \\xi_c$) the whole disc is active; at $r_1<r<r_2$ (where $\\Sigma>\\Sigma_\\mathrm{a}$) a dead zone is sandwiched between two active layers; and at $r > r_2$ (where $\\Sigma< \\Sigma_\\mathrm{a}$) the whole disc becomes active again. This model was proposed by \\citet{gammie96} to explain the FU~Ori outbursts.  FU~Ori stars are young stellar objects whose optical brightness occasionally increases by $3-4$~mag on a time-scale of $1-10$~yr, and remains at the increased level for $10-100$~yr. It is generally believed that these events are related to the variations of the accretion rate in circumstellar discs. To explain the observed increase in luminosities the accretion rate has to change from $\\sim 10^{-7}$ to $\\sim 10^{-4}$~M$_\\odot$yr$^{-1}$, implying that as much as $10^{-2}$M$_\\odot$ must be accreted onto the central star during the outburst \\citep{hk96}.  The FU~Ori phenomenon may be caused by a thermal instability of the protoplanetary disc (see e.g. \\citet{lc04}, \\citet{bl94} and references therein), analogous to the one operating in accretion discs around primary components of cataclysmic binaries \\citep{mm-h81}. However, the thermal instability model has difficulties with explaining long duration of the outbursts, since the high accretion rate during the outburst quickly removes most of the material from the inner part of the disc where the instability occurs \\citep{hk96}. A transition of the very inner disc to the outburst state due to the thermal instability was studied by \\citet{kl99} with the help of the similar two-dimensional code as used in this work.  \\citet{gammie96} suggested that this problem may be solved by the layered model, in which the dead zone serves as a mass reservoir for the outburst. In the layered disc the accretion rate can be an increasing function of $r$. When this is the case, an annulus centred at $r_0$ receives more mass per unit time from the disc exterior to it $(r>r_0$ than it loses to the disc interior to it $(r<r_0)$, causing the matter to accumulate in the dead zone. If this process lasts long enough, the total surface density of the disc may become so large that once the accretion was somehow triggered in the dead zone it would release enough heat to restore the coupling between the matter and the magnetic field. Such an \"ignition\" of the accumulated matter would then start a self-sustaining outflow of mas from the dead zone into the inner disc, substantially increasing the accretion rate.  \\citet{alp01} assume that the triggering factor is the gravitational instability. Their model shows repeating periods of vigorous accretion separated by quiescent intervals lasting $\\sim 10^5$~yr, in qualitative agreement with the observed properties of FU~Ori objects. However, our radiation-hydrodynamic simulations of layered discs (\\citealt*{wkr05}, hereafter Paper~I) indicate, that it may be rather difficult for the dead zone to store $\\sim10^{-2}$M$_\\odot$ of the disc matter necessary to feed an FU~Ori outburst. The reason is the ring instability described in Paper~I, which, long before the required mass is reached, causes the dead zone to split into well-defined rings that become unstable to the instability discovered by \\citet{pp85}.  In the present paper we extend the simulations reported in Paper~I onto a broader range of disc parameters, with the principal aim to find in which conditions the dead zone could become massive enough to account for FU~Ori outbursts. For reasons explained in \\S2, along with models whose dead zones are entirely free of viscosity we explore models with some residual viscosity acting in the region which formally should be declared dead. We also explore the effect of different viscosity prescriptions.  The paper is organized as follows. In \\S2 we briefly describe the numerical methods and input physics. The results of simulations are reported in \\S3. In \\S4 we give a simple analytical description of processes at the inner boundary of the dead zone, and derive conditions for ignition. In \\S5 we summarize our results and discuss the limits of the mass accumulation process. ",
        "conclusions": "%  The two-dimensional radiation-hydrodynamic simulations of the layered disc we carried out showed that the behaviour of the disc strongly depends on the presence of a small residual viscosity in the dead zone $\\alpha_\\mathrm{DZ}$. In the models with $\\alpha_\\mathrm{DZ} = 0$ the ring instability develops as described in Paper~I. The  rings may eventually decay which would temporarily increase the accretion rate onto the central star by a factor of several. On the other hand, in the models with $\\alpha_\\mathrm{DZ}$ set to 10\\% of the viscosity in the active parts of the disc, the dead zone remains contiguous, but its inner boundary oscillates.  These oscillations are more or less regular with period between several tens and several hundreds years depending on model parameters. They can be accompanied by variations of the accretion rate on a time-scale as short as ten years, at a relative amplitude reaching from a few per cent to 200 -- 300 per cent in the high viscosity Model 8. These variations would be rather difficult to observe, however it does not seem absolutely impossible. Since the properties of these oscillations are quite sensitive to physical conditions in the disc, such observations could help to constrain the set of models of protoplanetary discs.  Similar variations of the accretion rate were also found by \\citet{stepinski99} using an one-dimensional hydrodynamic code. However, the period was much longer than in our model ($10^3-10^4$~yr). The reason for that may be that in his models the ignition of the inner part of the dead zone is much more difficult due to the absence of the radiation transport (and, generally speaking, due to oversimplified description of the interactions between the inner hot $\\alpha$-disc and the layered region).  We give an analytical description of these oscillations based on the vertically averaged model modified for the case of layered disc. It illustrates the mechanism of the dead zone oscillations and tests the numerical model.  Finally, we perform a simple 1-D model of the layered disc based on the analytical solution to follow the evolution of the whole layered region for a long time. We found that with $\\alpha_\\mathrm{DZ} \\neq 0$ a stationary layered region may exist at higher radii (this was also observed in numerical Model~8). Part of this stationary region is still combustible, i.e. it may be made active by increasing the temperature by some external event, and part of its mass is available for accretion.  We compare the amount of mass stored in the combustible region to the mass accreted during a FU~Ori type outburst following the idea that the dead zone can serve as a mass reservoir for such event. We found, that these masses are comparable for the case of massive disc with the high accretion rate, which is in agreement with the observation that the FU~Ori objects are very young systems."
    },
    "0512/astro-ph0512364_arXiv.txt": {
        "abstract": "We have modelled the process of reionization of the intergalactic medium (IGM) by photoionization by galaxies, in order to learn what galaxy formation in the framework of the CDM model predicts for the epoch of reionization.  We use a sophisticated semi-analytic model of galaxy formation to track the formation of these galaxies, their influence on the IGM, and the back-reaction of the state of the IGM on further galaxy formation. Our study represents a much more complete and physically consistent modelling of reionization than has been conducted in the past. In particular, compared to previous work by ourselves and others, our new calculations contain significant improvements in the modelling of the effects of reionization of the IGM on the collapse of baryons into dark matter halos (this is now computed self-consistently from the properties of model galaxies), and in the model for the cooling and condensation of gas within halos (our new model includes photoheating from a self-consistently computed ionizing background and also includes cooling due to molecular hydrogen).  We find that reionization can be achieved by $z\\sim 10$--$20$ in a $\\Lambda$CDM cosmological model with $\\sigma_8\\approx 0.9$.  However, a cosmological model with a running spectral index is only able to achieve reionization before $z\\approx 9$, and thus be consistent with an optical depth of $0.1$, if very extreme assumptions are made about the physics of feedback at high redshifts.  We also consider the specific galaxy formation model recently discussed by Baugh et al., which includes a top-heavy IMF in starbursts, and find that it is able to reionize the Universe by $z\\approx 12$. The previous results assume that all of the ionizing photons produced by stars in galaxies are able to escape and ionize the IGM. If this is not the case, then the redshift of reionization could be substantially reduced. We find that extended periods of partial reionization and double reionizations can occur in models in which the first stars formed via cooling by \\H2 molecules, are very massive, and in which the escape fraction of ionizing photons $\\sim $10--30\\%.  Such models do not fully reionize until $z\\approx$ 6--7, but predict an electron scattering optical depth as large as $0.15$. Models with lower $\\sigma_8=0.7--0.8$ as suggested by the recent Wilkinson Microwave Anisotropy Probe (WMAP) three year data have reduced redshifts of reionization, but can be consistent with the lower optical depth also suggested by the WMAP three year data. ",
        "introduction": "While gas in the intergalactic medium (IGM) is highly ionized in the present day Universe, it is known to have been mostly neutral in the past. Evidence for this comes from observations of the Gunn-Peterson effect in the spectra of high redshift quasars \\cite{becker01,djorg01,fan06} and from the value of the optical depth to scattering measured from the cosmic microwave background (CMB; \\pcite{spergel03}). Thus, there was a period in the history of the Universe when it transitioned from being neutral to ionized---the epoch of reionization.  Measurements of the CMB polarization made during the first year of operation of the Wilkinson Microwave Anisotropy Probe (WMAP) satellite\\footnote{While this paper was being completed, results based upon the first three years of data from WMAP became available (Spergel et al. 2006). These show a much lower optical depth to reionization. We will comment upon these results in \\S\\protect\\ref{sec:disc}.} have been used to estimate a value $\\tau=0.17 \\pm 0.04$ for the electron scattering optical depth to reionization \\cite{kogut03}, which, in the simple case that the IGM reionized instantaneously, implies a reionization redshift $z_{\\rm reion}=17\\pm 5$ \\cite{spergel03}.  This reionization redshift is significantly higher than was previously thought on the basis of the observed Gunn-Peterson effect and on the basis of theoretical modelling (see, for example, \\pcite{haiman97,go97,baugh98,vs99,bruscoli00,chiu00,ciardi00,benson01}).  Although the WMAP measurement could be explained in scenarios in which reionization happened twice, or happened only partially for some time \\cite{cen03,wylo03,nasel04},\\footnote{The physical plausibility of these scenarios has been questioned by \\protect\\scite{furl05}. We present evidence, in \\S\\protect\\ref{sec:enhance} that such double or partial reionization may in fact occur under plausible physical conditions.} the simple interpretation (that the Universe was fully reionized at $z=17$ and remained so ever since) implies the presence of a significant population of ionizing sources at high redshift. Coupled with the hints that the effective spectral index of the matter power spectrum is changing with scale (i.e. a ``running spectral index''), this scenario becomes hard to explain in the standard cold dark matter (CDM) scenario of structure and galaxy formation.\\footnote{The running spectral index reduces the amplitude of fluctuations on small, galactic scales relative to that on large scales. As a result there are many fewer haloes present at high redshift with mass sufficient to form galaxies.}  This tension has been explored recently by several authors \\cite{hh03,rs03,yoshi03b,ciardi03,onken04}, who generally reach the conclusion that $z_{\\rm reion}=17$ is incompatible with the running spectral index model, and hard to achieve even in a power-law model. One possible escape from this apparent problem is to reionize the Universe in some other way, i.e. without relying on ionizing photons from stars in high redshift galaxies. One such mechanism which has received attention recently is a contribution to reionization from particle decays during the cosmic dark ages \\cite{chen04,hansen04,kasuya04,kasuya04b}, another is a contribution from ``mini-QSOs'' (see, for example, \\pcite{rico04}).  In this paper, we undertake detailed calculations of the galaxy formation process at high redshift to investigate whether galaxies alone can ionize the Universe sufficiently early to explain the WMAP measurement of the optical depth. For this purpose we employ the {\\sc galform} semi-analytic model of galaxy formation \\cite{cole00,benson02}. The advantage of our model over previous attempts to address this question is that it incorporates much more of the actual physics of galaxy formation (e.g. it follows the cooling of gas within haloes and incorporates a model of star formation and feedback), even if this physics is treated in a simplified way. The model used in this paper is significantly improved over that employed in our earlier study of reionization in \\scite{benson01}. The effects of reionization on subsequent galaxy formation (due to the heating of the IGM and the subsequent rise in the Jeans mass, and due to photoheating of gas in haloes) are now self-consistently computed within our model. Furthermore, we now include a detailed treatment of cooling due to molecular hydrogen in our calculations. Thus, we are able to provide a more reliable answer as to what the CDM model of structure formation predicts for the epoch of reionization.  The remainder of this paper is arranged as follows. In \\S\\ref{sec:model} we outline the model used to follow the reionization of the Universe. In \\S\\ref{sec:results} we examine the reionization histories predicted by our models and assess the importance of the various physical ingredients of the models.  In \\S\\ref{sec:2res} we present results for the reionization epoch and the imprint of reionization on the cosmic microwave background. Finally, in \\S\\ref{sec:disc} we assess the implications of our results. Several Appendices describe further details of the modelling, and present additional results. ",
        "conclusions": "\\end{figure*}  \\begin{figure*} \\addtocounter{figure}{-1} \\caption{\\emph{(cont.)} The evolution of key physical parameters in our models. The three columns show mean star formation rate density, ionizing luminosity and \\HII\\ filling factor as indicated at the top of each panel. Each row corresponds to a different galaxy formation model, as indicated on the right-hand side of the figure. The three lines in each panel correspond to three different cosmological models: WMAP RSI (red), WMAP PL (blue) and Baugh et al. (2005; green---not shown for models MetFree, TopHeavy, WeakFB and H$_2$WeakFB-VMS). In the \\HII\\ filling factor columns, the horizontal dashed line indicates the point of complete reionization, while dashed coloured lines for the H$_2$WeakFB-VMS model show the contribution of Pop~I and II stars to the total filling factor.} \\end{figure*}   \\subsection{Does H{\\boldmath $_2$} Cooling Help?} \\label{sec:h2help}  We can examine whether the inclusion of \\H2\\ cooling channels in our calculations actually aids in pushing the epoch of reionization to higher redshifts. Figure~\\ref{fig:FfillH2} shows the \\HII\\ filling factor as a function of redshift for models with and without cooling due to \\H2. It can clearly be seen that the addition of \\H2\\ cooling channels has almost no effect on the epoch of reionization. At earlier redshifts, \\H2\\ cooling \\emph{does} result in an enhancement in the filling factor (e.g. by a factor of $10^4$ at $z=35$ in the WMAP PL cosmology). However, by the time $F_{\\rm fill}$ is approaching unity, the ionizing emissivity is dominated by emission from haloes which are able to cool through atomic processes, and so the addition of \\H2\\ cooling channels does not help to increase $z_{\\rm reion}$ significantly.  \\begin{figure} \\psfig{file=Ffill_H2.ps,width=80mm} \\caption{The \\HII\\ filling factor as a function of redshift for models with (dot-short dashed lines) and without (solid lines) \\H2\\ cooling. Colours indicate the cosmological model as indicated in the figure. Results are shown to higher redshifts and lower filling factors than in Fig.~\\protect\\ref{fig:summary}.} \\label{fig:FfillH2} \\end{figure}  To understand the effects of cooling through molecular hydrogen, we can explore the cooled baryonic fraction in haloes as a function of their mass. The upper panels of Fig.~\\ref{fig:fcond_nophoto} show the condensed baryonic fractions of haloes at redshifts 5 (left-hand panel) and 10 (right-hand panel) for models without the effects of photo-ionization and photo-dissociation included. Blue circles indicate the cooled fraction, $f_{\\rm cond}$ (i.e. the mass of material which has been able to cool and condense into the galactic phase), for a model with no molecular hydrogen cooling, while red circles show results for a model with molecular hydrogen cooling. We indicate, with a vertical dashed magenta line, the halo mass for which the virial temperature corresponds to the supernova feedback parameter $V_{\\rm hot}$ in our semi-analytic model. Vertical, dashed cyan lines indicate the filtering masses for the two models. In the lower panels of Fig.~\\ref{fig:fcond_nophoto} are curves which show the cooling measure $t_{\\rm age}/(1+\\beta)t_{\\rm cool}$, the ratio of the age of the Universe at the specified redshift $t_{\\rm age}$ to the cooling time $t_{\\rm cool}$ (calculated at 200 times the mean density of the Universe, which is approximately the mean density of haloes), and divided by $1+\\beta$, where $\\beta$ (given by eqn.~\\ref{eq:SNfeedback}) is the ratio of the mass ejection rate due to supernova feedback to the star formation rate in model galaxies. Solid blue lines are for the model without molecular hydrogen cooling and without Compton cooling, while dashed red lines include molecular hydrogen cooling (and Compton cooling). We show curves for primordial gas and for gas of half Solar metallicity (the upper curves at large halo masses).  Our cooling measure, the quantity $t_{\\rm age}/(1+\\beta)t_{\\rm cool}$, is interesting since it roughly determines whether a halo has had time to turn most of its supply of gas into stars. Roughly speaking, a galaxy will accrete (from its surrounding hot atmosphere) all available gas in a time $t_{\\rm cool}$. Assuming haloes to have survived for a period roughly equal to $t_{\\rm age}$, the ratio of these two timescales then determines whether or not a halo has been able to accrete all available gas. However, much of this gas will be ejected once again due to feedback. This gas may be subsequently re-accreted. Consequently, a galaxy must accrete its hot halo $\\sim(1+\\beta)$ times before it can turn the majority of the available gas into stars.\\footnote{Assuming the cooling/reheating loop to occur $n$ times (and assuming cooling and reheating to happen instantaneously), the fraction of available gas eventually bound into stars will be $1-(\\beta/(1+\\beta))^n$. For $n=1+\\beta$ this converges, for large $\\beta$ to $1-1/{\\rm e}\\approx 0.63$.}  When $t_{\\rm age}/(1+\\beta)t_{\\rm cool}$ is greater than one (the point $t_{\\rm age}/(1+\\beta)t_{\\rm cool}=1$ is indicate by green horizontal lines in Fig.~\\ref{fig:fcond_nophoto}), the effective cooling time is short compared to the age of the Universe and we do not expect the cooling function to have strongly determined the shape of the $f_{\\rm cond}$ curve. When it is less than unity, the effective cooling time is long and so the shape of the $f_{\\rm cond}$ curve should have been determined strongly by the shape of the cooling curve.  \\begin{figure*} \\psfig{file=color_fcond_nophoto.ps,width=160mm,bbllx=0mm,bblly=75mm,bburx=190mm,bbury=270mm,clip=} \\vspace{-15mm} \\caption{Measures of the ability of baryonic material to cool into the galactic phase are shown for $z=5$ (left-hand panels) and $z=10$ (right-hand panels) for models \\emph{without} the effects of photo-ionization and photo-dissociation feedback. Vertical dashed magenta lines indicate the halo mass for which the virial temperature corresponds to the feedback parameter $V_{\\rm hot}$. Vertical, dashed cyan lines indicate the filtering masses for the two models. \\emph{Upper panels:} The average fraction of baryonic material in haloes which has been able to cool and condense into galaxies, as a function of halo mass in the Standard model in the WMAP RSI cosmology. Blue circles indicate the cooled fraction for a model with no molecular hydrogen cooling, while red circles show results for a model with molecular hydrogen cooling. \\emph{Lower panels:} Cooling measures for metallicities $Z=0$ and $Z=0.5 Z_{\\odot}$ (the upper curves at large halo masses): solid blue lines are for the model without molecular hydrogen cooling, while dashed red lines include molecular hydrogen cooling (and Compton cooling). The cooling measure is the ratio $t_{\\rm age}/(1+\\beta)t_{\\rm cool}$. Here $\\beta$ is the ratio of rate of mass ejection due to supernova feedback to the star formation rate in model galaxies. The horizontal green lines indicate the point where the cooling measure equals 1 (i.e. the point above which cooling is efficient in the sense described in the text.} \\label{fig:fcond_nophoto} \\end{figure*}  From Fig.~\\ref{fig:fcond_nophoto} we can see that the inclusion of \\H2\\ cooling makes no difference to the fraction of condensed mass in haloes with virial temperatures above about $2 \\times 10^4$K. \\H2\\ cooling does however, result in a tail of low mass haloes (with virial temperatures below $10^4$K) which are able to condense some of their gas, unlike in the case of cooling through purely atomic processes. However, the amount of gas condensed is minimal, with $\\lsim 10$\\% of the total available gas being condensed (a number which falls rapidly to zero in lower mass haloes). This inefficient condensation occurs because $t_{\\rm age}/t_{\\rm cool}(1+\\beta) \\ll 1$ (typically $10^{-3}$--0.1) for these haloes due mostly to the fact that $\\beta \\gg 1$ (typically, $\\beta \\gsim 100$ for these haloes)---this makes star formation a very inefficient process in these galaxies, with the majority of the gas which was able to condense being re-ejected due to feedback.  Further plots of the condensed fraction for all cosmologies and models can be found in Appendix~\\ref{app:fcond}.  \\subsection{Effects of including photoionization and photodissociation feedback} \\label{sec:feedmechs}  We have run all models both with and without the effects of photoionization feedback (i.e. the suppression of accretion due to the IGM pressure and photoheating of gas in haloes by an ionizing background) and photodissociation feedback (i.e. the destruction of H$_2$ molecules by the dissociating background). In Fig.~\\ref{fig:phot} we show the resulting star formation rates and \\HII\\ filling factors as a function of redshift. Photoionization and photodissociation feedback reduce the redshift of reionization in general due to the suppression of galaxy and star formation. Fig.~\\ref{fig:phot} also illustrates once more that the inclusion of \\H2\\ cooling does little to increase the redshift of reionization. In fact, it appears that including \\H2\\ cooling actually reduces the redshift of reionization in the curves plotted here. This is a numerical artifact, arising from the fact that the most massive haloes simulated in our calculations are of lower mass in models which include \\H2\\ cooling (see \\S\\ref{sec:ulim}). Once this is corrected for, reionization redshifts are almost identical whether or not \\H2\\ cooling is included (see \\S\\ref{sec:epoch}).  \\begin{figure*} \\psfig{file=photo_effect_paper.ps,width=170mm,bbllx=0mm,bblly=75mm,bburx=195mm,bbury=270mm,clip=} \\caption{Star formation histories (blue and cyan lines) and \\HII\\ filling factors calculated using ionized spheres (red and magenta lines). Cyan and magenta lines show results for models where photoionization and photodissociation feedback are neglected, while blue and red lines show results for models including these feedback mechanisms. Results are shown for Standard and \\H2\\ models in WMAP PL, WMAP RSI and Baugh et al. '05 cosmologies as indicated for each column and row.} \\label{fig:phot} \\end{figure*}  Figure~\\ref{fig:fcond_photo} shows the same information as Fig.~\\ref{fig:fcond_nophoto} but for models which now include the effects of photo-ionization and photo-dissociation feedback. By comparing this figure with Fig.~\\ref{fig:fcond_nophoto}, it can be seen from this figure that the main causes of the suppression of galaxy formation are photoheating of hot gas in haloes and photodissociation of H$_2$. The filtering mass is approximately the same in models with and without photoionization and photodissociation feedback included\\footnote{Note that, even when photoionization feedback is not included the filtering mass still rises due to the heating of the IGM by ionizing photons emitted by galaxies. All that is left out of the calculation is the effect of the filtering mass on subsequent galaxy formation.} (it is not until after reionization that the filtering mass has the chance to respond to the increase in the temperature of the gas). At $z=5$ we can clearly see that the cooling times in haloes with virial temperatures below around $10^5$K are greatly increased due to this photoheating and photodissociation. At $z=10$ the effect is much weaker, due both to the lower background at that redshift and the fact that gas in haloes is denser at higher redshifts and so less affected by photoheating.  \\begin{figure*} \\psfig{file=color_fcond_photo.ps,width=160mm,bbllx=0mm,bblly=75mm,bburx=190mm,bbury=270mm,clip=} \\vspace{-15mm} \\caption{As Fig.~\\protect\\ref{fig:fcond_nophoto} but for models \\emph{including} the effects of photo-ionization and photo-dissociation feedback.} \\label{fig:fcond_photo} \\end{figure*}  It can be seen that the inclusion of photoionization and photodissociation feedback has only a small effect on the redshift of reionization. This is not surprising, as the mechanisms through which photoionization feedback works (photoheating of gas in haloes by the ionizing background and an increase in the filtering mass) become important only \\emph{after} reionization. Figure~\\ref{fig:feedmechs} shows the evolution of the \\HII\\ filling factor\\footnote{The ``spikes'' seen in the \\HII\\ filling factor are an artifact of the semi-analytic model. As in Paper~I, we compute the reionization process by drawing a set of halo masses from the halo mass function at some redshift $z$ and constructing merger trees for each halo back to very high redshifts. The galaxy formation calculation then proceeds in each tree. Our merger tree algorithm is not completely accurate in the sense that, when used to produce merger trees that span large redshift ranges, it does not exactly reproduce the expected distribution of progenitor halo masses. This problem becomes more severe as the redshift range spanned increases. Consequently, we repeat our calculations for several starting redshifts $z_i$ (ensuring that at each $z_i$ we have the correct distribution of halo masses). The reionization histories shown are constructed by splicing together results from each calculation, using quantities between $z_i$ and $z_{i+1}$ from each calculation. At each $z_i$ we find discontinuities in the predicted properties due to this splicing process.}, filtering mass and ionizing background for the standard model in the WMAP RSI cosmology with and without the effects of photoionization feedback included. The filtering mass begins to rise shortly before reionization occurs, when photoheating first dominates over cooling of gas due to the expansion of the Universe, but does not achieve any significant value until after the Universe is already reionized. Similarly, the ionizing background cannot build up to any significant level (and thereby efficiently photoheat gas in haloes) until the Universe has already been ionized and thereby become transparent to ionizing radiation. \\scite{benson02} found that it was the filtering mass which caused the greatest suppression of galaxy formation, seemingly in contradiction with the above result. \\scite{benson02} were, however, considering the effects of photoionization feedback on the galaxy population at $z=0$, by which time the filtering mass had had time to fully respond to the change in the IGM temperature during reionization. When considering the effects of photoionization feedback on the epoch of reionization itself, as we do here, the filtering mass is much less important as it does not have time to fully respond to the heating of the IGM until well after reionization is over. The ionizing background, however, rises very rapidly as reionization occurs, allowing photoheating to have some (albeit small) effect on the redshift of reionization.  \\begin{figure} \\psfig{file=photo_effect_mechs.ps,width=90mm} \\caption{Indicators of the strength of photoionization feedback mechanisms as a function of redshift for the Standard model in the WMAP RSI cosmology. Results are shown for this single model and cosmology only as they are typical of all of the models we consider. The solid red (blue) line in the upper panel shows the filtering mass (scale on left-hand axis) when photoionization feedback is (is not) included. In the lower panel red (blue) circles show the ionizing background intensity (scale on left-hand axis) when photoionization feedback is (is not) included. The dashed red line in both panels indicates the \\HII\\ filling factor (scale on right-hand axis) for a model including photoionization feedback.} \\label{fig:feedmechs} \\end{figure}   \\subsubsection{Effects of $H_2$ Dissociation?}  In general, we find that stars formed via $H_2$ cooling contribute little to the reionization of the Universe, if they form with a normal IMF. Therefore, it is not surprising that the inclusion or otherwise of $H_2$ dissociation in our model makes no significant difference to the end results on the ionized filling factor, as shown in Fig.~\\ref{fig:nodissoc}.  \\begin{figure} \\psfig{file=Ffill_nodissoc.ps,width=80mm} \\caption{The \\HII\\ filling factor as a function of redshift for model H$_2$ with (solid line) and without (dashed line) dissociation of $H_2$ molecules included in our calculations. Results are shown for the WMAP RSI cosmology only, but are typical of all other cosmologies and models.} \\label{fig:nodissoc} \\end{figure}   \\subsection{Enhanced Emission from Early Generations of Stars} \\label{sec:enhance}  In model ``H$_2$WeakFB-VMS'' we enhance the UV emission from stellar populations which form via \\H2\\ cooling (i.e. those which form in dark matter haloes with virial temperatures below $10^4$K) by a factor of 20 relative to a standard IMF, to account for the possibility that these stars might be super-massive ``pop~III'' stars.  The left-hand panel of Figure~\\ref{fig:VMS} shows the resulting \\HII\\ filling factors vs. redshift for this model and for model H$_2$WeakFB (which is identical apart from having no enhancement of UV emission due to supermassive stars) for the WMAP cosmologies. The results are plotted for an escape fraction chosen to give a double reionization: 25\\% (WMAP RSI) and 15\\% (WMAP PL). This results in all models producing full reionization at $z\\approx 6$. Note however, that model H$_2$WeakFB-VMS displays an extended period of partial reionization back to higher redshifts (from $z\\approx 8$ for the WMAP RSI cosmology to $z\\approx 17$ for the WMAP PL cosmology). As a consequence, model H$_2$WeakFB-VMS has an increased electron scattering optical depth as shown in the right-hand panel of Figure~\\ref{fig:VMS}.  For example in the WMAP PL cosmology we find $\\tau = 0.15$, compared to $0.08$ for model H$_2$WeakFB. Values of $\\tau$ are lower for the other cosmologies, but still come significantly closer to the WMAP measurement of $\\tau=0.17\\pm 0.04$. As such, model H$_2$WeakFB-VMS is in much better agreement with the data obtained from the WMAP satellite \\cite{kogut03}. Enhanced UV emission from stellar populations formed via \\H2\\ cooling is therefore able to reconcile the apparent disparity between a large optical depth and evidence for reionization at $z\\approx 6$ in the WMAP PL cosmology, by inducing a period of partial reionization.  The dotted line in Fig~\\ref{fig:VMS}(a) shows the contribution to the \\HII\\ filling factor from stars forming via \\H2\\ cooling in the H$_2$WeakFB-VMS model for the WMAP PL cosmology. Before $z\\approx 20$ these stars dominated the reionization of the Universe, as the steepness of the dark matter halo mass function results in there being very few haloes hot enough to cool via atomic processes at these redshifts. At lower redshifts the abundance of haloes with virial temperature above $10^4$K increases, while galaxy formation in lower temperature haloes becomes inefficient due to increased cooling times and the rising filtering mass\\footnote{The filtering mass in the H$_2$WeakFB-VMS model in the WMAP PL cosmology rises above the mass corresponding to a $10^4$K halo at $z=15$.}. Therefore, haloes hotter than $10^4$K rapidly become the dominant contributors to the \\HII\\ filling factor.  Obtaining a period of extended partial reionization, or a double reionization, requires an interplay of several key factors. Firstly, the cosmological model must produce enough haloes with virial temperatures below $10^4$K to permit a significant \\HII\\ filling factor to be obtained from the early generation of stars (assumed to be supermassive). The enhancement factor (taken to be 20 in this work) is also important: if it is too high then reionization may happen once only, if it is too low the early generation of stars will not be able to produce a significant \\HII\\ filling factor. The decline in the contribution from early generation stars as time progresses must also be sufficiently rapid (otherwise the filling factor will continue to rise, leading to a single reionization event)---requiring either that the filtering mass grows rapidly or that cooling times for haloes cooler than $10^4$K increase rapidly shortly after the first reionization. Despite this requirement that many factors combine to produce a double reionization, we find, for quite reasonable values of the UV enhancement factor, that double reionization is a natural outcome of our H$_2$WeakFB-VMS model in all cosmologies considered. An extensive survey of parameter space would be required to determine how common such double reionizations are, but this is beyond the scope of this work.  While our modelling of this enhanced UV emission from early stars is very simplified (e.g. we assume that supernova feedback is the same for Very Massive Stars and for normal stellar populations), and based on uncertain assumptions (e.g. that all haloes which cool by H$_2$ produce Very Massive Stars), it is still relatively sophisticated compared to previous studies of double reionization \\cite{cen03,wylo03} to the extent that ours is the only model which follows the full hierarchical merging history of each forming galaxy and, furthermore, computes the fraction of baryons able to cool and the fraction able to turn into stars using physically motivated prescriptions rather than simply making these fractions parameters of the model.  \\begin{figure*} \\begin{tabular}{cc} \\psfig{file=Ffill_VMS.ps,width=80mm} & \\psfig{file=tau_VMS.ps,width=80mm} \\end{tabular} \\caption{\\emph{Left-hand panel:} The \\HII\\ filling factor as a function of redshift for models H$_2$WeakFB-VMS (solid lines) and H$_2$WeakFB (dashed lines) in our three cosmologies (see key in right-hand panel). The dotted line shows, for the WMAP PL cosmology only, the contribution to the \\HII\\ filling factor from stars forming in haloes with virial temperatures below $10^4$K. Escape fractions have been chosen so as to produce double reionizations in each model, apart from the Baugh et al. (2005) model (for which we do not compute the H$_2$WeakFB-VMS model) in which case we simply choose $f_{\\rm esc}=1$. The escape fractions required are: $f_{\\rm esc}=0.25$, 0.15 and 1.0 for the WMAP RSI, WMAP PL and Baugh et al. (2005) cosmologies respectively. \\emph{Right-hand panel:} The electron scattering optical depth as a function of redshift for the same models.} \\label{fig:VMS} \\end{figure*}"
    },
    "0512/astro-ph0512078.txt": {
        "abstract": "New optical images of the young supernova remnant (SNR) G292.0+1.8, obtained from the 0.9-m telescope at CTIO, show a more extensive network of filaments than had been known previously.   Filaments emitting in \\oiii\\  are distributed throughout much of the 8 arcmin diameter shell seen in X-ray and radio images, including a few at the very outermost shell limits. In addition to the extensive \\oiii\\ filaments, we have detected four small complexes of filaments that show \\sii\\  emission along with the oxygen lines. In a single long-slit spectrum we find variations of almost an order of magnitude in the relative strengths of oxygen and sulfur lines, which must result from abundance variations.  None of the filaments, with or without \\sii, shows any evidence for hydrogen, so all appear to  be fragments of pure supernova ejecta.   The \\sii\\  filaments provide the first evidence for undiluted products of oxygen burning in the ejecta from the supernova that gave rise to G292.0+1.8.    Some oxygen burning, either hydrostatic or explosive, must have occurred, but the paucity of \\sii -emitting filaments suggests that either the oxygen burning was not extensive or that most of its products have yet to be excited. Most of the outer filaments exhibit radial, pencil-like morphologies that suggest an origin as Rayleigh-Taylor fingers of ejecta, perhaps formed during the explosion. Simulations of core-collapse supernovae predict the development of such fingers, but these have never before been so clearly observed in a young SNR. Following careful subtraction of the stars in the field, we have measured the total flux in \\oiii \\lam 5007 as $5.4 \\times 10^{-12} \\FLUX$.   Using a distance of 6 kpc and an extinction correction corresponding to $E(B-V) = 0.6$\\ mag (lower than previous values but more consistent both with our data and with X-ray and radio measurements of the hydrogen column density)  leads to a luminosity of $1.6 \\times 10^{35} \\LUM $\\ in the 5007 \\AA\\ line. ",
        "introduction": "The canonical picture of a supernova remnant (SNR) that results from the core collapse and supernova explosion of a massive star includes a shell produced by an expanding shock, fragments of ejecta that may appear as optical filaments rich in oxygen and (often) other heavy elements, a central neutron star that manifests itself as a pulsar, and a synchrotron-emitting pulsar wind nebula that surrounds and is powered by the pulsar.   Yet there is only a single SNR in the Galaxy that has been found to display all these properties:  G292.0+1.8, also known as MSH 11--5{\\it 4}, and its associated pulsar PSR J1124--5916 \\citep{hughes01, camilo02, hughes03}. This system thus provides a unique opportunity to study both the compact and ejected remains from a massive star and their interaction with the local environment.  Although G292.0+1.8 (henceforth more succinctly G292) was originally detected and identified as a SNR through its radio properties  \\citep{mills61, milne69}, it first attracted significant attention when  \\citet{goss79} discovered optical filaments with spectra dominated by lines of oxygen and neon.   Hydrogen Balmer lines, and also lines of \\sii\\ and \\nii\\  typically seen in SNRs, were conspicuous by their absence.   \\citet{murdin79} showed these filaments to have high radial velocities, $-700 \\kms \\lesssim v_{rad} \\lesssim +1300 \\kms$, indicating that they are undecelerated (or minimally decelerated) ejecta from the supernova explosion. G292 thus joined the prototype Cas A in the class of ``oxygen-rich'' SNRs---an exclusive group that still comprises only about 8 members.  The optical filaments with extremely metal-rich spectra that characterize the O-rich remnants are fragments of nearly pure ejecta that were launched from the core of the progenitor star during its explosion and  that remain virtually uncontaminated through interaction with interstellar or circumstellar material.    Subsequent spectra of G292 by \\citet{dopita84} and \\citet{sutherland95a} have shown no lines other than those of oxygen and neon.  G292 has been the subject of numerous X-ray studies, culminating in the spectacular high-resolution image from {\\it Chandra} \\citep{park02}. The X-ray spectrum is dominated by K-shell lines of O, Ne, Mg, Si, and S.  Non-equilibrium ionization analyses require large enhancements (relative to solar) in abundances for O, Ne, and Mg, with lesser enhancements for Si, S, and Fe \\citep{hughes94, gonzalez03}. By comparing the inferred abundances with the integrated yields predicted by  models for core-collapse supernovae, they estimated a progenitor mass of $\\sim20-40\\:M_\\sun$. In a more detailed study of individual X-ray features,  \\citet{park04} found that different knots have very different  compositions and suggested that  these represent clumps of ejecta from different zones of the progenitor.  The {\\it Chandra} image of G292 also revealed a compact central source, surrounded by what appeared to be a pulsar wind nebula \\citep{hughes01}. Shortly thereafter, \\citet{camilo02} discovered the radio pulsar PSR J1124--5916 within or very near G292.   \\citet{hughes03} then showed that the compact X-ray source is pulsed with the same period, confirming that it must be the compact remnant of the star that produced G292. The period of 135 ms and spin-down age of 2900 yr are roughly consistent with the age of the G292 SNR, estimated by \\citet{chevalier05} as 2700--3700 yr based on properties of the pulsar-wind nebula, and more directly as 3000--3400 yr from the kinematic study by \\citet{ghavamian05}.  Detailed radio images of G292  obtained by \\citet{gaensler03} show a bright central core around the pulsar, surrounded by a fainter plateau about 8\\arcmin\\ in diameter with sharp outer edges.   The outer edge of the radio plateau coincides closely with the outer extent of X-ray emission seen in the Chandra image, and must delineate the SNR's primary shock.  Based primarily on the \\hi\\ absorption profile, \\citet{gaensler03} also give what is probably the most reliable distance estimate to G292:  $6.2\\pm 0.9\\;$ kpc, which implies a diameter of about 15 pc for the SNR.   In this paper we present the first CCD images that cover the full extent of G292. The narrow-band  \\oiii\\ image, especially after we subtract a matched continuum image, shows a network of filaments far more extensive than indicated in photographic studies \\citep{goss79, tuohy82}, and even more wide-spread than those in the Fabry-Perot studies of \\citet{ghavamian05}. In addition, we have discovered a few filaments that show emission lines of \\sii\\ in addition to oxygen and neon---the first evidence for the presence of oxygen-burning products in the ejecta-dominated optical filaments. ",
        "conclusions": "\\subsection{Reddening, Spectral Diagnostics, and \\oiii\\  Luminosity}  \\citet{goss79}, in the paper where  G292 was first identified as an oxygen-rich SNR, estimated that the reddening was $E(B-V) = 0.9$\\ mag, arguing for consistency in the relative strengths of the oxygen lines. Although this argument depends entirely on highly uncertain models for the emission from shocked material composed almost entirely of metals, this original value for the reddening has persisted in the literature. An alternative approach to obtaining the reddening is through the hydrogen column density, $N_H$, which \\citet{predehl95} have shown to be closely correlated with optical extinction. They found that these quantities to be related by $N_H/A_V = 1.79 \\times 10^{21} \\;{\\rm cm^{-2}mag^{-1}}$, equivalent to $E(B-V) = 5.55 \\times 10^{21} \\;{\\rm cm^{-2}mag^{-1}}$.  There have been a number of recent measurements of $N_H$\\ based on X-ray and radio data.  \\citet{gonzalez03} found a value $(5 \\pm 1) \\times 10^{21}\\; {\\rm cm^{-2}}$\\ based on an average of fits to the spectra from numerous X-ray knots in G292, suggesting $E(B-V) = 0.9 \\pm 0.2$\\ mag.  \\citet{park04} fit the spectra of a number of individual X-ray knots and found values for $N_H$\\ ranging from 3 to 7.7 $\\times 10^{21}\\; {\\rm cm^{-2}}$, equivalent to $E(B-V)$\\ from 0.5 to 1.4 mag.  However, modeling of thermal X-ray plasmas with abundances that may be far from solar is difficult, and the absolute accuracy of fits based on these models is hard to assess.  For this reason, we believe the most reliable X-ray determination of the column density is that by \\citet{hughes01}, who fit  the pulsar spectrum with an absorbed power law to obtain $N_H \\approx (3.17 \\pm 0.15) \\times 10^{21}\\; {\\rm cm^{-2}}$, close to the minimum value derived by \\citet{park04}.  This lower value agrees well with the column obtained using radio data by \\citet{gaensler03}, who integrated the \\hi\\ emission over a range of velocities and inferred that $N_H \\approx 3.3 \\times 10^{21}\\; {\\rm cm^{-2}}$, equivalent to $E(B-V)=0.6$\\ mag.  The most successful models for the optical emission from young, oxygen-rich SNRs are those of \\citet{sutherland95b}, who calculated the emission resulting when a strong (reverse) shock enters an oxygen-rich fragment of material from a progenitor core. They considered both direct post-shock emission and also the indirect effects of  X-rays that can  produce an ionization front propagating into the fragment ahead of the shock itself. They used a grid of models considering either or both of these processes, with a range of shock velocities, and calculated diagnostic ratios for various oxygen lines. Comparing  their models with our observed spectra,  we find that dereddening as much as $E(B-V) = 0.9$\\ leads to intensities for  \\oii \\lam3727 that are stronger than any of their models predict. \\citet{sutherland95b} had precisely the same problem for G292, and indeed their own observed intensity ratios \\citep[from][where they also used  $E(B-V) = 0.9$]{sutherland95a} are almost exactly the same as ours for the bright Filament 1 East. But decreasing the reddening to $E(B-V) = 0.6$\\ gives ratios very similar to their model OSP150, which includes both direct post-shock emission and preionization from a $150 \\kms$\\ shock, and also similar to ratios they report for a number of filaments in the O-rich remnant N132D in the Large Magellanic Cloud. For example, the most reddening-sensitive ratio in their analysis is \\oii \\lam7325:\\oii \\lam3727.  For $E(B-V) = 0.6$, our dereddened intensity ratio for both the East and West regions of Filament 1 (see Table 3) is 0.028, almost exactly that predicted by the OSP150 model, while for $E(B-V) = 0.9$\\ it is half this value and substantially lower than has been observed in dereddened spectra of other O-rich SNRs \\citep[see e.g.][Fig.\\ 13c]{sutherland95b}.   Here and in the subsequent discussion we adopt the value $E(B-V)=0.6$\\ mag, which is consistent with most X-ray and radio measurements of the column density and gives relative line strengths that are consistent with the best models for a plasma composed entirely of heavy elements.  This in turn allows us to estimate the luminosity of G292 in \\oiii \\lam 5007.  Using our estimate (\\S\\thinspace 2) of the integrated flux, $F_{5007} = 5.4 \\times 10^{-12} \\FLUX$, and our adopted value for the extinction implies $L_{5007} = 1.6 \\times 10^{35}\\:d_6^{\\:2}\\; \\LUM$, where $d_6$\\ is the distance modulo 6 kpc. This luminosity value has an uncertainty $\\sim 20\\%$\\ attributable to our flux estimate, in addition to any uncertainty in the extinction.  One other O-rich SNR for which the \\oiii\\ luminosity has been reported is the extremely bright SNR in the Magellanic-type irregular galaxy NGC 4449, at a distance of about 3.9 Mpc \\citep{hunter99}.  For this case \\citet{blair83}\\footnote {The 4959 \\AA\\ and 5007 \\AA\\ lines in the NGC 4449 SNR are so broad that they cannot be resolved.  Note also that \\citet{blair83} used a distance of 5 Mpc and thus found even higher values for the luminosity.} give $F_{4959+5007} = 2.02 \\times 10^{-13} \\FLUX$\\@.  If we attribute $3/4$\\ of the total flux to the 5007 \\AA\\ line and correct after for extinction we find  $L_{5007} = 5.2 \\times 10^{38} \\LUM$, 300 times higher than G292. It has been recognized since its discovery, however, that the NGC 4449 SNR is extraordinary; its X-ray and radio luminosities are also higher than any other known remnants. It should be straightforward to obtain \\oiii\\ luminosities for most or all of the other O-rich SNRs, but no values are readily available in the literature.  \\subsection{\\sii\\ Emission and Nucleosynthesis Implications}  The detection of \\sii\\ emission, even if in only a handful of filaments, is nevertheless significant.   Sulfur resulting from massive-star supernovae stems almost entirely from oxygen burning, either hydrostatically in a shell surrounding the core during a short time before core collapse, and/or explosively during the supernova itself \\citep[e.g.][]{arnett96,woosley02}. The presence of S in fragments of pure SN ejecta (which these filaments must be since they are devoid of H) represents definitive evidence that some oxygen burning has occurred in the G292 progenitor.  The apparent absence of S or other O-burning products in optical spectra from G292 had been something of a puzzle, since models for the yields from core-collapse SNe all  predict from 0.04 to 0.2 $M_\\sun$\\ of S \\citep[e.g.][]{rauscher02}.  An analysis of the abundances in  filaments composed entirely of heavy elements represents a challenge beyond the scope of this paper; %\\citep[see][for a discussion of this subject]{sutherland95b}. nevertheless, the detection of only relatively faint \\sii\\ emission in a small minority of G292's filaments clearly indicates that S is not a prominent constituent in the optical filaments of this SNR\\@. Either there has been only a little oxygen burning, or most of the O-burning products have yet to be excited. Comparing G292 with other O-rich SNRs, only two of these have strong \\sii\\ lines in their optical spectra: Cas A \\citep[e.g.][]{chevalier78, fesen01} and 0540--69.3 in the Large Magellanic Cloud \\citep{dopita84, kirshner89}.   In the other members of this small society--- Puppis A \\citep{winkler85,sutherland95a}, N132D in the LMC and 1E0102--7219 in the SMC \\citep[e.g.][for both]{blair00}, the very bright SNR in NGC 4449 \\citep{blair83}, and the remnant of SN1957D in M83 \\citep{long89}--- \\sii\\ is extremely weak or absent entirely in the ejecta-dominated filaments.  \\subsection{Optical--X-Ray Comparison}  Comparison of the optical images with high-resolution X-ray images from {\\it Chandra} \\citep{park02} shows little detailed correspondence in features.  The brightest optical feature, filament 1 where we obtained our spectrum, has no exact X-ray counterpart; it lies $\\sim 30\\arcsec $\\ north of the eastern end of the bright equatorial X-ray belt.   This X-ray belt does not show enhanced abundances, and \\citet{park02} have interpreted it as a region where the primary supernova shock is encountering a ring of circumstellar material, seen approximately edge-on. The proximity of such a bright, ejecta-dominated  filament to one end of this projected ring may result from a strong reverse shock, produced indirectly by the encounter with the circumstellar ring, propagating  into ejecta expanding roughly transverse to our line of sight.  The low radial velocity observed for this filament, $v_r \\approx 100 \\kms$, is consistent with near-transverse motion.  \\citet{park02} show equivalent-width (EW) maps for strong X-ray lines of several metals. For oxygen, the highest equivalent widths are found in a complex within the southeast quadrant of G292.    The largest number of O-rich optical filaments are located along a broad arc extending  through this quadrant roughly from NNE to SSW (Fig.\\ 1). The Ne (He $\\alpha$) and Mg X-ray EW maps also show strong emission along a similar arc.  This suggests that pure ejecta fragments similar in composition to the ones we now observe have provided the material for the X-ray plasma in this region.  The {\\it Chandra} data partially complement the optical picture of oxygen-burning products inhomogeneously distributed around G292.   The \\citet{park02} EW map for Si shows a string of prominent X-ray knots extending almost two-thirds of the way across the northern portion of the shell.   The two optical filaments where the S/O ratio is highest are both found in this general region.  Filament 4 lies near the center of a relatively bright X-ray feature, designated as Region 5 by \\citet{park04} who found this region to be relatively more enhanced in Si and S than other portions of the SNR\\@. Our filament 3 lies about 40\\arcsec\\ to the NW, near the edge of another X-ray complex. Fig.\\ 8 shows a detailed optical/X-ray comparison for this portion of G292. On the version of the {\\it Chandra} image by \\citet{gonzalez03}, filaments 4 and 3 are located within their regions 16 and 17, respectively.   These contribute to the global abundances they calculated, but they did not report line fluxes or abundance estimates for individual regions.  A particularly interesting set of  O-rich filaments is at the extreme south, outside the field of the {\\it Chandra} ACIS image and at the extreme edge of X-ray emission detected from the {\\it ROSAT} HRI, as pointed out in \\S\\thinspace 2. Fragments of ejecta launched in this direction have traveled farther than any others, so these are probably the fastest such fragments in G292. It appears that  the primary shock is encountering a low-density region to the south, leading to more rapid expansion in this direction.  Perhaps pre-supernova mass loss has swept out the ISM in this direction, perpendicular to the equatorial ring seen in X-rays.   One might suspect a bipolar flow, leading to similar phenomena to the north, but as yet there is no evidence for any emission beyond the near-circular northern portion of the shell.   It would be interesting to obtain optical spectra of one or more of the extreme southern filaments and to investigate the X-ray structure here in more detail than can is possible with the {\\it ROSAT} data.  It is interesting to compare the total X-ray luminosity of G292 with that in \\oiii. \\citet{seward05} give values for the X-ray flux based on the {\\it Chandra} ACIS data, viz.\\ $F_{0.3-2.1\\:keV} = 1.80 \\times 10^{-10}\\; \\FLUX$, based on fits using a multi-component model with the absorption column density fixed at $N_H=6.16 \\times 10^{21}\\;{\\rm cm^{-2}}$\\@.  As we argue in \\S4.1, a column of $3.3 \\times 10^{21}\\;{\\rm cm^{-2}}$\\ is more consistent all the data for G292.  It is not a simple matter to adjust the absorption, since changing $N_H$\\ would change other parameters of the fit, but to first order we estimate that the lower absorption would give an unabsorbed flux of $F_0 \\approx  8 \\times 10^{-10} \\FLUX$, and an X-ray luminosity $L_{0.3-2.1\\: keV} \\approx 3 \\times 10^{36}\\:d_6^{\\:2}\\; \\LUM$, roughly 20 times the luminosity in \\oiii \\lam 5007\\@.   Comparing G292 with the extraordinary SNR in NGC 4449, \\citet{patnaude03} find for the latter $L_X = 2.4 \\times 10^{36} \\LUM$, just under half its luminosity in \\oiii\\@. \\footnote{The X-ray flux and luminosity given by \\citet{patnaude03} are for energy range 0.5--2.1 keV,  slightly different from the one used by \\citet{seward05} for G292.} The \\oiii/X-ray luminosity ratio for G292 is probably more typical of O-rich SNRs, but it would be interesting to do a quantitative comparison for others.   \\subsection{Filament Morphology---Rayleigh-Taylor Fingers}  The most striking feature of the outer filaments in G292 is that most of them display an elongated, finger-like morphology directed generally outward from the center of the remnant. A common feature of two- and three-dimensional hydrodynamic models for core-collapse supernovae is Rayleigh-Taylor instabilities, which form at interfaces between layers of different composition as the supernova shock pushes outward through the outer core of the progenitor \\citep[e.g.,][]{fryxell91, burrows95, nagataki98, kifonidis00, kifonidis03}. These instabilities produce mixing of material from different layers, as observed in SN1987A \\citep[e.g.,][]{arnett89, mccray93, wang02} and eventually may lead to SNR filaments of highly varied composition, as in Cas A \\citep[e.g.,][]{chevalier78, reed95, hurford96, fesen01}. The geometry of most of the outer filaments in G292 is very suggestive of Rayleigh-Taylor instabilities: thin, irregular fingers oriented radially, as illustrated by several examples shown in detail in Fig.\\ 9, taken from locations indicated in Fig.\\ 10.\\footnote{\\citet{ghavamian05} have suggested that the scalloped morphology of the bright spur filament may result from R-T instabilities that have resulted from the recent encounter of a reverse shock with an overdense shell of ejecta, and they have noted a similarity to filaments in Cas A \\citep{fesen01}. Our discussion here refers primarily to the outer filaments, whose finger-like geometry may be a remnant of R-T instabilities during the explosion.} The arrows in each panel of the figure are directed toward the geometric center of the outer radio shell, as defined by \\citet{gaensler03}:  $\\alpha(2000) = 11^{\\rm h}24^{\\rm m}34\\fs2,\\ \\delta (2000) = -59\\degr 15\\arcmin 54\\arcsec$\\@. \\citep[The pulsar is  located about 45\\arcsec\\ to the SE relative to this center, but it could well be moving rapidly, as] [have noted.]{hughes01} These filaments could have originated as R-T fingers, formed when the shock pushed outward through the oxygen zone in the early stages of the supernova explosion, that have been ballistically expanding since.  The ejecta clumps would have only recently become visible, once they entered the reverse shock and became excited. \\citet{blondin01} have shown that the fingers should stay near the forward shock front throughout the time that the reverse shock is in the power-law portion of the ejecta density profile---which can extend up to several thousand years.  \\notetoeditor{If possible we would like to have Figs. 9 and 10 on the same or facing pages where they can easily be compared.}"
    },
    "0512/astro-ph0512152_arXiv.txt": {
        "abstract": "Quantitative analyses of extreme helium stars to date face the difficulty that theory fails to reproduce the observed helium lines in their entirety, wings {\\it and} line cores.  Here, we demonstrate how the issues can be resolved using state-of-the-art non-LTE line formation for these chemically peculiar objects. Two unique B-type objects are discussed in detail, the pulsating variable V652\\,Her and the metal-poor star HD\\,144941. The improved non-LTE computations for helium show that analyses assuming LTE or based on older non-LTE model atoms can predict equivalent widths, for the He\\,{\\sc i}\\,10\\,830\\,{\\AA} transition in particular, in error by up to a factor $\\sim$3. Our modelling approach also succeeds in largely resolving the general mismatch for effective temperatures of EHe stars derived from ionization equilibria and from spectral energy distributions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512478_arXiv.txt": {
        "abstract": "We present Submillimeter Array observations of the Herbig Ae star HD169142 in 1.3~millimeter continuum emission and $^{12}$CO J=2-1 line emission at $\\sim1\\farcs5$ resolution that reveal a circumstellar disk. The continuum emission is centered on the star position and resolved, and provides a mass estimate of $\\sim0.02$~$M_{\\odot}$ for the disk. The CO images show patterns in position and velocity that are well matched by a disk in Keplerian rotation with low inclination to the line-of-sight. We use radiative transfer calculations based on a flared, passive disk model to constrain the disk parameters by comparison to the spectral line emission. The derived disk radius is $235$~AU, and the inclination is $13^{\\circ}$. The model also necessitates modest depletion of the CO molecules, similar to that found in Keplerian disks around T Tauri stars. ",
        "introduction": "Millimeter-wave observations of disks surrounding pre-main sequence stars are a promising probe of the planet formation process \\citep{bec96}. While most studies have concentrated on the circumstellar disks around low mass T Tauri stars, far less is known about the disks around their intermediate mass counterparts, the Herbig Ae/Be stars. Better characterization of the circumstellar disks around these higher mass stars will help address how stellar mass affects the development of planetary systems.  Detailed modeling of the spectral energy distributions of Herbig Ae/Be stars has generally supported the protoplanetary disk paradigm \\citep[e.g.][]{hil92,chi01,dom03}. But out of the large sample of Herbig Ae/Be star disk candidates recognized by infrared excess from circumstellar dust \\citep{the94,mal98}, only a handful have been spatially and spectrally resolved through direct imaging observations to show clear signatures of rotationally supported disks of Solar System size \\citep{man97,man00,pie03,cor05}.  The emission line star HD169142 (SAO 186777, MWC 925; spectral type A5) with infrared excess (IRAS 18213-2948) was included in the \\citet{the94} catalog of 287 Herbig Ae/Be candidates. Subsequent photometric and spectroscopic observations filled out its spectral energy distribution at infrared and submillimeter wavelengths \\citep{syl96,dun97,mal98}. At $\\sim145$~pc distance, HD169142 is notably isolated from any molecular cloud or extensive reflection nebulosity, though \\cite{kuh01} used differential polarimetry to detect a $\\sim200$~AU ($\\sim1\\farcs5$) diameter circular halo of scattered light. While some doubt lingers as to the pre-main-sequence nature of HD169142, it is among isolated Herbig Ae stars selected by \\cite{mee01} and classified as ``group Ib'' with a strong mid-infrared component, and a conspicuous absence of 10~$\\mu$m silicate emission \\citep{bou01}. This star also has prominent polycyclic aromatic hydrocarbon emission features \\citep{hab05}.  The spectral energy distribution of HD169142 is well fit by a passive irradiated disk model, though the near-infrared flux seems low, perhaps because the inner disk rim is viewed pole-on \\citep{dom03}, or perhaps because extensive inner disk clearing is underway, possibly together with planet formation.  The relative isolation of this star from any molecular clouds may indicate an advanced pre-main-sequence age. A circumstellar molecular component is detected via relatively strong emission in the CO J=3-2 line, with the narrow CO line profile modeled with a near face-on disk \\citep{gre00,den05a}.  Here we present high resolution observations of HD169142 from the Submillimeter Array\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA) that clearly resolve a circumstellar disk in Keplerian rotation. We derive the basic disk properties from arcsecond resolution images of the $^{12}$CO J=2-1 line using radiative transfer calculations based on a flared, passive disk model \\citep{dal05}. ",
        "conclusions": "We have imaged the Herbig Ae star HD169142 in 1.3~millimeter continuum emission and $^{12}$CO J=2-1 line emission at $\\sim1\\farcs5$ resolution. These observations resolve a disk of molecular gas in Keplerian rotation around the star. The presence of the disk is strong evidence that HD169142 is a {\\it bona fide} pre-main-sequence star, despite its distance from any molecular cloud. The star's spectral energy distribution, isolation, and the regularity of its disk velocity field suggest an advanced pre-main-sequence evolutionary state. The millimeter dust emission indicates a mass of $\\sim0.02$~$M_{\\odot}$ for the disk. Modeling the CO line emission shows that the disk has an outer radius of 235~AU and is viewed close to face-on, compatible with an associated scattered light nebula seen with differential polarimetry. The disk mass, size, and CO depletion are in line with those found for resolved Keplerian disks around other young stars."
    },
    "0512/astro-ph0512222_arXiv.txt": {
        "abstract": "Results from detailed spectroscopic analyses of stars hosting massive planets are employed to search for trends between abundances and condensation temperatures. The elements C, S, Na, Mg, Al,  Ca, Sc, Ti, V, Cr, Mn, Fe, Ni and Zn are included in the analysis of 64 stars with planets and 33 comparison stars. No significant trends are evident in the data. This null result suggests that accretion of rocky material onto the photospheres of stars with planets is not the primary explanation for their high metallicities. However, the differences between the solar photospheric and meteoritic abundances do display a weak but significant trend with condensation temperature. This suggests that the metallicity of the sun's envelope has been enriched relative to its interior by about 0.07 dex. ",
        "introduction": "Spectroscopic analyses of stars with planets (SWPs) have revealed that they, as a group, exhibit higher metallicity than randomly selected samples of nearby stars \\citep{l03, s04, FV05}. \\citet{g03} reviewed the three explanations proposed to account for this difference. One of these, the \"self enrichment\" hypothesis, posits that SWPs accreted sufficient mass of metals to enhance the metallicity of their photospheres by observable amounts.  \\citet{g97} suggested that accretion of metals by a star might be seen via its effects on the correlation between its photospheric abundances and the elemental condensation temperatures (T$_{\\rm c}$). This would be expected if, for example, a star accretes rocky material that formed significantly inside the water ice radius in the protoplanetary disk. Such material would be relatively more depleted in volatiles (e.g., C, N and O) than in refractories (e.g., Al, Ca and Ti). Material in the inner disk could find its way into the central star, for example, from gravitational scattering by a migrating giant planet.  \\citet{s01} were the first to search for anomalous trends between abundances and T$_{\\rm c}$ among SWPs. They compared the slopes from some 15 elements in 30 SWPs and 102 comparison stars. They found 5 or 6 SWPs with significantly higher abundance-T$_{\\rm c}$ slopes than the comparison stars. \\citet{g03} also noted that these stars tended to be hotter. However, \\citet{t01} and \\citet{s02} analyzed 14 and 12 SWPs, respectively, finding no significant correlations between abundances and T$_{\\rm c}$.  \\citet{gr01} studied six common proper motion main sequence stars. For one pair, HD 219542 A and B, they found a significant trend between the differences in their abundances and T$_{\\rm c}$. \\citet{s03} also studied the pair, but they failed to find such a trend.  Taken together, these studies leave the question of abundance-T$_{\\rm c}$ trends unresolved. Over the last two years, many more SWPs have been discovered and analyzed using detailed spectroscopic methods. In this study we revisit the question of abundance-T$_{\\rm c}$ trends among SWPs using the most recent data sets. We discuss the sample preparation, analysis and interpretation of results. We also search for a trend between the solar photospheric and meteoritic abundance differences and T$_{\\rm c}$. ",
        "conclusions": "Previous searches for anomalous correlations between photospheric elemental abundances and T$_{\\rm c}$ values among SWPs have suffered from either small or heterogeneous samples. We have revisited this topic using a homogeneous set of published abundance data. We fail to find any significant anomalies with respect to T$_{\\rm c}$ among a sample of 64 SWPs and 33 comparison stars.  Nevertheless, we encourage others to revisit this topic as more SWPs are discovered and subjected to detailed spectroscopic analyses. The signature of accretion may reveal itself either as a general trend of abundance-T$_{\\rm c}$ slopes with T$_{\\rm eff}$ or as anomalously high abundance-T$_{\\rm c}$ slopes in a few hotter stars. It would also be helpful to extend the comparison stars sample above [Fe/H] $= 0.2$.  Relative to the chondritic abundances, the solar photospheric abundances do display a weak trend with T$_{\\rm c}$. It is seen in two recent compilations of Solar Systems abundances. The implications of this are far ranging, and we encourage additional work on improving photospheric abundance determinations."
    },
    "0512/astro-ph0512014_arXiv.txt": {
        "abstract": "Near comoving wavenumber $k$, the gravitational-wave background (GWB) from inflation carries information about the physical conditions near two moments in cosmic history: the moment when $k$ ``left the horizon'' during inflation, and the moment when it ``re-entered the horizon'' after inflation.  We investigate the extent to which this information can be extracted if the GWB is measured by a combination of cosmic-microwave-background (CMB) polarization experiments on large scales and space-based laser-interferometer experiments on small scales.  To disentangle this information, we derive a new gravitational-wave transfer function that incorporates a number of physical effects that were treated less accurately, less generally, or were missing altogether in previous treatments.  In particular, it incorporates: {\\it (i)} dark energy with time-varying equation-of-state $w(z)$; {\\it (ii)} tensor anisotropic stress due to free-streaming relativistic particles in the early universe; and {\\it (iii)} a variety of physical effects that cause deviations from the standard equation-of-state $w=1/3$ during the radiation era.  Based on this transfer function, we consider the degree to which the GWB can be used to test inflation and to probe the ``primordial dark age'' between the end of inflation and the electroweak phase transition. ",
        "introduction": "\\label{intro}  Inflation \\cite{inflation} generates tensor perturbations (gravitational waves) with a nearly scale-invariant primordial power spectrum \\cite{early_gwave}.  This gravitational-wave background (GWB) contributes to the cosmic microwave background (CMB) temperature anisotropy \\cite{early_cmb,cmb_temp}, and also produces a characteristic curl (or ``B-mode'') fingerprint in the CMB polarization \\cite{cmb_polar,curlmodes,bmodes}.  Several CMB experiments are being developed to pursue this signal \\cite{cmb_experiments}.  The GWB also persists as a sea of relic gravitational radiation filling the universe today \\cite{turner_bound,SmithKamCoor}.  Direct detection of this relic radiation has received considerable attention over the past year or so, since it has been realized that space-based laser interferometers operating in the frequency range $0.1~{\\rm Hz}<f<1~{\\rm Hz}$ might achieve the necessary sensitivity and foreground subtraction \\cite{DECIGO,Phinney}.  In particular, two futuristic experiments have been proposed --- NASA's ``Big Bang Observer'' (BBO) and the Japanese ``Deci-hertz Interferometer Gravitational Wave Observatory'' (DECIGO) --- and are currently being investigated \\cite{BBO,Ungarelli,Crowder}.  The gravitational wave spectrum generated by inflation carries important information about the conditions during inflation.  But the spectrum also receives corrections, both large and small, from the subsequent evolution and matter content of the universe after inflation.  In this paper, we identify various post-inflationary physical effects which modify the GWB, and show how they may be encoded in the gravitational-wave transfer function that relates the primordial tensor power spectrum to the gravitational-wave spectrum at a later point in cosmic history.  It is necessary to properly understand and disentangle the post-inflationary effects in order to optimally extract the inflationary information in the GWB.  But these modifications are also interesting in their own right, since they offer a rare window onto the physical properties of the high-energy universe during the ``primordial dark age'' between the end of inflation and the electroweak phase transition.  The same features that make the inflationary GWB difficult to detect --- namely its small amplitude and gravitational-strength coupling to matter --- also make it a clean cosmological probe.  First, because of their tiny amplitude, the gravitational waves obey linear equations of motion, so that their evolution may be predicted analytically with high precision.  (By contrast, density perturbations grow after horizon entry, and perturbations shorter than $\\sim10$~Mpc have already grown non-linear; so it is impossible to study their evolution analytically, and even numerical analysis is challenging.)  Second, because of their ultra-weak interactions with matter, the gravitational waves have been free streaming since the end of inflation --- in contrast to neutrinos (which began streaming roughly a second later) and photons (which began streaming several hundred thousand years later). The gravitational waves carry unsullied information from the early universe, and subsequent modifications of the gravitational-wave spectrum are not washed out by thermal effects (since the gravitons are thermally decoupled).  The gravitational-wave spectrum near a given wavenumber $k$ is primarily sensitive to two ``moments'' in cosmic history: (1) the moment when that mode ``left the horizon'' ({\\it i.e.}, became longer than the instantaneous Hubble radius during inflation), and (2) the moment when the mode ``re-entered the horizon'' ({\\it i.e.}, became shorter than the instantaneous Hubble radius once again, after the end of inflation).  The first moment imprints information about inflation itself, while the second moment imprints information about post-inflationary conditions.  The CMB is sensitive to long-wavelength modes that re-entered at relatively low temperatures (well after BBN), corresponding to relatively well-understood physics.  By contrast, laser interferometers are sensitive to shorter wavelengths that entered the horizon at high temperatures ($T\\sim10^{7}$~GeV), well above the electroweak phase transition.  The physical conditions at such high energies, which are considerably beyond the reach of particle accelerators, are a mystery, so that any information about this epoch from the GWB would be very valuable.  The outline of our paper is as follows.  In section \\ref{fundamentals}, we review the basics of inflationary gravitational waves.  We take special care to clearly establish our conventions, and to point out where conventions diverge and become confused in the literature.  In section \\ref{transfer_fn} we present a new calculation of the gravitational-wave transfer function which includes effects not considered in previous calculations \\cite{SmithKamCoor,TurnLidsWhite,Bashinsky}.  We separate the transfer function into three factors, each with a distinct physical meaning.  The first factor accounts for the redshift-suppression of the gravitational-wave amplitude after horizon re-entry.  Among other things, this factor accomodates a dark energy component with a time-varying equation-of-state parameter $w(z)$.  The second factor captures the behavior of the background equation-of-state parameter $w$ near the time of horizon re-entry.  The third factor accounts for the damping of tensor modes due to tensor anisotropic stress from free-streaming relativistic particles in the early universe.  This damping effect was recently pointed out by Weinberg \\cite{Weinberg}, and the damping on CMB scales due to free-streaming neutrinos has been studied in \\cite{Weinberg,PritchKam,Bashinsky,DicusRepko}.  In this paper, we stress that it is also necessary to consider this damping effect on laser-interferometer scales, which re-entered the horizon when free-streaming particles were an unknown fraction $f$ of the background energy density.  We present accurate expressions for the damping effect as a function of $f$.  We also observe that Weinberg's analysis, which originally focused on a single fermionic particle (the neutrino), extends in a simple way to the more general case of a mixture of free-streaming bosons and fermions with different temperatures and decoupling times.  We identify six physical effects which can modify the relic GWB by causing the equation-of-state parameter $w$ to deviate from its standard value ($w=1/3$) during the radiation-dominated epoch. Furthermore, although it is often treated as a stationary random process, the inflationary GWB is actually highly non-stationary (as emphasized by Grishchuk \\cite{GrishchukSidorov}; also see \\cite{AllenFlanaganPapa,Albrecht}).  Thus, our transfer function keeps track of the coherent phase information that it contains.  Finally, in section \\ref{discussion}, we discuss the implications of these results for future CMB and laser-interferometer experiments, and consider what we might learn by measuring the GWB in both ways. ",
        "conclusions": "\\label{discussion}  \\begin{figure} \\begin{center} \\includegraphics[width=3.1in]{boylefig2.eps} \\end{center} \\caption{The black solid curve represents the present-day gravitational-wave energy spectrum, $\\Omega_{gw}^{}(f,\\tau_{0})$, for the inflationary model $V(\\phi)=(1/2)m^{2}\\phi^{2}$.  The red dotted curve shows the damping effect due to (three ordinary massless species of) free-streaming neutrinos.  The green dot-dashed curve shows the damping effect which arises if free-streaming particles make up fifty percent of the background energy density at the time $\\tau_{BBO}^{}$ when the modes probed by BBO/DECIGO re-enter the horizon.  As shown in the figure, the particles begin free-streaming sometime before $\\tau_{BBO}^{}$, and decay sometime after $\\tau_{BBO}^{}$, but prior to electroweak symmetry breaking.  Finally, the blue dashed curve shows the effect of a conformal anomaly in the early universe that slightly reduces the equation of state from $w=0.33$ to $w=0.31$ above the electroweak phase transition.  The spectrum will also be modified on comoving scales that re-enter the horizon during the reheating epoch after inflation; but the range of scales affected by reheating is unknown.  Finally, note that the correlated BBO interferometer proposal claims a sensitivity that extends beyond the bottom of the figure (down to roughly $\\Omega_{gw}^{}\\sim10^{-17}$) in the frequency range from $10^{-1}$~Hz to $10^{0}$~Hz.} \\label{energy_plot} \\end{figure}  Fig.~\\ref{energy_plot} illustrates some of the effects discussed in section \\ref{transfer_fn}.  In this figure, the solid black curve represents the present-day energy spectrum, $\\Omega_{gw}(f, \\tau_{0}^{})$, generated by a particular inflationary model --- namely, a quadratic potential $V(\\phi)=(1/2)m^{2}\\phi^{2}$.  The red dotted curve illustrates the damping due to tensor anisotropic stress from free-streaming neutrinos.  We have assumed that the free-streaming fraction is $f=0.4052$, which is the $f$ value for three standard neutrino species which decouple around the time of BBN. The green dot-dashed curve represents the damping due to tensor anisotropic stress from various particle species ($X$ particles) which begin free-streaming before the scales detected by BBO/DECIGO re-enter the horizon and then decay after the scales re-enter, but prior to the electroweak phase transition.  As an example, we have assumed that the free-streaming fraction is $f=0.5$.  Finally, the blue dashed curve represents damping due to a trace anomaly that is present above the electroweak scale.  For illustration, we have assumed that this anomaly, through Eq.\\ (\\ref{trace_anom}), reduces the equation of state from $w_{r}^{}=1/3$ by $\\delta w_{r}^{}=-0.02$.  This reduction may be achieved by various combinations of the number of colors $N_{c}^{}$, the number of flavors $N_{f}^{}$, and the gauge coupling $g$; but the point is that we have not chosen an unreasonable large value for $\\delta w_{r}^{}$, given the large gauge groups that are often theorized to be present at high energies.  The key point conveyed by Fig.~\\ref{energy_plot} is that there are a variety of plausible post-inflationary effects that can produce rather large modifications of the gravitational-wave spectrum on laser-interferometer scales, without modifying the spectrum on CMB scales.  This is tantalizing, since the modifications on laser-interferometer scales reflect the primordial dark age between the end of inflation and the electroweak phase transition, at energies beyond the reach of terrestrial particle accelerators.  Let us look ahead to future observations.  First, as measurements of the scalar power spectrum improve, we will be restricted to considering inflationary models that match precisely the measured scalar perturbation spectrum on CMB scales.  This is not enough to specify a unique model because there is a range of inflaton potentials whose scalar perturbation spectra are indistinguishable on CMB scales. For example, two such models are shown in Fig.~\\ref{primordial}. \\begin{figure} \\begin{center} \\includegraphics[width=3.1in]{boylefig3.eps} \\end{center} \\caption{The precise predictions of the primordial tensor power spectrum for two inflationary models whose parameters have been chosen so that they make identical predictions for the scalar fluctuation amplitude and spectral index on CMB scales.  The two models are the quartic monomial potential $V(\\phi)=\\lambda\\phi^{4}$ (top), and the axion (or ``natural inflation'') potential $V(\\phi)=M_{I}^{4} [1+{\\rm cos}(\\phi/\\mu)]$.} \\label{primordial} \\end{figure} The two curves correspond to two different inflationary models whose parameters have been chosen so that they produce the same scalar amplitude and tilt on CMB scales, and match the observed scalar amplitude.  From the figure, though, we see that their primordial tensor power spectra are distinguishable.  The upper solid curve corresponds to the simple quartic monomial potential $V(\\phi)=\\lambda\\phi^{4}$, and the lower solid curve corresponds to the axion (or ``natural inflation'' \\cite{NaturalInflation}) potential $V(\\phi)=M_{I}^{4}[1+{\\rm cos}(\\phi/\\mu)]$.  Note that the two spectra ``curve'' significantly downward over the broad range of wavenumbers shown in Fig.~\\ref{primordial}, so that they are not well described by a power-law (straight-line) approximation over this range.  In particular, a power-law extrapolation based on CMB measurements will give a fairly poor estimate of the real spectrum on laser-interferometer scales.  In fact, since the real spectra tend to redden increasingly on smaller scales ({\\it i.e.}\\ as the end of inflation approaches), the CMB power-law extrapolation will tend to overestimate the tensor spectrum on laser-interferometer scales.  Second, note that the inflationary consistency relation (\\ref{consistency}) implies that a gravitational-wave spectrum with a larger tensor/scalar ratio $r$ on CMB scales also tends to have a ``redder'' (more negative) tilt $n_{t}$.  For gravitational waves with detectably-large amplitudes on CMB scales, this leads to the ``convergence effect'' in Fig.~\\ref{primordial}: if two different inflationary models produce scalar perturbations with the same amplitude and tilt on large (CMB) scales, then their tensor spectra tend to approach and even cross one another on small (laser-interferometer) scales, as illustrated by the two solid black curves in Fig.~\\ref{primordial}.  As a result, once one has measured the scalar and tensor perturbations on CMB scales, one expects the tensor spectrum on laser-interferometer scales to lie within a relatively narrow and predictable range.  This is both good news and bad news for laser-interferometer science. The bad news is that the convergence effect tends to make laser interferometers less effective at breaking any remaining degeneracies between inflationary models because their predictions for the tensor spectrum are so similar on small scales.  The degree to which laser interferometers can break such degeneracies has been studied in \\cite{SmithKamCoor}.  The good news is that the measurements can potentially provide a model-independent test of a key inflationary prediction: that the tensor spectrum observed at long wavelengths really does extend to small wavelengths, as expected from the consistency relation (\\ref{consistency}).  Furthermore, precise laser-interferometer measurements can serve as sensitive probes of new post-inflationary physics that modifies the transfer function on small scales.  In this light, let us now imagine some time in the future in which we have successfully measured the tensor amplitude on CMB scales, and consider four possible outcomes of a laser-interferometer experiment (like BBO).  First, suppose that BBO detects the gravitational-wave amplitude at the expected value, as exemplified in Fig.~\\ref{primordial}.  This would be a quantitative confirmation of inflation generally, and of the inflationary consistency condition (\\ref{consistency}) in particular, since the large difference in wavenumber between CMB and laser-interferometer scales provides a large lever-arm to measure the tensor tilt $n_{t}^{}$.  It would also suggest an upper bound on the size of any exotic effects in the transfer function on laser-interferometer scales.  Second, suppose that BBO detects the gravitational-wave amplitude significantly {\\it below} its expected value.  Then the interpretation is less straighforward.  On the one hand, we could interpret the suppression as being due to transfer-function effects, such as those discussed above.  With this interpretation, the suppression would imply a rare opportunity to measure physical properties of the early universe, at temperatures above the electroweak scale, when the relevant modes re-entered the horizon.  On the other hand, we could interpret the suppression to mean that inflation is more complicated than we expected: perhaps there is a feature in the inflaton potential which suppresses the primordial tensor spectrum on small scales relative to our expectations, or perhaps the consistency condition itself is violated (as it is in some multi-field models of inflation). Further measurements of the tensor spectrum at intermediate scales are needed to resolve this issue.  Third, suppose that BBO detects the gravitational-wave amplitude significantly {\\it above} its expected value.  At this point it is worth noting that virtually all of the transfer-function effects mentioned in section \\ref{transfer_fn} (including tensor anisotropic stress, the late decay of a massive relic species, bulk viscosity, a conformal anomaly, or standard reheating with $w\\leq 1/3$) {\\it suppress} the gravitational-wave spectrum on small scales, relative to large scales.  In order to {\\it enhance} the spectrum on laser-interferometer scales, one must invoke an even more exotic effect, such as a reheating epoch with a low reheat temperature and an unusual equation of state ($w>1/3$); or the production of gravitational waves {\\it after} inflation due to the collision of bubbles from a phase transition; or perhaps extra-dimensional physics.  Fourth, suppose that BBO detects nothing at all.  This would most likely indicate a fundamental problem with inflation, since an period of accelerating expansion produces a broad and nearly-flat spectrum of primordial tensor perturbations quite generically (regardless of the particular model that drives inflation).  One way out of this conclusion would be to invoke an extreme suppression effect in the transfer function on BBO scales --- {\\it e.g.} due to a reheating epoch with equation of state $w=0$ and a reheat temperature well below $10^{7}$~GeV.  Determining which interpretation is correct would require digging deeper, with more sensitive experiments on BBO scales."
    },
    "0512/astro-ph0512146_arXiv.txt": {
        "abstract": "We evaluate the effect of variations in the electron-impact excitation cross sections on the non-LTE line formation for hydrogen in early-type stars. While the Balmer lines are basically unaffected by the choice of atomic data, the Brackett and Pfund series members allow us to discriminate between the different models. Non-LTE calculations based on the widely-used approximations of Mihalas, Heasley \\& Auer and of Johnson fail to simultaneously reproduce the observed optical and IR spectra over the entire parameter range. Instead, we recommend a reference model using data from {\\em ab-initio} calculations up to principal quantum number $n$\\,$\\leq$\\,7 for quantitative work. This model is of general interest due to the ubiquity of the hydrogen spectrum. ",
        "introduction": "The quantitative interpretation of the hydrogen line spectrum is one of the foundations of modern astrophysics. Being the most abundant and most basic element in the universe hydrogen imprints its signature on the spectra of stars, nebulae and accretion phenomena. For decades the focus laid in the modelling of the first members of the Balmer series for deriving the physical properties of astronomical objects. In the meantime developments in instrumentation have opened the IR window to routine observation, which will even gain in importance in the future, with {\\sc Crires} and {\\sc Visir} on the {\\sc Vlt} being important cornerstone projects. The next generation of ground-based large telescopes and the next large space telescope will focus on this wavelength range, primarily to study the high-$z$ universe. But it will also allow to investigate local objects in otherwise inaccessible environments, e.g. ultra-compact H\\,{\\sc ii} regions, the Galactic centre and dust-enshrouded nearby starburst galaxies. The Brackett and Pfund lines are among the key diagnostics at these wavelengths. Here we report on the findings of our reinvestigation of the non-LTE line-formation problem for hydrogen~\\cite{PrBu04}. We confront H\\,{\\sc i} model atoms of different degree of sophistication with observations of early-type stars in order to derive a set of reference data which is, of course, of much broader interest than for stellar analyses alone. ",
        "conclusions": ""
    },
    "0512/astro-ph0512370_arXiv.txt": {
        "abstract": "Published $B$ and $V$ fluxes from nearby Type Ia supernov\\ae\\ are fitted to light-curve templates with 4-6 adjustable parameters.  Separately, $B$ magnitudes from the same sample are fitted to a linear dependence on $\\bv$ color within a post-maximum time window prescribed by the {\\sc cmagic} method.  These fits yield two independent SN magnitude estimates $B_{\\rm max}$ and $B_{BV}$.  Their difference varies systematically with decline rate $\\Delta m_{15}$ in a form that is compatible with a bilinear but not a linear dependence; a nonlinear form likely describes the decline-rate dependence of $B_{\\rm max}$ itself.  A Hubble fit to the average of $B_{\\rm max}$ and $B_{BV}$ requires a systematic correction for observed $\\bv$ color that can be described by a linear coefficient ${\\cal R} = 2.59 \\pm 0.24$, well below the coefficient $R_B \\approx 4.1$ commonly used to characterize the effects of Milky Way dust.  At 99.9\\% confidence the data reject a simple model in which no color correction is required for SNe that are clustered at the blue end of their observed color distribution.  After systematic corrections are performed, $B_{\\rm max}$ and $B_{BV}$ exhibit mutual rms intrinsic variation equal to $0.074 \\pm 0.019$ {mag}, of which at least an equal share likely belongs to $B_{BV}$.  SN magnitudes measured using maximum-luminosity or {\\sc cmagic} methods show comparable rms deviations of order $\\approx$0.14 {mag} from the Hubble line.  The same fit also establishes a 95\\% confidence upper limit of 486 km s$^{-1}$ on the rms peculiar velocity of nearby SNe relative to the Hubble flow. ",
        "introduction": "\\label{intro}  Well before the Universe was found to be accelerating \\citep{Perlmutter:1998, Garnavich:1998, Schmidt:1998, Riess:1998, Perlmutter:1999}, Type Ia supernov\\ae\\ (SNe), the study of which underpinned that discovery, were thought to be standard candles.  Even before the Type Ia subclass was identified, \\citet{Kowal:1968} studied the distribution of Type I SN luminosities, and \\citet{Pskovskii:1977} observed that brighter SNe exhibit slower decline rates.  The standardization of Type Ia SNe sharpened when \\citet{Phillips:1993} introduced a brighter-slower correction that was linear in the decline-rate parameter $\\Delta m_{15}$ (the change in magnitude from maximum light to 15 rest-frame days thereafter).  That correction was refined by \\citet{Hamuy:1996a} and by \\citet{Phillips:1999} (henceforth Phil99), who added a quadratic term.  \\citet{Tripp:1998}, \\citet{Tripp:1999}, and \\citet{Parodi:2000} fit a correction that was linear both in $\\Delta m_{15}$ and in a second parameter, $\\bv$ color.  In these two-parameter studies the observed brighter-bluer correlation was compensated by a fit color coefficient ${\\cal R} \\approx 2.5\\,$.  \\citet{Perlmutter:1997} and \\citet{Goldhaber:2001} (henceforth Perl97 and Gold01) stretched the time axis of a fixed template to approximate different SN light curves; that linear stretch factor substituted for $\\Delta m_{15}$ in their alternative brighter-slower correction.  Given an adequate training set -- here an ensemble of SNe in the nearby Hubble flow whose deviations $\\Delta_i$ from a fiducial absolute SN magnitude are deduced from their redshifts -- the problem of estimating the deviation $\\Delta$ of a single SN, using a set $\\{x_j\\}$ of its measured parameters, is amenable to general solution.  In principle the $x_j$ can include any parameter that may add to knowledge of $\\Delta$: fluxes in various bands at various phases, spectral line amplitudes, {\\it etc}.  One example of such a solution is given in Appendix \\ref{appS}.  Present progress along these more general lines is exemplified by the {\\sc mlcs}2k2 package and its predecessors described by \\citet{Riess:1995,Riess:1996}, \\citet{Riess:1998}, and \\citet{Jha:2002}.  As those authors found, multivariate approaches to SN standardization are severely limited by low training statistics.  For example, the {\\sc mlcs}2k2 training set did not adequately constrain the color-correction coefficient $R_V$, for which an external value was imposed.  This paper aims to extend understanding of the systematic corrections that help standardize Type Ia supernova luminosities, and of the intrinsic variations which blur that standard.  Large resources now are devoted to projects that use standardized SNe to measure with exquisite precision the parameters determining the Universe's recent acceleration, and even larger efforts will be needed to track the evolution of its equation of state, revealing the sources of that acceleration.  Even small advances in such understanding can have big implications for the realization of these projects.  In this paper, we make use of only a few properties of nearby SNe -- luminosity, color and decline rate -- that are simple to characterize.  In that sense our approach to SN standardization is traditional, in the vein of work cited in the first paragraph.  Nevertheless, we do bring fresh tools to this study.  By taking the difference between the two simplest measures of SN magnitude -- luminosity at peak, and luminosity at fixed color using the {\\sc cmagic} method \\citep[][henceforth Wang03]{Wang:2003} -- we impose a new, more precise constraint on the (nonlinear) form of their decline-rate dependence, and we set a robust lower bound on the intrinsic variation  suffered by at least one of these measures.  By applying a newly extended light-curve fitting method and standard error propagation consistently to each SN in the sample, we obtain a Hubble fit that sets a significant upper bound on the rms peculiar velocity of SNe relative to the Hubble flow.  In Section \\ref{fits} our light-curve ({\\sc lc}) and {\\sc cmagic} fitting methods are introduced.  Their outputs are shown to be independent, and the corrections made to them for extinction and time dilation are described.  Section \\ref{fitPhoto} identifies the photometric data sample that is processed by these methods, and the subsample of SNe to which Hubble fits are applied.  Section \\ref{DeltaE} reveals the new information gained by comparing the {\\sc lc} and {\\sc cmagic} fit magnitudes.  In Section \\ref{Bavg} the result of a Hubble fit to the average of these outputs is presented; this is the final step in a global fit to the sample.  Section \\ref{Bmax} uses the global fit parameters to quantify the deviations of the {\\sc lc} and {\\sc cmagic} outputs from the Hubble line.  Section \\ref{blue} explores properties of the subset of SNe assumed to be largely free of extinction by host-galactic dust.  The paper concludes with a summary and discussion. ",
        "conclusions": ""
    },
    "0512/astro-ph0512416_arXiv.txt": {
        "abstract": "We investigate the dependence of dark matter halo clustering on halo formation time, density profile concentration, and subhalo occupation number, using high-resolution numerical simulations of a LCDM cosmology.  We confirm results that halo clustering is a function of halo formation time at fixed mass, and that this trend depends on halo mass.  For the first time, we show unequivocally that halo clustering is a function of halo concentration and show that the dependence of halo bias on concentration, mass, and redshift can be accurately parameterized in a simple way: $b(M,c|z) = b(M|z) \\brel(c|M/M_*)$. Interestingly, the scaling between bias and concentration changes sign with the value of $M/M_*$: high concentration (early forming) objects cluster more strongly for $M \\lsim M_*$, while low concentration (late forming) objects cluster more strongly for rare high-mass halos, $M \\gsim M_*$.  We show the first explicit demonstration that host dark halo clustering depends on the halo occupation number (of dark matter subhalos) at fixed mass, and discuss implications for halo model calculations of dark matter power spectra and galaxy clustering statistics.  The effect of these halo properties on clustering is strongest for early-forming dwarf-mass halos, which are significantly more clustered than typical halos of their mass.  Our results suggest that isolated low-mass galaxies (e.g. low surface-brightness dwarfs) should have more slowly-rising rotation curves than their more clustered counterparts, and may have consequences for the dearth of dwarf galaxies in voids.  They also imply that self calibrating richness-selected cluster samples with their clustering properties might overestimate cluster masses and bias cosmological parameter estimation. ",
        "introduction": "\\label{section:intro}  The spatial distribution of galaxies is now well established to be dependent on several of their internal properties: stellar mass, luminosity, color, star formation rate, Hubble type, and several others \\citep[e.g.,][]{hubble:36,dressler:80,norberg_etal:01,zehavi_etal:02}. In the current paradigm for galaxy formation, this can be understood as a combination of the fact that dark matter halos with different masses and formation histories cluster differently and host different galaxy populations.  A full understanding of these trends is one of the primary goals of modern cosmology, as it is likely to provide insight into the physical process that govern galaxy formation and aid in the use of observed galaxy clustering as a probe of fundamental cosmological parameters.  Many models of galaxy formation and methods of calculating galaxy clustering statistics make two related simplifying assumptions.  The first is that the number and properties of galaxies within a host dark matter halo depend solely on the mass of the host, independent of halo environment or other properties of the dark halo.  The second is that the clustering properties of dark matter host halos are a function only of their masses and that dark matter halos are otherwise ignorant of their larger environments.  The latter assumption forms part of the basis of the excursion-set formalism for galaxy clustering \\citep{bond_etal:91}, at least in its simplest and most common implementation \\citep{lacey_cole:93,somerville_kolatt:99}.  This implementation assumes that halo formation is a Markov process with no correlations between different spatial scales, which then implies that future halo accretion is independent from past history, and that halo histories are independent of environment (see also discussion in \\citealt{white:96} and \\citealt{sheth_tormen:04}).  The first place that these assumptions are made is in semi-analytic models of galaxy formation.  A basic assumption of the technique is that the properties of galaxies depend only on the mass and formation time of the host halo.  In many implementations, galaxy clustering is calculated by filling simulated halos at a fixed redshift with halo formation histories that are calculated analytically using the excursion-set formalism \\citep[e.g.][]{kauffmann_etal:97,benson_etal:00,wechsler_etal:01,zentner_etal:05}. By construction, any dependence of galaxy properties on environment in these models must come only from the extent to which they populate halos of different masses.  Note that this assumption is avoided to some extent in many modern implementations which use halo merging histories extracted directly from N-body simulations \\citep[e.g.][]{springel_etal:01, helly_etal:03, hatton_etal:03, springel_etal:05, kang_etal:05, croton_etal:05, delucia_etal:06, bower_etal:05}.  In these implementations semi-analytic recipes depend on mass and formation history as before, but correlations between the galaxy properties and environment may now come from the extent to which they populate halos of different masses {\\em and formation histories}.  If there are physical effects that otherwise depend on the larger scale environment these would not be included.  The second place these assumptions are commonly made is in the standard halo model of galaxy clustering (e.g., \\citealt{seljak:00,peacock_smith:00,scoccimarro_etal:01}; \\citealt*{berlind_weinberg:01,bullock_etal:02}; \\citealt{cooray_sheth:02}).  The halo model is a framework for calculating clustering statistics of objects by associating them with dark matter halos, which have well-studied abundances and clustering properties.  The clustering statistics of a galaxy population can be computed after specifying the clustering properties of the dark matter halos, the probability distribution for the number of galaxies in a host halo as a function of halo mass, and the distribution of these galaxies within their host halos.  In principle, the halo bias, the probability distribution for the number of galaxies in a halo of fixed mass, and the spatial distribution of galaxies within their host halos can depend on other properties of the halo, but the standard assumption is that they depend only on halo mass. \\citet{abbas_sheth:05} have recently described a modification of the halo model that incorporates dependencies on local densities.  \\citet{lemson_kauffmann:99} tested the assumption that halo properties are independent of environment using numerical simulations of cosmological structure formation and found no dependence of halo clustering on formation time or on several other properties of the dark halos.  More recent theoretical studies have also indicated that any trends of halo occupation on environment have only a relatively small net effect on large-scale clustering statistics, at least at the level that can be measured in relatively small computational volumes \\citep{berlind_etal:03,zentner_etal:05,yoo_etal:05}.  However, a recent study by \\citet{avila-reese_etal:05} found environmental trends with halo concentration, spin, shape, and internal angular momentum.  An early indication of a relationship between halo formation histories and halo clustering properties was demonstrated by \\citet[][see also \\citealt{wechsler:01}]{sheth_tormen:04}.  Recently, \\citet*{gao_etal:05} showed convincingly that the clustering of low-mass halos is a strong function of their formation times.  In this mass regime, early-forming halos are significantly more clustered than their late-forming counterparts.  \\citet{harker_etal:06} provided confirmation of these results using statistics of marked point distributions similar to those that we employ below.  The trend they identified is strongest for low-mass halos, which have only recently been well resolved in numerical simulations.  Many properties of dark matter halos correlate well with halo formation time, so it is natural to investigate whether these trends with formation time extend to other halo properties.  In particular, it is natural to expect trends with the concentrations of dark halo density profiles, which \\citet*{navarro_etal:97}, \\citet{wechsler_etal:02}, and \\citet{zhao_etal:03} have shown to correlate well with halo formation time.  In addition, several studies have convincingly demonstrated that the number of satellite halos within a host halo of fixed mass is a function of halo formation time \\citep[][]{gao_etal:04,zentner_etal:05,vandenbosch_etal:05,taylor_babul:05} and halo concentration \\citep[][]{zentner_etal:05}.  If satellite halos are to be associated with satellite galaxies in groups and clusters, this indicates that the probability distribution for the number of galaxies per halo, known as the halo occupation distribution (HOD), is also a function of these variables and, by extension, may likely be a function of halo environment.  Given the correlations between these halo properties and formation time, it is interesting to determine whether or not they relate to environment in a similar way.  This does not have to be the case; if the relations between these halo properties and formation time are themselves a function of environment, their trends with clustering could in principle be quite different.  Moreover, as we discuss below, concentration and halo occupation have much more direct consequences for the halo model and its application to constraints on cosmological parameters, and may have a more direct impact on tests of galaxy clustering.  We focus our study on these halo properties.  In the present study we use two large, high-resolution dissipationless cosmological simulations to study the dependence of the clustering of dark matter halos on halo properties other than mass.  We show that halo clustering depends on halo formation time, and present a clear demonstration that halo clustering is a function of both halo concentration and halo occupation number.  We show how these properties change with halo mass, and present the first investigation into how they change with redshift.  We present a simple fitting formula for our concentration-dependent clustering results that will enable estimates of the strength of these effects for various applications in the context of the halo model.  We discuss several implications of these results for outstanding issues in galaxy formation, including the clustering of dwarf galaxies and the concentrations of low surface-brightness galaxies, and for the estimation of cosmological parameters, including self-calibration of cluster masses.  We begin with a description of our methods in \\S~\\ref{section:methods}.  Specifically, we describe our numerical simulations in \\S~\\ref{sub:sim}, and the statistics that we employ in \\S~\\ref{sub:mcf}.  In \\S~\\ref{sub:cvir} and \\S~\\ref{sub:ac} we discuss our definitions and measurements of halo concentration and halo formation time respectively.  In \\S~\\ref{section:results}, we explore the clustering dependence of both halo formation time and halo concentration, and present a model for the relative bias of halos as a function of concentration.  In this section we also update previous results on the correlations between both halo formation time and halo concentration and the number of satellite halos contained within a host halo of fixed mass.  Following this, we give the first explicit demonstration that host halo clustering is a function of the occupation number of satellite halos.  In \\S~\\ref{section:implications}, we discuss the implications of our results for the halo model.  We conclude with a summary of our primary results and a discussion of their implications for galaxy formation models and for cosmological constraints derived from galaxy clustering in \\S~\\ref{section:disc}. ",
        "conclusions": "\\label{section:disc}  We have investigated the clustering of dark matter halos as a function of several internal halo properties, namely formation time, concentration, and occupation number.  We have confirmed that halo clustering is a function of halo formation time and have shown that the effect is scale-dependent.  We have also demonstrated that halo clustering is a strong function of halo concentration, and that the strength {\\em and sign} of this trend is a function of mass.  Of relevance to studies of galaxy clustering statistics, we also find the clearest indication yet that host halos of fixed mass cluster in a way that is dependent upon the number of subhalos that reside in them.  Our primary results can be summarized as follows. \\begin{enumerate}  \\item Halo clustering is a strong function of formation time for fixed mass halos.  This effect strengthens with decreasing halo mass and with decreasing separations, and is an increasingly strong function of mass as halos become less massive than $\\mstar$.  These results are in broad agreement with the recent results of \\citet{sheth_tormen:04}, \\citet{gao_etal:05} and \\citet{harker_etal:06}.  \\item We have presented the first definitive measurement showing that the clustering of dark matter halos is a function of halo concentration.  This effect is a strong function of halo mass, and can be characterized over a range of mass and redshift as a function of halo mass scaled by the typical collapsing mass, $\\Mvir/\\mstar$. Below $\\Mvir/\\mstar$, halos of high concentration are more clustered than halos of low concentration and this trend strengthens with decreasing halo mass.  For halos more massive than $\\mstar$, the trend changes sign, and halos of low concentration become more strongly clustered than their high concentration counterparts.  \\item The dependence of halo bias on concentration, mass, and redshift can be parameterized in a simple way: $b(M,c|z) = b(M|z) \\brel (c|M/M_*)$.  We provide a fitting function for $\\brel(c|M/M_*)$ that can be used to estimate the importance of these effects in various regimes.  In \\S~\\ref{section:implications}, we demonstrate how this relative bias can be incorporated into a halo model formalism and discuss the effects of concentration- and formation time-dependent bias on estimates of matter and galaxy correlation statistics.  \\item We confirm and update earlier results \\citep[e.g.,][]{gao_etal:04,zentner_etal:05,vandenbosch_etal:05,taylor_babul:05} that the occupation number of satellite halos is strongly correlated with both concentration and formation time. We present the first detection of a trend between clustering and halo occupation, showing that at high mass, high-occupation (late-forming) halos are more clustered than their low-$N$ counterparts. \\end{enumerate}  \\citet{sheth_tormen:04}, \\citet{gao_etal:05} and \\citet{harker_etal:06} have emphasized that the trend with formation time indicates that using the halo model to estimate clustering can be problematic.  We have quantified this explicitly, by investigating how these trends extend to two variables that are directly relevant to such calculations, namely concentration and halo occupation.  There is a weak indication that the trends we see with concentration and halo occupation are slightly stronger than would be predicted simply by their dependence on formation time and the trends with formation time itself, however, this may be do to the larger measurement error in formation time.  Larger simulations will be required to determine to higher accuracy whether the strength and nature of the trends with concentration and occupation have features that are not represented by the global correlations between these variables and formation time.  As we emphasized in \\S~\\ref{section:implications}, the dependence of clustering on concentration implies corrections to halo model calculations of the {\\em dark matter} power spectrum, as well as corrections to halo model calculations of galaxy correlations.  These results prescribe caution when using the standard halo model assumptions to calculate galaxy clustering statistics and infer cosmological parameters precisely, but it is not clear that they will have a large effect for samples that are selected by mass or, as in observational samples, by luminosity.  This is especially true for galaxies around $L_*$.  For example, \\citet{yoo_etal:05} and \\citet{zentner_etal:05} have both shown that shuffling the host halos of such a sample results in less than about $5 \\%$ effects in clustering statistics, albeit in relatively small volumes.  We expect that the trends we have shown here will have stronger effects on high- or low-mass samples (when compared to $\\mstar$) that are selected on some property that is directly connected to either formation time, concentration, or the number of satellite galaxies, and we speculate on a few of these below.  The fact that the clustering of low-mass halos is strongly correlated with formation time may have interesting implications for the so-called ``void phenomenon'', the tendency of low-mass galaxies to avoid the voids defined by larger galaxies.  The currently-favored \\LCDM\\ model predicts substantial {\\em mass} in voids, so the absence of galaxies in these regions has been suggested to be a potentially serious problem for the prevailing paradigm \\citep{peebles:01}. Whether this is indeed a problem for these models isn't clear; for example, \\citet{mathis_white:02} have investigated the clustering of low luminosity galaxies in LCDM simulations combined with a galaxy formation model and found that all galaxies avoided the voids defined by the brighter galaxies.  \\citet{benson_etal:03} found a similar result but emphasized that more data was needed to fully test the issue.  Recently, \\citet{furlanetto_piran:06} compared predictions for void sizes based on the excursion-set formalism with observations from SDSS, and saw indications that the observed voids are somewhat bigger than the model predicts.  A related piece of evidence for differential clustering of low-luminosity galaxies has been emphasized by \\citet{tully_etal:02}, namely, the difference between the luminosity function in clusters, which rises steeply toward low-luminosity dwarf galaxies, and the luminosity function in voids, which is substantially shallower and appears to be entirely bereft of a dwarf galaxy population.  It has been suggested \\citep*[][]{bullock_etal:00} that the discrepancy between the abundance of dwarf satellites observed in the Local Group and the abundance of relevant dwarf mass dark halos ($v_{\\rm max} \\lesssim 50\\ {\\rm km\\, s^{-1}}$) expected by CDM can be resolved by suppressing galaxy formation in halos that form after the universe is re-ionized.  This would bias luminous dwarf galaxies to be associated with late-forming, small dark matter halos.  As we have shown, the clustering trend with formation time is quite strong in low-mass halos of this type.  If the photoionization significantly affects galaxy formation in dwarf-size galaxy halos, the dwarf galaxies they host would be substantially more clustered than typical halos in the same mass range, and this trend can be enhanced by additional environmental factors \\citep*{kravtsov_etal:04b}. If we extrapolate the results of Eq.~(\\ref{eq:brel}) to $M \\sim 10^{-3}\\, \\mstar$, this implies that the correlation length for the earliest forming quartile should be about a factor of $4$ higher than that of the typical halos of that mass.  This may provide a natural explanation for the lack of observed dwarf galaxies in voids.  More generally, any astrophysical effect that biases small galaxies to lie in early-forming halos would produce the same effect.  There is observational evidence that star formation timescales are quite long ($\\sim 10$Gyr) in small galaxies and relatively short ($\\sim 1$Gyr) in large galaxies \\citep[e.g.,][Weiner et al., in preparation]{searle_etal:73,juneau_etal:05,willmer_etal:05}.  One implication of this is that the total stellar mass accumulated in low-mass halos will be much more sensitive to halo formation time compared to high-mass halos.  Such a formation-time bias in the observed properties of low-mass galaxies may not only help to explain the void phenomenon, but will likely be important in attempts to construct conditional luminosity functions that extend to low-luminosity galaxies.  The strong clustering of early-forming low-mass halos may also play a part in the observed trend that dim red galaxies are substantially more clustered than their intermediate luminosity counterparts \\citep[e.g.][]{norberg_etal:02, hogg_etal:03}, although the basic effect can be explained if the the majority of these galaxies are satellites \\citep{berlind_etal:05}.  Of course, any model which forms galaxies using formation histories from a N-body simulation will include the halo clustering effects implicitly, but the extent to which this effects galaxy clustering will depend on the efficiency of galaxy formation in low-mass halos.  There is another interesting implication of our results.  Numerous observational signatures indicate that the central densities and concentrations of low surface-brightness (LSB) galaxies are significantly lower than the standard \\LCDM\\ paradigm predicts \\citep[e.g.,][]{bosch_etal:00,debattista_sellwood:00,keeton:01, vandenbosch_swaters:01,alam_etal:02,zentner_bullock:02,zentner_bullock:03, mcgaugh_etal:03,vandenbosch_etal:03,kuhlen_etal:05,simon_etal:05}.  The dependence of clustering on halo concentration may also have implications for the interpretation of the concentrations of LSBs. LSBs have been shown to be notably less clustered than typical galaxies (\\citealt{mo_etal:94, rosenbaum_bomans:04}, but see also \\citealt{peebles:01}).  This might imply that they reside in a biased population of late-forming, low-concentration halos.  If the typical host halos for these galaxies are around $\\sim 0.01 \\mstar$, and LSB galaxies are about $60\\%$ less biased than typical galaxies, this would imply that they have concentration values that are about $1\\sigma$ below the mean.  This would reduce the tension between the predicted and observed concentrations of the halos hosting these galaxies.  Finally, our results also have implications for estimates of the cluster mass function using optically-selected cluster samples.  For a richness-selected sample of clusters, or any sample where the primary observable that clusters are selected on is correlated with formation time or halo concentration, mass estimates from ``self-calibration'', using the clustering of clusters \\citep{majumdar_mohr:04, lima_hu:05}, may bias results towards higher masses.  Note that because these effects correlate with concentration, and probably halo shape and merger history, they could affect SZ- or X-ray-selected samples as well.  Similar biases could potentially manifest in weak lensing measurements, and could lead to over-estimates of the mass-to-light ratios in methods that use clustering to constrain the HOD or the conditional luminosity function \\citep[e.g.][]{bosch_etal:03}.  They could also impact clustering-based mass estimation for other populations that live in high $M/M_*$ halos, e.g., high-redshift bright galaxies or quasars.  Further work is needed to make quantitative estimates of these effects at high masses, but this trend may be measurable in current optical cluster samples (e.g., from the SDSS, Koester et al., in preparation). The trend between halo occupation and formation time implies that clusters with a given number of galaxies will be a mix of high-concentration, early-forming, high-mass halos and low-concentration, late-forming, low-mass halos.  If one can find an observable measure that correlates with formation time (for example, the difference between the luminosity of the first and second brightest cluster galaxies, or the star formation histories of the satellite galaxies), these correlations make specific predictions.  At fixed richness or $\\Ngal$, the early-forming sample should be more massive and more concentrated (because concentration varies much more slowly with halo mass than with formation time: $c \\sim m^{0.1}$, \\citealt{bullock_etal:01}), but less clustered than expected for typical halos of that mass.  Our results indicate that the Universe is somewhat more complicated than our simplest models. However, the complication should be viewed as an opportunity rather than an obstacle, as we can potentially learn a great deal about details of galaxy formation in halos and their evolutionary histories from the trends discussed in this work.  The current and upcoming large galaxy surveys (e.g., SDSS, \\citealt{adelman_etal:06}; DEEP2, \\citealt{coil_etal:04}; DES, \\citealt{abbot_etal:05}; LSST, and SNAP, \\citealt{aldering_etal:04}) should be able to accurately evaluate such effects and test the predicted trends."
    },
    "0512/astro-ph0512620_arXiv.txt": {
        "abstract": "{V617 Sgr is a V Sagittae star -- a group of binaries thought to be the galactic counterparts of the Compact Binary Supersoft X-ray Sources -- CBSS.} {To check this hypothesis, we measured the time derivative of its orbital period.} {Observed timings of eclipse minima spanning over 30,000 orbital cycles are presented.} {We found that the orbital period evolves quite rapidly: $P/\\dot{P} = 1.1\\times10^{6}$ years. This is consistent with the idea that V617 Sgr is a wind driven accretion supersoft source. As the binary system evolves with a time-scale of about one million years, which is extremely short for a low mass evolved binary, it is likely that the system will soon end either by having its secondary completely evaporated or by the primary exploding as a supernova of type Ia.} {} ",
        "introduction": "Compact Binary Supersoft X-ray Sources (CBSS) are a class of objects that share in common a set of properties. They are luminous ($\\sim$Eddington luminosity) sources of soft (15-70 eV) X-ray photons and were initially discovered in the Magellanic Clouds by the Einstein observatory and ROSAT. The CBSS are thought to be cataclysmic binaries in which the secondary is more massive than the primary star. In this situation, when the secondary fills its Roche lobe a dynamical instability occurs and the mass transfer takes place on the thermal time-scale, which is about 10 million years for donor stars of 1--1.5 M$_{\\sun}$. This produces accretion rates 100 times larger than in normal cataclysmic variables and causes hydrostatic nuclear burning on the surface of the white dwarf (see Kahabka \\& van den Heuvel \\cite{kah} for a review).  Only two CBSS (\\object{MR Vel} and \\object{QR And}) are found in the Galaxy, where one should find about a thousand. This is presumably due to the absorption of their soft X-ray emission by the interstellar gas in the Galactic plane. V Sagittae stars (Steiner \\& Diaz \\cite{stei98}) were proposed as a new class of binaries that display properties quite similar to those of CBSS, but are not detected as supersoft sources. They may be the galactic counterpart of the CBSS. The soft photons are either absorbed by the stellar wind or by the interstellar medium (or both). In case this hypothesis is correct, these two classes should share a number of properties in common. For example, the time variation of the orbital period should be high and similar in the two situations  - and this could be a critical test for the hypothesis of the CBSS -- V~Sge connection.  What do we expect in terms of the orbital period time derivative? In the scenario of dynamical instability, we expect that the orbital period decreases with time. There is only one such object for which the period derivative has been measured: \\object{V~Sge}. Its period, in fact, decreases with a time-scale of 5 million years (Patterson et al. \\cite{patter}). However, this scenario only predicts the existence of orbital periods longer than 6 hours (Deutschmann \\cite{deutsch}; King et al. \\cite{king2}). For periods smaller than this limit, the mass transfer is too small for nuclear burning to occur. This limitation on the orbital period imposes a problem to the interpretation of the short orbital period systems among CBSS (\\object{SMC~13} and \\object{RX J0537.7-7034}) and among V~Sge stars (\\object{V617~Sgr}), which have orbital periods shorter than 5 hours.  It has been proposed that in such objects the mass transfer occurs because of radiation induced wind from the secondary (van Teeseling \\& King \\cite{tees}). If a low mass secondary star ($M_{2}\\leq0.6 M_{\\sun}$) looses mass adiabatically, that is, on a time-scale short compared to the thermal time-scale, it expands as  \\eq \\frac{\\dot{R_{2}}}{R_{2}} = \\zeta \\frac{\\dot{M_{2}}}{M_{2}} \\eeq  \\noindent where $\\zeta=-1/3$ is the effective mass-radius index of the secondary star.  In this case, the expansion forces mass transfer through the internal Lagrangean point at a high rate, producing a CBSS. In this situation, the orbital period increases with time (van Teeseling \\& King \\cite{tees}).  We have, thus, two distinct scenarios for CBSS as well as for V~Sge stars: systems with initial orbital periods longer than 6 hours evolve with negative derivative while systems with shorter periods should evolve with positive derivatives, both with high absolute values. For wind driven supersoft binaries, the orbital period increases with time and may become longer than 6 hours. The total possible orbital period range is 2--30 hours (van Teeseling \\& King \\cite{tees}). So, for a given object, how to decide to which of the two above paradigms the object belongs? In principle one could determine the mass ratio or the secondary's spectral type. But no secondary star has been detected so far in any of the CBSS or V~Sge stars. Emission lines are strongly contaminated by the wind and are not reliable dynamical indicators.  For eclipsing systems there is an alternative way of finding whether a given star belongs to one class or to the other: by measuring its orbital period time derivative. A critical test could be provided by V617~Sgr. A positive period derivative with time-scale of about a million years is predicted. This should be measurable, given the time base of our eclipse timings. A similar study was proposed for \\object{T~Pyx}, a recurrent nova with an orbital period of 1.8 hour (Knigge et al. \\cite{knigge}) .  V617~Sgr (Steiner et al. \\cite{stei99}; Cieslinski et al. \\cite{cies}) was identified as a V~Sge star (Steiner \\& Diaz \\cite{stei98}) with an orbital period of 4.98 hours. It presents a light curve with two maxima and two minima (see Fig. 3 in Steiner et al. \\cite{stei99}) of unequal depths. The system presents high and low photometric states, like V~Sge itself and most CBSS. Timings of the main (deepest) minima provide an opportunity to measure the orbital period with accuracy as well as its time derivative. ",
        "conclusions": "%  An important system parameter is the mass of the white dwarf. This is usually derived from dynamical measurements. In the case of this system, however, the observed strong wind complicates this kind of determination (Cieslinski et al. \\cite{cies}). Steiner et al. (\\cite{stei99}) considered that the mass of the white dwarf must be quite low. This, however, was based on the erroneous assumption that the system is a CBSS with a secondary more massive than the primary and that accretion is driven by the thermal expansion of the donor star.  One alternative way of estimating this mass is by measuring the time-scale of decline from outburst maxima (Southwell et al. \\cite{south}). V617~Sgr displays such outbursts as most of other well observed V~Sge and CBSS systems do. One such event can be seen in Fig. 4 of Steiner et al. (\\cite{stei99}). Eclipse cycle 4069 was observed when the system was at high state. Four days later cycle 4088 was observed at about one magnitude fainter. The decay timescale is, therefore, 4 days or shorter.  From the outburst decline we derive a high mass for the white dwarf: $M_1 = 1.2 M_{\\sun}$ or higher. If one assumes that the secondary is in the main sequence (this is not obvious, as it is not necessarily in equilibrium), $M_2 = 0.48 M_{\\sun}$ and the mass ratio is $q= 0.40$.  In the wind driven supersoft X-ray binary scenario, the orbital period time derivative is given by the formula (Knigge et al. \\cite{knigge})  \\eq \\frac{P}{\\dot{P}} = \\frac{2M_2}{(3\\zeta -1)(1+g)\\dot{M}_{w2}} \\eeq  \\noindent where  \\eq \\dot{M}_{acc}= -g \\dot{M}_{w2}% \\simeq 1.2 \\times 10^{-6} g^{2} \\phi^{2} \\eta_a \\eta_s \\left(\\frac{q^{5/2}}{1+q}\\right)^{2/3} m_1~M_{\\sun} yr^{-1} \\eeq  \\noindent and  \\eq g=\\frac{(6\\beta_2 + 2q)-(5+3\\zeta)(1+q)}{(1+q)(5+3\\zeta-6q)} \\eeq  here $q=M_2/M_1$ is the mass ratio, $\\dot{M}_{w2}$ is the wind mass loss rate from the irradiated secondary star, $\\dot{M}_{acc}$ is the accretion rate, $m_1=M_{1}/M_{\\sun}$, $\\eta_a$ measures the luminosity per gram of matter accreted relative to the value for nuclear shell burning, $\\eta_s \\sim 1$ for CBSS (the efficiency of the primary's spectrum in producing ionizing photons and driving a wind), $\\beta_2$ is the specific angular momentum loss of the secondary star, and $\\phi$ is an efficiency factor parameterizing the fraction of the companion's face which is irradiated.  In the present case the period derivative and the mass ratio are self-consistent if $\\beta_2 = 1.2$. That is the ratio of the  specific momentum of the wind relative to that of the secondary star. The observed period variation suggests that V617~Sgr is, indeed, a wind-driven supersoft X-ray source. This adds more evidence to the hypothesis that V~Sge stars are the galactic counterparts of the CBSS objects.  This binary system evolves with a time-scale of about one million years, which is extremely short for this kind of low-mass evolved binary. As the white dwarf is quite massive and if it accretes half of the mass of the secondary star it may soon reach the Chandrasekhar limit and, eventually, explode as a Supernovae type Ia. Such binary systems may occur in old populations as demanded by SN Ia statistics, contrary to thermal time-scale mass transfer systems that require relatively young populations. The other possibility is that the secondary will evaporate completely. In this situation one would expect the orbital periods to increase up to 30 hr (van Teeseling \\& King \\cite{tees}).  V617~Sgr is a system that has left completely the standard CV evolutionary track and will probably destroy itself. This is a channel that may remove this kind of system from the general CV population.  V617~Sgr is a member of a group of 3 known stars that are wind driven supersoft x-ray binaries; the other members of the group are SMC~13 and T~Pyx (van Teeseling \\& King \\cite{tees}; Knigge et al. \\cite{knigge}). It is the first system in which the predicted evolutionary trend has been clearly observed, thanks to the existence of a relatively deep eclipse in the orbital light curve."
    },
    "0512/astro-ph0512550_arXiv.txt": {
        "abstract": "We report on integral field spectrocopy observations, performed with the PPAK module of the PMAS spectrograph, covering a field-of-view of $\\sim$74''$\\times$64'' centered on the core of the galaxy cluster Abell 2218. A total of 43 objects were detected, 27 of them galaxies at the redshift of the cluster. We deblended and extracted the integrated spectra of each of the objects in the field using an adapted version of {\\tt galfit} for 3D spectroscopy ({\\tt galfit3d}). We use these spectra, in combination with morphological parameters derived from deep HST/ACS images, to study the stellar population and evolution of galaxies in the core of this cluster. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512440.txt": {
        "abstract": "{ We report new simultaneous near-infrared/sub-millimeter/X-ray observations of the SgrA* counterpart associated with the massive 3--4$\\times$10$^6$\\solm ~black hole at the Galactic Center. }{ The main aim is to investigate the physical processes responsible for the variable emission from SgrA*. }{ The observations have been carried out using the NACO adaptive optics (AO) instrument at the European Southern Observatory's Very Large Telescope\\footnote{Based on observations at the Very Large Telescope (VLT) of the European Southern Observatory (ESO) on Paranal in Chile; Program: 271.B-5019(A).} and the ACIS-I instrument aboard the \\emph{Chandra X-ray Observatory} as well as the Submillimeter Array SMA\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} on Mauna Kea, Hawaii, and the Very Large Array\\footnote{The VLA is operated by the National Radio Astronomy Observatory which is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} in New Mexico. }{ We detected one moderately bright flare event in the X-ray domain and 5 events at infrared wavelengths. The X-ray flare had an excess 2 - 8 keV luminosity of about 33$\\times$10$^{33}$~erg/s. The duration of this flare was completely covered in the infrared and it was detected as a simultaneous NIR event with a time lag of $\\le$10~minutes. For 4 flares simultaneous infrared/X-ray observations are available. All simultaneously covered flares, combined with the flare covered in 2003, indicate that the time-lag between the NIR and X-ray flare emission is very small and in agreement with a synchronous evolution. There are no simultaneous flare detections between the NIR/X-ray data and the VLA and SMA data. The excess flux densities detected in the radio and sub-millimeter domain may be linked with the flare activity observed at shorter wavelengths. }{ We find that the flaring state can be explained with a synchrotron self-Compton (SSC) model involving up-scattered sub-millimeter photons from a compact source component. This model allows for NIR flux density contributions from both the synchrotron and SSC mechanisms. Indications for an exponential cutoff of the NIR/MIR synchrotron spectrum allow for a straight forward explanation of the variable and red spectral indices of NIR flares. } ",
        "introduction": "\\label{section:Introduction}  Over the last decades, evidence has been accumulating that most quiet galaxies harbor a massive black hole (MBH) at their centers. Especially in the case of the center of our Galaxy, progress could be made through the investigation of the stellar dynamics (Eckart \\& Genzel 1996, Genzel et al. 1997, 2000, Ghez et al. 1998, 2000, 2003a, 2003b, 2005, Eckart et al. 2002, Sch\\\"odel et al. 2002, 2003, Eisenhauer 2003, 2005). At a distance of only $\\sim$8 kpc from the sun (Reid 1993, Eisenhauer et al. 2003, 2005), the Galactic Center allows for detailed observations of stars at distances much less than 1~pc from the central black hole candidate, the compact radio source Sgr~A*.  Additional compelling evidence for a massive black hole at the position of Sgr~A* is provided by the observation of variable emission from that position both in the X-ray and recently in the near-infrared wavelength domain (Baganoff et al. 2001, 2002, 2003, Eckart et al. 2003, 2004, Porquet et al. 2003, Goldwurm et al. 2003, Genzel et al. 2003, Ghez et al. 2004a, and Eisenhauer et al. 2005). Throughout the paper we will use the term 'interim-quiescent' (or IQ) for the phases of low-level, and especially in the NIR domain - possibly continously variable flux density at any given observational epoch. This state may represent flux density variations on longer time scales (days to years). This is especially true for the NIR source (Genzel et al. 2003, Ghez et al. 2004a, Eckart et al. 2004).  Simultaneous observations of SgrA* across different wavelength regimes are of high value, since they provide information on the emission mechanisms responsible for the radiation from the immediate vicinity of the central black hole.  The first observation of SgrA* detecting an X-ray flare simultaneously in the near-infrared was presented by Eckart et al. (2004). They detected a weak 6$\\times$10$^{33}$~erg/s X-ray flare and covered its decaying flank simultaneously in the NIR.  Variability at radio through submillimeter wavelengths has been studied extensively, showing that variations occur on time scales from hours to years (Wright \\& Backer 1994, Bower et al. 2002, Herrnstein et al. 2004, Zhao et al. 2003).  Some of the variability may be due to interstellar scintillation.  The connection to variability at NIR and X-ray wavelengths has not been clearly elucidated.  Zhao et al. (2004) showed a probable link between the brightest X-ray flare ever observed and flux density at 0.7, 1.3, and 2 cm wavelength on a timescale of $<1$ day (see also Mauerhan et al. 2005).  In section 2 of the present paper we report on new simultaneous NIR/X-ray observations using \\emph{Chandra} and the adaptive optics instrument NACO at the VLT UT4. The new 8.6$\\mu$m and 19.5$\\mu$m observations of the central region were obtained during the commissioning of the ESO MIR VISIR camera. We also describe the new SMA and VLA data of SgrA*. These data give additional information on the flux density limit of SgrA* at millimeter and sub-millimeter wavelengths. In section 3 we discuss the NIR to X-ray variability of SgrA*, followed by a discussion of its MIR/NIR spectrum in section 4. In sections 5 and 6 we discuss the flux densities, spectral indices and flares observed in the NIR and X-ray domain. The physical interpretation in section 7 is then followed by the summary and discussion in the final section 8.    %__________________________________________________________________ ",
        "conclusions": "We have presented new, successful simultaneous X-ray and NIR observations of SgrA* in a flaring and the IQ low NIR flux  density state. We found 4 X-ray flares (2 definite flares and 2 putative events) and 5 NIR flares with  4 events covered simultaneously at both wavelengths. For the flares we observed simultaneously in both wavelength domains, the time lag between the flares at different wavelengths is less than 10 minutes and therefore consistent with zero. Combined with the information that the NIR flare spectra are very red with variable spectral indices (Eisenhauer et al. 2005, Ghez et al. private communication) we can successfully describe the flares by a SSC model in which a substantial fraction of the NIR emission is due to a truncated synchrotron spectrum. Inverse Compton scattering of the THz-peaked flare spectrum by the relativistic electrons accounts for the X-ray emission.  Our investigation also shows that the NIR K-band is the ideal wavelength band to study the flare emission from SgrA*. In combination with adaptive optics systems it provides the highest angular resolution at the lowest amount of contamination by dust emission. At wavelengths shorter than the K-band little emission is found because the flares are red (Eisenhauer et al. 2005) and at longer wavelength the angular resolution is lower and the dust contamination is high. Observations in the K-band allow us to measure the highest flare rate and are - in the framework of the presented physical model - ideally suited to observe both synchrotron and SSC flare emission. In addition the model also gives us the opportunity to perform polarization measurements which could provide additional information to study the relevant emission mechanisms.   The total number $\\Sigma$ of detectable flares can be obtained by integrating over the amplitude dependent flare rate $N(A)=\\kappa_0 (A)^{-\\zeta} \\kappa_1  \\kappa_2$ (see section \\ref{section:Flares}) as $\\Sigma=\\int_{A_{limit}}^{\\infty}N(A)dA$, with $A_{limit}$ being the detection limit of the flare emission. The model presented in section~\\ref{section:Interpretation} suggests that in the NIR domain the observed flares can be produced by a mixture of synchrotron and SSC emission, i.e. $\\Sigma_{NIR}=\\Sigma_{Synch,SSC}$. Since we can assume that the X-ray flares are predominantly produced by SSC emission rather than synchrotron emission, as also suggested by the very steep NIR flare spectra (Eisenhauer et al. 2005), it follows that $\\Sigma_{X-ray}\\sim\\Sigma_{SSC}$. As a consequence - and in good agreement with the observations - the total number of detected X-ray flares is smaller than that in the NIR $\\Sigma_{X-ray} \\le \\Sigma_{NIR}$. This does, however, not imply that the flux density distribution of flares dominated by SSC emission in both wavelength domains are the same. That distribution depends on the properties of the relativistic electron spectrum responsible for the emission at both wavelength regimes. These are reflected in parameters like the spectral index of the optically thin radio continuum and the exact location of the high and low energy cutoff frequencies of the scattered SSC spectrum. In addition, NIR flares may have contributions from both the synchrotron and the SSC part the flare spectrum and it may be difficult to discriminate between the SSC and synchrotron dominated flare activity. One can, however, expect that SSC dominated NIR flares are bluer than synchrotron dominated ones.  The description of the flare activity as a power-law under the assumption of a characteristic flare time implies that the IQ phase can be regarded as a sequence of frequent low amplitude flares of SgrA*. Such a model would predict phases of very low flux densities (see also Eckart et al. 2004 - Garching). In Fig. \\ref{Fig:nirsim} we show a simulation with the appropriate power-law spectral index. The observed flares may be the consequence of a clumpy or turbulent accretion. Evidence of a hot turbulent accretion flow onto SgrA* based on polarization measurements has been discussed by Bower et al. (2005). In this case the flare power spectrum is coupled to the power spectrum of accreted clumps or the turbulences in the accretion flow.   The red source component we identified close to the position of SgrA* at 3.8$\\mu$m, 8.6$\\mu$m, and 19.5$\\mu$m is probably contaminated significantly by thermal emission from a dust component along the line of sight towards SgrA* (Fig.~\\ref{Fig:dustblob} and \\ref{Fig:dustblobvisir}). The infrared flux density ratios of the emission from that region compared to values obtained from the mini-spiral and other discrete sources in the central parsec suggest that the emission is due to dust. Assuming that the gas and dust properties of this component are similar to the material in the northern arm we can obtain a first order estimate of its mass, which can be thought of as a structure which is thin with respect to its projected extent (e.g. Vollmer \\& Duschl 2000). Based on CO(7-6) measurements Stacey et al. (2004) derive a total gas mass of the northern arm of 5 to 50 \\solm. In projected size the dust component close to SgrA* covers about 1/250 of the areas comprised by the northern arm. This results in a gas mass of the order of 10$^{-2}$\\solm (If the dust temperature is substantially higher than $\\sim$200-400~K which are typical of the mini-spiral - see Cotera et al. 1999 - then the overall mass of this component can be considerably smaller; see Ghez et al. 2005).  Depending of the clumpiness of the gas distribution within that component on the source size scale of SgrA* this may result in a significant column density. The dust source is, however, most likely located behind SgrA*. If it were located in front of SgrA* the high velocity stars in the central cusp would also be affected. However, for other sources in the field, like S2, S12, and S14, their (A$_V$$\\sim$25$^m$) extinction corrected spectra are blue, and variable flux densities and colors have not been detected within the uncertainties of a few 0.1 magnitudes. The variable NIR spectral indices reported for the red flares (Eisenhauer et al. 2005, Ghez et al. private communication) suggest that source intrinsic emission processes are responsible for the NIR spectral shape rather than extrinsic processes like extinction. The fact that the dust source is most likely located behind SgrA* also suggests that is associated with the northern arm section of the mini-spiral which is assumed to approach the central stellar cluster from behind the plane of the sky in which SgrA* is located (Vollmer \\& Duschl 2000).  Finally our investigation shows that most of the MIR flux density seen towards the position of SgrA* is due to dust emission. This suggests that the overall spectral shape of SgrA* is significantly less peaked in the FIR wavelength domain as suggested by the the upper limits. Combined with the results from our SSC modeling we find that one can expect that the intrinsic spectrum of SgrA* is peaked at frequencies of a few THz. The radio and submillimeter data show clear indications for variability. In general the flux density variations are slow and occur on somewhat longer timescales than the X-ray and IR variations. The exact relation between the radio/sub-mm domain and the NIR/X-ray domain still remains uncertain, due to the lack of sufficient simultaneous coverage. However, the amplitudes and time scales indicated are consistent with a model in which the emitting material is expanding and cooling adiabatically.  Future observations will lead to improved statistics on the differences between simultaneous NIR and X-ray flares. Especially the coupling to the mm-domain is of importance. Here, no simultaneous data are available so far. Such observations will help to investigate whether individual mm-flare events are related to events in the NIR or X-ray regime. Upcoming simultaneous monitoring programs from the radio to the X-ray regime will be required to further investigate the physical processes that give rise to the observed IQ low NIR flux density state and flare phenomena associated with SgrA* at the position of the massive black hole at the center of the Milky Way.    %______________________________________________________________"
    },
    "0512/astro-ph0512190_arXiv.txt": {
        "abstract": "{We present mid-infrared (2--12\\,$\\mu$m) spectra of the microquasar SS\\,433 obtained with the Infrared Space Observatory (spectroscopic mode of ISOPHOT and ISOCAM). We compare them to the spectra of four Wolf-Rayet stars: WR\\,78, WR\\,134, WR\\,136, and WR\\,147 in the same wavelength range. The mid-infrared spectrum of SS\\,433 mainly shows \\ion{H}{i} and \\ion{He}{i} emission lines and is very similar to the spectrum of WR\\,147, a WN8(h)+B0.5V binary. The 2--12\\,$\\mu$m continuum emission of SS\\,433 corresponds to optically thin and partially optically thick free-free emission, from which we calculate a mass loss rate of $2-3\\times 10^{-4} M_\\odot \\,\\mathrm{yr}^{-1}$ if the wind is homogeneous and a third of these values if it is clumped. This is consistent with a strong stellar wind from a WN star. However, following recent studies concluding that the mass donor star of SS\\,433 is not a Wolf-Rayet star, we propose that this strong wind out flows from a geometrically thick envelope of material that surrounds the compact object like a stellar atmosphere, imitating the Wolf-Rayet phenomenon. This wind could also wrap the mass donor star, and at larger distances ($\\sim40$\\,AU), it might form a dust envelope from which the thermal emission, detected with ISOPHOT at 25\\,$\\mu$m and 60\\,$\\mu$m, would originate. This wind also probably feeds the material that is ejected in the orbital plane of the binary system and that forms the equatorial outflow detected in radio at distances $>100$\\,AU. ",
        "introduction": "SS\\,433 was the first microquasar (an X-ray binary with relativistic jets) discovered in the 1970's, when that designation had not even been invented yet. It was first discovered as a star with strong H$\\alpha$ emission lines and so was included in the Stephenson and Sanduleak catalog \\citep{stephsand77} as object number 433. It was later associated to a variable X-ray and radio source \\citep{zwitter89}.  Detailed study of its visible spectrum revealed that this source is unique. For a complete review of SS\\,433, we refer the reader to \\citet{margon84}, \\citet{zwitter89}, \\citet{verm96}, and to the introduction of \\citet{giesmcswain02}.  There are two sets of optical emission lines in the spectrum of SS\\,433: the first set corresponds to the so-called ``stationary'' lines \\citep{margon79} showing normal Doppler shift movements, including the strong H$\\alpha$ lines, with a period of 13~days. The second set regroups the ``moving'' lines, those that are less intense (by about 1/3) but that show huge Doppler shifts ($+50\\,000\\,$Km.s$^{-1}$ and $-35\\,000\\,$Km.s$^{-1}$ at the maximum elongation) corresponding to relativistic velocities with a period of about 162~days \\citep{margon79}. The ``kinematic model'' is generally the one chosen to explain these unusual observations: SS\\,433 is an X-ray binary system orbiting in 13~days and emitting relativistic jets where the moving lines are formed. These jets undergo a precession movement in 162 days. The parameters of this system were recently re-calculated by \\citet{eiken01}, who found: $P_\\mathrm{orb}=13.08$\\,d, $P_\\mathrm{prec}=162.375\\pm0.011$\\,d, velocity of the ejections $v=0.2647\\pm0.0008$\\,c, inclination of the jet axis to the line of sight $i=78.05\\degr\\pm0.5\\degr$, and opening angle of the precession cone $\\theta=20.93\\degr\\pm0.08\\degr$. The ejections and their precession movement are indeed observed in the radio images; see e.g. \\citet{hjellming81} and \\citet{verm93}. The source has always been observed ejecting material, meaning that this is not only the first microquasar discovered but also the only one with continuous ejections, and the only object known to show evidence of \\emph{ions} accelerated to relativistic velocities (0.26\\,c).   Despite intensive studies, the nature of the two stars in this binary system remains uncertain. No emission line  has been identified as clearly belonging to the donor star of this X-ray binary, some of them being suspected to be formed close to the compact object. Numerous mass ratios were estimated leading to a low-mass or high-mass X-ray binary with either a neutron star or a black hole candidate. The presence of Wolf-Rayet-like lines added to the very luminous continuum in the visible and near-IR ranges led several people to associate the donor star to a Wolf-Rayet or Of star \\citep{murdin80,vandenheuvel80,hut81}. However, the coordinated optical and X-ray light curves of SS\\,433 show that the X-ray eclipse corresponds to the minimum visual magnitude \\citep{stewart87}; thus a thick accretion disc is inferred, which is more luminous in the visible range than the donor star. Only recently, \\citet{gieshuang02} detected absorption features in the blue spectrum  suggesting that the donor star is an A-type evolved star. In the more detailed spectroscopic study of \\citet{hillwig04}, the donor star is consistent with an \\mbox{A3-7\\,I} type with $M=10.9\\pm3.1\\,M_\\odot$, and the compact object is a low-mass black hole candidate with $M_\\mathrm{X}=2.9\\pm0.7\\,M_\\odot$.   The distance of SS\\,433 has also remained uncertain for a long time. From complete observations of the radio emitting materials ejected in the relativistic jets and after taking  the kinematic model into account, \\citet{blundell04} eventually determined this distance as $d=5.5\\pm0.2$\\,kpc.  As the origin of the spectrum (lines and continuum) of SS\\,433 in the visible and near-infrared ranges appears quite complicated for understanding clearly, the idea was to look for characteristic lines in the mid-IR range which could help to understand the nature and the origin of the emissions better and possibly to constrain the nature of the donor star further. In this article we present observations of SS\\,433 with the Infrared Space Observatory (ISO) in both spectroscopic and photometric modes in the 2--12\\,$\\mu$m range at different epochs and at 25\\,$\\mu$m \\& 60\\,$\\mu$m. We compare the emission line spectra to the spectra of four Wolf-Rayet star of WN type (Sect.~3). We also study the continuum emission of SS\\,433 in Sect.~4 and calculate the mass loss and radius of the corresponding free-free emission. In Sect.~5 we discuss the constraints on the nature of the mass donor star provided by our mid-IR observations, we describe the phenomenon imitating a Wolf-Rayet star, and we consider the possible large scales behaviour of the strong mass outflow. ",
        "conclusions": "The mid-IR spectra of SS\\,433 obtained with ISO in 1996 and 1997 show H and He emission lines similar to the ones observed in the spectrum of the Wolf-Rayet star WR\\,147, a WN8(h)+B0.5V binary system with colliding wind. The spectrum of SS\\,433 is thus compatible with the presence of a WN8 star.  The 2--12\\,$\\mu$m continuum emission of SS\\,433 corresponds to free-free emission, optically thin or intermediate between optically thin and thick, depending on both the wavelength and time of observation. At 25\\,$\\mu$m and 60\\,$\\mu$m, the emission may be due to dust at $T\\sim150K$ surrounding the binary system at a large distance ($\\gtrsim 8000\\,R_\\odot$). Assuming that the free-free emission is emitted by a geometrically thick homogeneous wind, we calculated the corresponding mass loss of $\\sim 2-3 \\times 10^{-4} \\ M_\\odot \\,\\mathrm{yr}^{-1}$. If the wind is clumped, this result is a factor of 3 times lower: $\\sim 6-10 \\times 10^{-5} \\ M_\\odot \\,\\mathrm{yr}^{-1}$, thus compatible with a strong wind from a WN star.  However, considering recent results discarding a WR star for the nature of the mass donor star in SS\\,433, we propose that the WR-like wind observed in the IR is outflowing from an envelope of material enshrouding the region of the compact object and maybe also the donor star. This envelope, heated, ionized, and expelled by the X-ray emission of the compact object, is thus imitating a Wolf-Rayet star. This wind is probably the source that provides material which forms the possible dust emitting in the far-IR and, at larger distance ($>100$\\,AU), the equatorial outflow observed in radio.\\\\"
    },
    "0512/astro-ph0512159_arXiv.txt": {
        "abstract": "Large scale structure introduces two different kinds of errors in the luminosity distance estimates from standardizable candles such as supernovae Ia (SNe) -- a Poissonian scatter for each SN and a coherent component due to correlated fluctuations between different SNe. Increasing the number of SNe helps reduce the first type of error but not the second. The coherent component has been largely ignored in forecasts of dark energy parameter estimation from upcoming SN surveys. For instance it is commonly thought, based on Poissonian considerations, that peculiar motion is unimportant, even for a low redshift SN survey such as the Nearby Supernova Factory (SNfactory; $z = 0.03 - 0.08$), which provides a useful anchor for future high redshift surveys by determining the SN zero-point. We show that ignoring coherent peculiar motion leads to an underestimate of the zero-point error by about a factor of $2$, despite the fact that SNfactory covers almost half of the sky. More generally, there are four types of fluctuations: peculiar motion, gravitational lensing, gravitational redshift and what is akin to the integrated Sachs-Wolfe effect. Peculiar motion and lensing dominates at low and high redshifts respectively. Taking into account all significant luminosity distance fluctuations due to large scale structure leads to a degradation of up to $60 \\%$ in the determination of the dark energy equation of state from upcoming high redshift SN surveys, when used in conjunction with a low redshift anchor such as the SNfactory. The most relevant fluctuations are the coherent ones due to peculiar motion and the Poissonian ones due to lensing, with peculiar motion playing the dominant role. We also discuss to what extent the noise here can be viewed as a useful signal, and whether corrections can be made to reduce the degradation. ",
        "introduction": "\\label{intro}  The problem of the cosmological constant, or more generally dark energy, is one of the deepest problems in cosmology today. While there are by now multiple lines of evidence for the existence of dark energy \\cite{evidence}, the evidence from type Ia supernovae (SNe) was historically what convinced a large fraction of the cosmology community that this enigmatic form of energy should be taken seriously \\cite{Riess:1998,Perlmutter:1999}. Upcoming and ongoing SN surveys \\cite{SNsurveys}, with vastly improved statistics, promise to constrain the equation of state of dark energy to unprecedented precision, thus shedding light on the issue of whether the apparent acceleration of the universe is caused by the cosmological constant, a dynamical scalar field or departure from Einstein gravity \\cite{gravity}.  There has been much recent work on projections for the determination of dark energy properties from these SN surveys. By and large, they focus on the following aspects of the error budget: intrinsic statistical error, systematic error and gravitational lensing induced scatter (e.g. \\cite{josh0,hutererturner,josh,linder,kim,hutererkim,hl} and references therein). The intrinsic statistical error refers to the intrinsic spread in SN luminosity even after suitable standardizing corrections have been applied. It is typical to assume that the intrinsic spread in magnitude has a (root-mean-squared; rms) size of $\\sigma^{\\rm intr.} = 0.1 - 0.15$ for each SN \\cite{sigintr}. This kind of intrinsic statistical error can be beaten down by having a large number of SNe. There are several sources of systematic error, such as Malmquist bias, luminosity evolution, imperfect corrections for dust extinction, and so on. They are not necessarily diminished by having a large number of SNe, although a large sample often helps in identifying and characterizing them. Lastly, gravitational lensing by intervening structures introduces fluctuations in the observed flux of SNe. So far, the focus has been on how gravitational lensing introduces a Poissonian scatter rather analogous to the intrinsic spread. This kind of error can likewise be reduced by having a large sample of SNe \\cite{josh0,hl}.  The existing discussion can be improved in two ways. First of all, gravitational lensing by large scale structure introduces not only a Poissonian scatter to the individual SN flux, but also correlated flux fluctuations between different SNe. One can view the correlated fluctuations as a consequence of the large scale coherence of the intervening structures. Second, large scale structure introduces fluctuations beyond that captured by gravitational lensing, and like lensing, these fluctuations have a Poissonian component as well as a correlated or coherent component. It is worth noting that an expression for all the first order fluctuations in the luminosity distance -- first order in metric and energy-momentum perturbations -- has been worked out for quite some time e.g. \\cite{sasaki,pyne1,pyne2} (with minor corrections; see below). The full implications for current and future SN surveys, however, have not been explored, with an important exception (pointed out to us by Dragan Huterer) -- Sugiura, Sugiyama \\& Sasaki \\cite{sss99} computed the anisotropies (the dipole and beyond) in luminosity distance and investigated the implications for measurements of the decceleration parameter $q_0$.  As we will see, to first order, there are four sources of luminosity distance (or magnitude) fluctuations: gravitational lensing, gravitational redshift, peculiar motion and an effect akin to the integrated Sachs-Wolfe (ISW) effect. We will see that for most practical purposes, it is sufficient to consider gravitational lensing and peculiar motion. They become important at high ($z \\gsim 1$) and low ($z \\lsim 0.1$) redshifts respectively.  What is particularly interesting, and perhaps surprising, is that peculiar motion plays a significant role in the degradation of dark energy errors. There is a widespread perception that the effects of peculiar motion are negligible as long as the median redshift is greater than $0.05$ or so. Let us take as an example the Nearby Supernova Factory (SNfactory; other low redshift surveys include the CfA Supernova Program, Carnegie Supernova Project and LOTOSS, see \\cite{SNsurveys}), whose redshift range $z = 0.03 - 0.08$ was chosen in the hope of making the effects of peculiar motion negligible. Such a perception seems at first sight quite reasonable: typical peculiar velocities are of the order of $300$ km/s, and so the ratio of peculiar flow to Hubble flow at $z = 0.055$ is about $300 / (3 \\times 10^5 \\times 0.055) \\sim 0.02$. Translating this into fluctuations in magnitude (details are given in \\S \\ref{dLf}), we have $\\delta m \\sim 2.17 \\times 0.02 \\sim 0.04$, which is quite a bit smaller than the intrinsic spread in SN magnitude ($0.1 - 0.15$), apparently suggesting we can ignore peculiar motion (recall that different sources of errors are to be added in quadrature).  What such an argument misses is that coherent peculiar flows introduce correlations in magnitude fluctuations between different SNe. While it is true that peculiar motion introduces a negligible Poissonian scatter compared to the intrinsic scatter, the correlated component cannot be ignored as it turns out. One can intuitively understand it as follows. As the number of SNe ($N$) becomes large, the intrinsic statistical error is beaten down to be quite small in the usual root-$N$ fashion. Correlated errors, such as that due to correlated/coherent peculiar flows, are not reduced by $N$ at all, and so there must be some $N$ beyond which the correlated errors become dominant. We will see that this is indeed the case for the SNfactory. Since a low redshift survey such as the SNfactory plays an important role in constraining the SN zero-point, dark energy determination from higher redshift surveys (where peculiar motion is less of an issue) is affected indirectly by these considerations as well.  Coherent large scale flows (i.e. bulk flows) have of course been the subject of research for a long time (see \\cite{strauss} for a review). Particularly relevant to our investigation are discussions of the peculiar velocity monopole, or what is sometimes referred to as the local Hubble bubble, which incidentally made use of SNe Ia \\cite{shi,shiturner,idit}. We will see later that for a survey like the SNfactory, which covers roughly half of the sky, fluctuations in the lower velocity multipoles (the monopole, dipole, etc) contribute significantly to the dark energy error budget.  The rest of the paper is organized as follows. In \\S \\ref{prelim}, we set the stage by describing the parameters of interest and how the fluctuations in magnitude are related to the parameter errors. Some details on how to go from the Fisher matrix to actual errorbars of various types are given in Appendix \\ref{app:fisher}. We describe in \\S \\ref{flux} how the average magnitude in a given redshift bin fluctuates, and how these fluctuations can be divided into a Poissonian component and a correlated/coherent component, in effect defining the magnitude covariance matrix. To keep the discussion simple, the derivation is relegated to Appendix \\ref{app:poisson}. In \\S \\ref{frw}, we derive an expression for the first order luminosity distance fluctuations (\\S \\ref{dLf}), and work out explicitly their implications for the magnitude covariance matrix in terms of the mass power spectrum (\\S \\ref{dLf2}). To keep the discussion simple in the main body, details of these two steps are relegated to Appendices \\ref{app:dL} and \\ref{app:poisson2}. Appendix \\ref{app:dL} might be interesting to the more theoretically inclined: the explicit expression given here for the luminosity distance fluctuation corrects a minor error in earlier expressions in the literature. Appendix \\ref{app:poisson2} contains expressions for the velocity window function for an arbitrary survey geometry that might be of interest to observers who are interested in making predictions for their own surveys. In \\S \\ref{forecasts}, we finally put everything together to make error forecasts. It proves illuminating to first focus separately on the contributions to errors from peculiar motion and lensing (\\S \\ref{vel} and \\ref{lens}), and in \\S \\ref{everything} we make forecasts for a number of ongoing/planned/proposed SN surveys. The key results are summarized in Fig. \\ref{fig:testhistoB}, \\ref{fig:testhistoBnoprior}, \\ref{fig:contour.1.combo} and Table \\ref{table}. We conclude in \\ref{discuss} with a brief summary of major results and a discussion of several issues that naturally arise, some of which are worth exploring further: \\begin{itemize} \\item whether peculiar motion degrades the {\\it current} constraints on dark energy (the answer is: not significantly); \\item how the exact survey geometry impacts dark energy errors; \\item whether internal motion that could add to the peculiar velocity is important (the answer is no); \\item the issue of systematic errors, and how they might change our conclusions; \\item whether realistic redshift measurements are accurate enough for us to have to worry about peculiar velocities (the answer is yes); \\item whether corrections can be made for peculiar motion and lensing to reduce the dark energy errors (the answer is: probably difficult for (the Poissonian part of) lensing, but maybe yes for peculiar motion); \\item whether the noise that we refer to here from peculiar motion and lensing can in fact be turned into a useful signal (the answer, for lensing, is that the signal is not competitive with the lensing of galaxies). \\end{itemize}  A comment on our terminology: we refer to the fluctuations of interest in this paper as large scale structure induced. The term large scale structure should be viewed as synonymous with departure from homogeneity. Some of the fluctuations discussed here, such as the Poissonian lensing scatter, are in fact dominated by structures on relatively small scales (galactic scales or smaller). Also, even though much of the discussion in this paper is phrased in terms of SNe as standard candles, most of our expressions of course apply equally well to any other distance indicators.  While this paper was in preparation, two preprints \\cite{cooray} appeared in the electronic archive that partially overlap with ours, specifically concerning lensing covariance as noise and signal. See also the preprint by \\cite{dodelson} on SN lensing as a potentially useful signal. There is also a preprint by \\cite{bdg} that discusses fluctuations of the luminosity distance in general as a useful signal (see also \\cite{schucker}). ",
        "conclusions": "\\label{discuss}  Let us summarize the main lessons.  \\begin{itemize}  \\item Large scale structure induced magnitude fluctuations have a significant impact on dark energy measurements from a whole array of ongoing or future SN surveys. For instance, the degradation, due to large scale structure, in the error for the equation of state $w_{\\rm pivot}$ ranges from $10 \\%$ to $60 \\%$ depending on surveys and assumptions (Fig. \\ref{fig:testhistoB}, \\ref{fig:testhistoBnoprior}). It appears difficult to measure from SNe alone the equation of state to better than about $7 - 10\\%$ (depending on assumptions), unless one has a survey considerably more ambitious than SNAP $+$ SNfactory (see Fig. \\ref{fig:contour.1.combo} and Table \\ref{table}).  \\item Of all possible large scale structure fluctuations, the dominant ones are due to peculiar motion and gravitational lensing. Peculiar motion is important at $z \\lsim 0.1$ (through a low redshift anchor such as the SNfactory) while lensing dominates at $z \\gsim 1$. The impact of peculiar motion is mainly through coherent/correlated large scale flows (Fig. \\ref{fig:testcvelB}) while the impact of lensing is mainly through Poissonian fluctuations (Fig. \\ref{fig:testcvellensB}). The Poissonian fluctuations can be reduced by increasing the number of SNe, while the coherent ones can only be suppressed by increasing the survey area. When a high redshift ($z \\gsim 0.1$) survey is combined with a low redshift anchor ($z \\lsim 0.1$), we find that peculiar motion mostly dominates over lensing as a source of error.  \\item What does the above mean for survey designs? As has been emphasized in the literature (e.g. \\cite{SNsurveys}), a low redshift anchor such as the SNfactory is very useful for reducing the eventual dark energy errors from a high redshift SN survey. For such a low redshift survey, one might have hoped to reduce the coherent peculiar motion induced fluctuations by either increasing the survey area or moving it to a higher redshift. {\\it All else being equal, neither will improve appreciably the precision on dark energy determination.} The SNfactory already covers half of the sky; going to full sky will not reduce the errors significantly. For instance, combining SNAP with an all-sky version of SNfactory instead of the half-sky one that we have been assuming, the marginalized error for $\\delta w_{\\rm pivot}$ would improve by only about $1 \\%$. Furthermore, moving SNfactory to a higher redshift, while useful in reducing peculiar motion induced fluctuations, shortens the lever arm that the combination of a high redshift survey and a low redshift anchor offers. The net effect is that moving the low redshift anchor to a higher redshift actually does not reduce the error on the equation of state $w_{\\rm pivot}$ appreciably (Fig. \\ref{fig:testcveldwB}).   \\item How about survey designs for a high redshift survey? For a high redshift survey (considered on its own) that does not extend beyond $z \\sim 1$ (e.g. DES, ESSENCE and SNLS all cover roughly $z \\sim 0.2 - 0.8$), neither peculiar motion nor lensing constitutes significant sources of errors. The only way to reduce dark energy errors is to increase the number of SNe, and suppress systematic errors. The precise survey area is of little importance for such a survey, as long as it is not too small (too small meaning $1$ square degree or less, see Fig. \\ref{fig:testcvellensB}). (Of course, peculiar motion does play a role in the eventual errors once one combines such a high redshift survey with a low redshift anchor, which as emphasized above, is generally a good idea.) For a high redshift survey that extends beyond $z \\sim 1$, gravitational lensing becomes a non-negligible source of errors. But because lensing's impact is mainly through the Poissonian fluctuations it introduces, increasing survey area (such as for SNAP) is not really necessary. The only instance in which a case can be made for increasing survey area is JEDI, which has a sufficiently small Poissonian error (due to its large number of SNe) that coherent/correlated lensing fluctuations actually play a role (see Fig. \\ref{fig:testhistoB} and \\ref{fig:testhistoBnoprior}).  \\end{itemize}  Our investigations in this paper naturally raise a number of questions and issues, some of which we address briefly here, and some require further research.  {\\bf 1.} Perhaps the most natural and interesting question is whether {\\it current} constraints on dark energy, which typically come from some combination of high and low redshift SNe ($z \\gsim 0.1$ and $\\lsim 0.1$), are already affected by peculiar motion. The short answer is: not very much. This is because the current number of low redshift SNe used (typically several 10's) is sufficiently small that the Poissonian error (due to simply intrinsic scatter) is quite a bit larger than the coherent velocity error. This can be inferred from Fig. \\ref{fig:testcvelB}: raising the dotted lines by a factor of $\\sim 10$ (due to dropping the number of SNe from $300$ as in the figure to $\\sim 30$) means the Poissonian intrinsic scatter constitutes a larger source of error than coherent peculiar motion (solid lines). Note that this argument assumes the existing low redshift SNe are selected from a large area of the sky e.g. \\cite{hamuy} (so that it is the lowest few solid lines of Fig. \\ref{fig:testcvelB} that is relevant). The conclusion could be quite different if this assumption does not hold. We urge SN experiments to clearly state the survey areas of their different samples (especially the low $z$ samples) when publishing their results.  {\\bf 2.} In our forecasts for a selection of representative SN surveys (Table \\ref{tabsurveys}), we have assumed simple geometries -- a contiguous circular region on the sky for the high $z$ surveys, and two separate patches (one in the north and one in the south) for the SNfactory. Realistic surveys are bound to be more complicated in shape, possibly with many holes or gaps. The exact geometry affects the size of the coherent fluctuations but not the Poissonian ones. Since for the most part the only coherent fluctuations we need worry about are those due to peculiar motion which is important only at low $z$'s, it is mainly the exact geometry of something like the SNfactory that concerns us. It is therefore worth repeating our calculations for the actual geometry of the SNfactory, including possible extra gaps for instance. For this purpose, we have given sufficiently general expressions for the relevant window functions ($W^{\\rm lens}_{ij}$ and $W^{\\rm vel.}_{ij}$) in eq. (\\ref{Cijfull}) and (\\ref{cijvellarge}), and in eq. (\\ref{WvellargeSNf}) and (\\ref{WvellargeGeneral}) in Appendix \\ref{app:poisson2}. Note that gaps almost always increase the importance of coherent fluctuations because of the introduction of high $k$ modes. For the high redshift surveys that extend beyond $z \\sim 1$, it would be useful to check that a realistic survey geometry does not make the coherent lensing fluctuations much more important (though we do not expect this to happen, as long as the survey area exceeds $\\sim 1$ square degree). For high $z$ surveys that stay within the redshift range $0.1 \\lsim z \\lsim 1$, neither velocity nor lensing fluctuations are expected to be important (unless the number of SNe is much larger than what has been considered), and so the survey geometry has a relatively minor impact.  {\\bf 3.} We have largely ignored internal motion in our discussions of velocity induced fluctuations. By internal motion we mean the motion of the SNe within galaxies, for instance due to the virialized motion of the SN progenitors, or even due to the orbital motion of the SN itself within the binary system that is its progenitor. Such motion could contribute to the overall peculiar velocity of the SNe. Ignoring internal motion is partially justified by the fact that in practice the redshifts are assigned based on the redshifts of the host galaxies. (We thank the referee for emphasizing this to us.) However, internal motion could still in principle modify the apparent luminosity. The important point to keep in mind is that such internal motion is not expected to be correlated between SNe in different galaxies, and so our calculation of the coherent/correlated velocity fluctuations remains valid. Internal motion can certainly increase the Poissonian velocity fluctuations. However, typical virialized motion is of the order of a few hundred km/s, similar to the typical large scale flow velocity, and so the Poissonian velocity fluctuations remain subdominant (Fig. \\ref{fig:testcvelB}). Also, the orbital motion internal to the binary progenitor is too slow to be of significance (see e.g. \\cite{whelaniben}).  {\\bf 4.} Our main focus in this paper is on statistical errors: from intrinsic scatter and from large scale structure induced fluctuations. We have investigated the effect of systematic error in some simple examples (see Table \\ref{table}, entries for 'SNAP $+$ SNf (all $+$ sys)' and 'SNAP $+$ SNf (all $+$ sys2)', and the associated discussion at the end of \\S \\ref{everything}; see also Fig. \\ref{fig:contour.1.combo}) -- we show that even in the presence of systematic error of the assumed magnitudes, large scale structure fluctuations remain a non-negligible source of errors for dark energy measurements. It would obviously be useful to investigate this further and explore a wider range of systematic errors suitable for each SN experiment.  {\\bf 5.} An implicit assumption in our calculations is that the redshift measurements of low $z$ SNe are sufficiently accurate for us to worry about their peculiar motion in the first place. Existing low $z$ measurements typically report an accuracy of $\\delta z \\sim 0.001 - 0.002$ (e.g. \\cite{hamuy96}). The spectral instrument of the SNfactory has a resolution of $1200$, corresponding to $\\delta z \\sim 0.001$. (Note that the actual redshift accuracy is likely to be better than the instrumental spectral resolution, so this is a conservative estimate.) Translating into velocities, we are talking about a velocity of $300 - 600$ km/s, or a magnitude fluctuation of $\\delta m \\sim 0.04 - 0.08$ (for $z \\sim 0.05$). This is still smaller than the intrinsic magnitude scatter that we assume: $\\sigma^{\\rm intr.} = 0.1$ or $0.15$. The redshift uncertainty adds to the Poissonian scatter, and our range of $0.1$ to $0.15$ can be thought of as accounting for this possibility already. It is important, however, that the redshift measurements do not suffer from a systematic bias (that affects all SNe in the same way). From Fig. \\ref{fig:testcvelB}, it can be seen that a systematic bias of $\\delta z \\sim 0.0003$ or $100$ km/s (at $z \\sim 0.05$) would have a comparable effect as coherent peculiar motion.  {\\bf 6.} An interesting question is: to what extent can corrections be made for the velocity and lensing fluctuations? For instance, one could imagine using galaxy weak lensing maps to correct for the magnification of SNe. This has been shown to be not viable, or not sufficiently accurate to be useful, by \\cite{dalal}. This is because galaxy weak lensing maps typically tell us the magnification on scales larger than are relevant for the Poissonian part of SN lensing. (These maps can be useful for correcting the non-Poissonian/coherent part of SN lensing, but this part of lensing is not very important for most SN surveys anyway.) More recently it was argued by \\cite{gunnarsson} that corrections for (the Poissonian part of) SN lensing can be made by modeling foreground galaxies as isothermal spheres or generalizations thereof. One should keep in mind that the Poissonian lensing fluctuations are sensitive to structures on relatively small scales ($k \\gsim 10$ h/Mpc), and a smooth halo profile does not necessarily capture all the relevant fluctuations. For instance, in the case of strong lensing, substructures are often invoked to explain the observed flux ratios \\cite{nealchris}. Most of the high $z$ SNe will not be strongly lensed, but a similar lesson applies here. Moreover, using the foreground galaxies to make a magnification correction inevitably involves assumptions about bias: how galaxies trace mass. As can be seen from Fig. \\ref{fig:testcvellensB}, $\\sigma^{\\rm lens}/\\sqrt{N} \\sim 0.01$ for SNe at $z \\sim 1.5$ and $N = 100$ (per $\\Delta z$ of $0.1$, as appropriate for SNAP for instance). For the lensing correction to be useful, it should therefore satisfy two criteria: first, the correction should be more accurate than $0.1$ (in magnitude) per SN; second, it should not introduce a systematic bias that is larger than $0.01$ (in magnitude; magnitude fluctuation $\\sim$ magnification fluctuation). Even if the first can be achieved, the second seems challenging.  How about corrections for velocity fluctuations? Here, the situation is slightly different: what needs to be corrected is the coherent part (i.e. large scale), not the Poissonian part, of the fluctuations. Roughly speaking, we need to know the low order multipoles of the peculiar flow at the redshift of e.g. the SNfactory ($z \\sim 0.055$). One option is to use peculiar velocity surveys (such as from the SNe themselves), but it should be kept in mind that to disentangle the Hubble flow from say the monopole, one needs a survey that has the same sky coverage as the SNfactory, but is deeper. This requires considerable resources. Another option is to use the galaxy spatial distribution as a guide, i.e using mass conservation to relate peculiar velocity to the galaxy overdensity. Such a procedure of course suffers from the uncertain biasing relation between galaxies and mass. Note also that one needs a galaxy survey that is deeper than the SNfactory to define the correct mean galaxy density. (For a recent paper that examines the peculiar motions predicted by the PSCz survey, see \\cite{john}; it focuses on peculiar flows at slightly lower redshifts than we need.) Whether either option allows us to take out the effect of bulk flows to sufficient accuracy is a question we would like to address in the future.  {\\bf 7.} Another natural and interesting question is: to what extent can the noise here, due to lensing and peculiar motion, be viewed as a useful signal? In the case of lensing, the issue is discussed in several recent papers \\cite{williams,wang,menard,cooray,dodelson}. In general, it is difficult for SNe to be competitive with galaxies as the sources for weak gravitational lensing experiments. Consider for instance the measurement of the convergence power spectrum: the shot-noise in the case of SNe is $(\\sigma^{\\rm intr.})^2/(4n)$ (the factor of $4$ comes from $\\delta m \\sim 2 \\kappa$ where $\\kappa$ is convergence), while the shot-noise in the case of galaxies is $(\\sigma^\\gamma)^2/(2n)$ (the factor of $2$ comes from the use of two components of shear to estimate $\\kappa$). Here the intrinsic magnitude scatter $\\sigma^{\\rm intr.}$ is roughly $0.1 - 0.15$, the shape noise $\\sigma^\\gamma$ is about $0.3$, and the surface density $n$ is approximately $0.04$ per square arcminute for SNe (taking numbers from SNAP), while $n \\sim 30$ per square arcminute for a typical weak lensing galaxy survey. The shot-noise from SNe is simply too big compared to that from galaxies. Nonetheless, the lensing of SNe is free from certain systematic errors that might affect the lensing of galaxies, such as intrinsic alignment, and so the SN method still provides a useful, though not terribly stringent, consistency test.  How about the SN peculiar motion as a signal? For some of the earlier work on this issue, see e.g. \\cite{hamuy96,riess95,riess97,idit}. Among the different methods for measuring peculiar velocities (see \\cite{strauss} for a review), SNe Ia constitute the most accurate distance indicator on an object by object basis. For instance, SNe Ia yield distances with an error of $\\sim 5 - 7\\%$, while Tully-Fisher distances are typically uncertain at the $15 - 20 \\%$ level. On the other hand, Tully-Fisher galaxy catalogs (e.g. \\cite{willick}) typically have significantly more objects than SN surveys, and therefore have perhaps more statistical power. Yet, SN surveys might suffer less from systematic errors that seem to have plagued at least some Tully-Fisher galaxy catalogs (e.g. \\cite{davis}), and SN surveys generally go deeper. It remains an interesting question to what extent competitive cosmological constraints can be obtained from the peculiar motion of SNe. We hope to explore this in the future."
    },
    "0512/astro-ph0512645_arXiv.txt": {
        "abstract": "We use the ROSAT All Sky Survey (RASS) to study the soft X-ray properties of a homogeneous sample of 46,420 quasars selected from the third data release of the Sloan Digital Sky Survey (SDSS).  Optical luminosities, both at rest-frame 2500\\AA\\, ($L_{2500}$) and in [OIII] ($L_{[\\rm{OIII}]}$) span more than three orders of magnitude, while redshifts range over $0.1<z<5.4$. We detect 3366 quasars directly in the observed 0.1--2.4 keV band. Sub-samples of radio-loud and radio-quiet objects (RLQs and RQQs) are obtained by cross-matching with the FIRST catalogue. We study the distribution of X-ray luminosity as a function of optical luminosity, redshift and radio power using both individual detections and stacks of complete sets of similar quasars. At every optical luminosity and redshift $\\log L_{2\\kev}$ is, to a good approximation, normally distributed with dispersion $\\sim 0.40$, at least brightwards of the median X-ray luminosity. This median X-ray luminosity of quasars is a power law of optical luminosity with index $\\sim 0.53$ for $L_{2500}$ and $\\sim 0.30$ for $L_{[\\rm{OIII}]}$.  RLQs are systematically brighter than RQQs by about a factor of 2 at given optical luminosity. The zero-points of these relations increase systematically with redshift, possibly in different ways for RLQs and RQQs.  Evolution is particularly strong at low redshift and if the optical luminosity is characterised by $L_{[\\rm{OIII}]}$. At low redshift and at given $L_{[\\rm{OIII}]}$ the soft X-ray emission from type II AGN is more than 100 times weaker than that from type I AGN. ",
        "introduction": "Quasars show strong emission at both ultraviolet (UV) and X-ray wavelengths. Indeed, many quasar catalogues have been constructed on the basis of their UV and X-ray properties (e.g. Schneider et al. 2003, 2005; Wolf et al. 2004). The relation between the UV and X-ray continuum emission is usually characterized by a spectral index, \\beq \\alpha_{OX} \\equiv -\\frac{\\log(L_{2\\kev}/L_{2500})} {\\log(\\nu_{2\\kev}/\\nu_{2500})} = -0.384\\log(L_{2\\kev}/L_{2500})\\,, \\eeq where $L_{2\\kev}$ and $L_{2500}$ are luminosities per unit frequency at wavelengths of 2\\,keV and 2500\\AA, respectively.  Previous investigations have indicated that the value of $\\alpha_{OX}$ changes systematically with optical luminosity $L_O$ (e.g. Zamorani et al. 1981; Avni \\& Tananbaum 1982, 1986; Bechtold et al. 1994; Pickering, Impey \\& Foltz 1994; Avni, Worrall \\& Morgan 1995; Green et al. 1995; Vignali, Brandt \\& Schneider 2003a; Anderson et al. 2004 and references therein). Most of these studies also concluded that $\\alpha_{OX}$ depends only weakly on redshift, but separation of the redshift and luminosity dependences is difficult because of the strong correlation between the two quantities in flux-limited samples (e.g. Anderson \\& Margon 1987; Bechtold et al. 2003). A dependence of $\\alpha_{OX}$ on $L_O$ implies that the relation between $L_X$ and $L_O$ is nonlinear, the trend in the observational data being represented by a power law, \\beq L_X\\propto L_O^e, \\eeq with $e<1$. However, this trend can be affected by a variety of observational factors, for example, sample definition, sample completeness, observational errors, and the adopted fitting method. Consequently, the exact form of the mean $L_X$-$L_O$ relation is still controversial (Franceschini et al.  1994; Yuan et al. 1998; Vignali et al. 2003b). In particular, as pointed out by La Franca et al. (1995) and Yuan, Siebert \\& Brinkmann (1997), for some earlier datasets it was conceivable that large photometric errors in the optical luminosities had biased an intrinsically linear relation into an apparently nonlinear one.  In order to establish the $L_O$-$L_X$ relation robustly, a large and homogeneous quasar sample with accurate optical and X-ray measurements is required. The quasar catalogue selected from the SDSS (York et al. 2000) third data release (DR3) includes more than 46,000 quasars (Schneider et al. 2005) and is currently by far the largest available (as of July 2005). The SDSS quasar selection algorithm is quite efficient, and the completeness at $z<3$ is at least 90 percent (Richard et al. 2002; Vanden Berk et al. 2005). Deep surveys which can be used to select X-ray AGN to faint limits cover small areas of the sky (e.g. Wolf et al. 2004) and so overlap at most a small fraction of the SDSS data (e.g. Risaliti \\& Elvis 2005). On the other hand, the ROSAT All-Sky Survey (RASS) is quite shallow, with an average exposure time of about $\\sim400$s, and does not detect most SDSS quasars individually.  It is possible, however, to detect these objects statistically by stacking their X-ray images. Such stacking is widely used in X-ray astronomy for objects which are individually below the detection limit (e.g. Wu \\& Anderson 1992; Schartel et al. 1996; Nandra et al. 2002; Georgakakis et al. 2003). Since the number of SDSS quasars is large, the detection limits can be greatly improved by stacking many objects with similar optical properties.  Different classes of quasars are observed to have different X-ray properties. For example, radio-loud objects are systematically brighter in X-rays than radio-quiet ones (e.g. Ku, Helfand \\& Lucy 1980; Zamorani et al. 1981; Bassett et al. 2004). Type II quasars are usually much weaker in the soft X-ray band than type I objects, presumably due to stronger absorption (e.g.  Zakamska et al. 2004; Vignali, Alexander \\& Comastri 2004). BL Lacertae (BL Lac) objects show strong and rapid variability in all bands. In combination, the RASS and the large quasar catalogue provided by SDSS enable a comparative study of quasars in these different categories.  In this paper, we use the RASS to study the soft X-ray properties of the SDSS DR3 quasars. Optical properties for the sample are derived from the SDSS photometry and spectroscopy. We will make particular use of the continuum luminosity at 2500\\AA\\, rest wavelength $L_{2500}$ and the [OIII]$\\lambda$5007 line luminosity $L_{[\\rm{OIII}]}$, both of which can be determined accurately for large samples of quasars. Our analysis considers X-ray fluxes both for individually detected quasars, and for stacks of quasars of similar redshift and optical luminosity. We develop statistical methods specifically designed to study the joint distribution of soft X-ray and optical properties in the SDSS-RASS samples. In addition, we use the FIRST survey (Becker, White \\& Helfand 1995) to split the SDSS quasars into radio-loud and radio-quiet classes, and we study the X-ray properties of these two sub-samples separately. The SDSS quasar catalogue does not include type II and BL Lac objects (Schneider et al. 2005), but a sample of nearby type II AGN has been constructed from the SDSS galaxy catalogues by Kauffmann et al.  (2003). We will use this sample to measure the average soft X-ray luminosities of such objects.  Our paper is organized as follows. In Section 2 we introduce the SDSS DR3 quasar sample and explain how we define their optical and radio properties. Section 3 describes our RASS detection technique, both for individual objects and for stacks. In Section 4, we use a variety of techniques to study the joint $L_X$ -- $L_O$ distribution. We derive mean relations from complete sets of X-ray detections of stacks spanning the full redshift and optical luminosity range of the sample. We use individual detections to study the high luminosity tail of the distribution of X-ray luminosity at given optical luminosity. We use complete sets of detections for stacks of objects which are {\\it not} individually detected to study the lower luminosity part of this distribution.  All approaches give consistent results. We also split the quasar sample in various ways to examine the dependence of the relation on redshift, on radio emission strength and on the presence or absence of broad optical emission lines.  We discuss our results and draw our conclusions in Section 5. ",
        "conclusions": "In this paper, we have studied the soft X-ray properties of quasars based on the SDSS DR3 quasar catalogue and the RASS. The FIRST catalogue enabled us to define subsamples of radio-loud and radio-quiet objects. We used both individual and stack detections to investigate the X-ray properties of quasars as a function of $L_{2500}$, of radio power, of $z$ and of $L_{[\\rm{OIII}]}$.  By stacking all QSOs with similar optical luminosity and redshift, we have shown that the average $L_{2\\kev}$ of both RLQs and RQQs depends significantly on redshift at given optical luminosity. At fixed UV continuum luminosity $L_{2500}$ and at low redshift ($z<0.5$) the RLQ population may evolve more strongly than the RQQs. At higher redshift, the X-ray brightening of RLQs and RQQs of given $L_{2500}$ are both considerably slower.  If [OIII] line luminosity $L_{[\\rm{OIII}]}$ is substituted for $L_{2500}$, then RLQs and RQQs show similar and strong evolution out to $z=0.8$. Assuming that the redshift and optical luminosity dependences of the $L_{2\\kev}$ distribution separate, we correct all X-ray luminosities to a fiducial redshift and analyse in more detail the shape of the distribution of the evolution-corrected X-ray luminosity $\\log L_{2\\kev,z_0}$ at given optical luminosity. This distribution is well approximated by a log-normal, at least for quasars more X-ray bright than the median.  Adopting this model we find a tight but nonlinear relation between $\\tilde{L}_{2\\kev}$ and optical luminosity ($L_{2500}$ or $L_{[\\rm{OIII}]}$). The dispersion in $\\log L_{2\\kev}$ is roughly $\\sim 0.40$, independent of optical luminosity. The typical X-ray luminosity of RLQs is twice that of RQQs of the same $L_{2500}$, and 1.6 times that of RQQs of the same $L_{[\\rm{OIII}]}$. In addition, we find the average soft X-ray emission of type II AGN to be more than 100 times weaker than that of radio-quiet type I quasars of the same $L_{[\\rm{OIII}]}$ and redshift.  Many of our results have been seen in previous studies. For example, the nonlinear relation between $L_{2\\kev}$ and $L_{2500}$ is consistent with that of Vignali, Brandt \\& Schneider (2003a); the result that RLQs have stronger soft X-ray emission than RQQs is consistent with that of Brinkmann et al. (2000); the result that soft X-ray emission from type II AGN is very weak is consistent with that of Zakamska et al. (2004). There are also a number of new results, for example, that $\\alpha_{OX}$ depends not only on optical luminosity but also on redshift. At given optical luminosity $L_{2500}$, the typical value of $\\alpha_{OX}$ decreases by about 0.3 from $z\\approx0$ to $z\\approx5$ [extrapolated from our best fits to equation (\\ref{zL5p})] both for RLQs and for RQQs. This is significant compared to earlier claims for the redshift independence of $\\alpha_{OX}$. For example, Strateva et al. (2005) quote $\\Delta\\alpha_{OX}\\approx0.03$ from $z\\approx$0 to 5. The apparent change in redshift dependence around $z\\sim0.5$ is very interesting. We note that rest-frame 2500\\AA\\, passes out of the SDSS spectral range at about this redshift. We obtain the 2500\\AA\\, flux for such low redshift quasars by extrapolation using the Vanden Berk et al. (2001) composite quasar spectrum, a procedure which could plausibly introduce systematic biases into $L_{2500}$. To check for such effects, we repeated our analysis replacing $L_{2500}$ by $L_{3500}$, the continuum luminosity at rest-frame 3500\\AA. This can be estimated directly from the SDSS spectrum for all our quasars with $z<1.63$. This substitution did not alter any of our conclusions about low-redshift evolution either for RLQs or for RQQs. Another possible systematic might arise from host-galaxy contributions to $L_{2500}$ in low-redshift ($z<0.5$) and low-luminosity quasars. Strateva et al. (2005) addressed this issue in their own study, concluding that the host-galaxy contribution does not exceed 20 percent.  The most detailed previous study of the relation between $L_{2\\kev}$ and $L_{2500}$ is that of Strateva et al. (2005), based on a complete sample of 228 radio-quiet non-BAL quasars more than 80\\% of which were detected in medium-deep pointed observations with ROSAT. The slope of $0.65\\pm 0.02$ which they quote for the $\\log L_{2\\kev}$--$\\log L_{2500}$ relation is in excellent agreement with the values we find for all our samples when we fit our ``total'' stacks without allowing for evolution [equation (\\ref{ARQQ})]. However, after allowing for evolution, our analysis predicts a shallower value, $0.53\\pm0.02$, both from our individual detections and from our ``total'' stacks. This is a consequence of the significant evolution we detect, together with the strong selection-induced correlation between optical luminosity and redshift in our samples. Moreover, since $\\log L_{2\\kev}$ is normally distributed with a scatter of about 0.40, the median $L_{2\\kev}$ at given $L_{2500}$ is smaller than the mean by a factor of about 1.7. The excellent apparent agreement in Fig. \\ref{All_RQQ} between the Strateva et al. relation and our own RQQ relation is thus misleading. The Strateva et al. result gives median $L_{2\\kev}$ as a function of $L_{2500}$ since they fit in logarithm space; on the other hand, our result gives the mean $L_{2\\kev}$ of QSOs of given $L_{2500}$. Thus the relations should be off-set by a factor of 1.7. Quasar variability could perhaps account for this discrepancy given the longer X-ray exposure times in the Strateva et al. (2005) data. QSOs vary on time-scales of $\\sim 10^4$ seconds at X-ray wavelengths with typical amplitudes of about 0.2 in $\\log L_{2\\kev}$ (e.g. Almaini et al. 2000; Manners, Almaini \\& Lawrence 2002). In our study such variability contributes to the scatter in the $L_{2\\kev}$ distribution at given optical luminosity, since exposure times for our quasars are shorter than $\\sim 10^4$s. Studies based on longer X-ray exposures should find a distribution with a higher median and lower scatter (but with the same mean). It is suggestive that Strateva et al.  find a smaller scatter in $\\log L_{2\\kev}$ than we do, 0.29 (see Fig. 14 in their paper) rather than 0.40.  La Franca et al. (1995) and Yuan, Siber \\& Brinkmann (1997) argued that photometric error in the optical luminosities might cause an apparently weaker than linear dependence of $L_{X}$ on $L_{O}$, even if the true relation is linear.  This explanation is no longer tenable for data of the quality analysed here [or indeed in Strateva et al. (2005)]. The relations we find really are much weaker than linear, $\\tilde{L}_{2\\kev}\\propto L_{2500}^{0.53}$(RQQs) and $\\tilde{L}_{2\\kev}\\propto  L_{[\\rm{OIII}]}^{0.33}$. The difference between the slopes for the two different optical measures of quasar activity is large and interesting.  The most important assumption in our study which still remains to be fully justified observationally is the log-normal form we adopt for the distribution of $L_{2\\kev}$ at given optical luminosity. As shown by Figures \\ref{simudata1} and \\ref{simudata2}, all the observational data (e.g. the fraction of individually detected objects, the distribution of their $L_{2\\kev}$, the mean $L_{2\\kev}$ both for stacks of all sources and for stacks of individually undetected sources) appear consistent with this assumption, but they do not constrain the shape of the distribution at low X-ray luminosities. Note, however, that the part of the distribution we can characterize well already corresponds to more than 50\\% of the sources and to about 75\\% of their total X-ray output. Our log-normal hypothesis will be more thoroughly tested by future X-ray studies which will simultaneously achieve wide area coverage and deep sensitivity limits, resulting in much greater detection completeness than is possible with the RASS."
    },
    "0512/astro-ph0512229_arXiv.txt": {
        "abstract": "Type Ia supernovae have played a crucial role in the discovery of the dark energy, via the measurement of their light curves and the determination of the peak brightness via fitting templates to the observed lightcurve shape. Two spectroscopic indicators are also known to be well correlated with peak luminosity. Since the spectroscopic luminosity indicators are obtained directly from observed spectra, they will have different systematic errors than do measurements using photometry. Additionally, these spectroscopic indicators may be useful for studies of effects of evolution or age of the SNe~Ia progenitor population. We present several new variants of such spectroscopic indicators which are easy to automate and which minimize the effects of noise. We show that these spectroscopic indicators can be measured by proposed JDEM missions such as \\snap\\ and JEDI. ",
        "introduction": "Type Ia supernovae are now recognized as one of the most important cosmological probes. Astronomical interest has been focused on Type Ia supernovae since it was recognized long ago \\citep{wilson39,kowal68} that they are good ``standard candles'' and hence are useful cosmological probes. Their use as cosmological beacons led to the discovery of the ``dark energy'' \\citep{riess_scoop98,garnetal98,perletal99} and they are an integral part of plans to further characterize the nature of the dark energy equation of state with a wide field satellite probe, the ``Joint Dark Energy Mission'' (JDEM) as well as ground-based studies. The reliability of SNe Ia as distance indicators improved significantly with the realization that the luminosity at peak was correlated with the width of the light curve \\citep{philm15} and hence that SNe~Ia are correctable candles in much the same way that Cepheids are \\citep{philetal99,goldhetal01,rpk95}. To date, cosmology with Type Ia \\supernovae has used photometry exclusively over spectroscopy, except that a spectrum is required for supernova type confirmation and can be used for the redshift determination, the luminosity at peak is always determined photometrically (by matching lightcurve templates). However, there exist spectroscopic luminosity indicators first discussed by \\citet{nugseq95}.  \\citet{garn99by04} showed that \\RSi\\ \\citep{nugseq95} is well-correlated with $\\Delta m_{15}$. Of course, the value of $M_B$ at maximum light is also well correlated with $\\Delta m_{15}$ \\citep{philetal99,philm15}. Thus, the spectroscopic luminosity indicator should be well correlated with the peak luminosity, making it a useful secondary luminosity indicator for cosmological studies. Even though it is correlated with peak luminosity via the photometric indicator $\\Delta m_{15}$, it is still useful as an independent check because it will have different systematic errors than light curve shape fitting via templates.  We study the correlation of the two spectroscopic luminosity indicators \\RSi\\ and \\RCa\\  with luminosity, both for observed \\supernovae and for synthetic spectra of the parameterized explosion model W7 \\citep{nomw7} calculated using \\phx \\citep{hbjcam99,hbmathgesel04} in LTE. We define new, more robust spectroscopic luminosity indicators and we show that these indicators can be used with good accuracy via proposed JDEM missions such as SNAP \\citep{snap_detf_wp05} and JEDI \\citep{jedi_detf_wp05}. ",
        "conclusions": "We have presented a way to automatically measure the value of the spectral indices \\RCa\\ and \\RSi, which should be robust. We have defined three new spectral indices \\RCaS, \\RSiSu, and \\RSiSuS, which are less sensitive to noise in the spectrum. We have shown that for the limited set of Branch Normal supernovae where public spectra are available near maximum light that there is a good correlation between the spectral indices and the luminosity at peak. Within the caveat of small size of our sample we found that there is not a general trend in the spectral indices with epoch as long as the epoch is close to maximum light. We have used synthetic spectral models to show qualitatively that there does appear to be a deviation from linearity in the spectral indices, particularly among the dim, fast-declining \\supernovae, such as SN~1991bg. This result is in agreement with that of \\citet{philetal99} who also found that there is a deviation from linearity in peak magnitude with $\\delta m_{15}$ as $\\delta m_{15}$ increases (see their Figure 8). It is also in agreement with some theoretical hydrodynamical results which suggest it may be difficult to produce a SN~Ia with such low nickel mass via a deflagration (W.~Hillebrandt, private communication), although \\citet{HGFS99by02} are able to vary the nickel mass by varying the density of the deflagration to detonation transition. It has been suggested that there is both a variation in \\SNeIa luminosity with galaxy type \\citep{blue96,hametal01} and that there are two populations of \\SNeIa progenitors \\citep{SB05,MDVP05} which could be related to the deviation in the linearity of the spectral indices.   Restricting ourselves to the sample of Branch Normal \\supernovae we find that the proposed JDEM mission SNAP could in fact use \\RSiSu\\ as a complementary luminosity indicator. While the results may appear somewhat circular in that we have restricted our calibrators to have a small dispersion and then shown that the correlation based on that result also recovers the small dispersion, the ability to use a \\emph{separate} observable from light-curve shape will provide for a complementary measurement of the luminosity distance as a function of redshift in the proposed JDEM wide-field high-redshift spaced based surveys such as \\snap\\ and JEDI. The lightcurve based method of template fitting will have a certain set of systematics that depends on exactly how ``like versus like'' is chosen, whereas the spectral index method will have a different set of systematics. In addition, the spectral index methods are very weakly sensitive to both the effects of dust and of the flux calibration of the spectrum, since they are defined in a very narrow wavelength range.  Our simulations show that the single spectrum obtained by \\snap\\ for type identification will have sufficient signal to noise to measure the spectral indices to the required precision. JEDI, which will obtain multiple spectra for each supernova, would be even better suited to using \\RSiSu\\ since they would have a larger sample of spectra very close to maximum light.  As more spectra from nearby \\supernovae with known peak luminosities become available we should be able to refine our calibration, and thus improve its accuracy for JDEM and other \\supernovae surveys. We note that our actual numerical results rely on both a very small sample and also on the value of the absolute magnitudes assigned by \\citet{reindl05}, although since we quote results with both Case A and B the sensitivity to the assigned absolute magnitudes can be gauged somewhat. All available distance estimates (and hence absolute magnitudes) to our sample will contain errors, and we have chosen to use the results of \\citet{reindl05} because they are derived in a specific somewhat homogeneous manner. Thus, we regard this work as a ``proof of concept'' and the specific numerical values as our best current estimate. Clearly when 300 Hubble-flow supernovae are available from current ground-based surveys (i.e. in the next 3-5 years), this analysis should be recomputed. Once accurate values are in hand, comparing absolute magnitude (determined from light-curve shape) and \\RSiSu\\ determined at high redshift may provide a useful measure of the evolutionary effects in the high-redshift sample. We discuss this further in a forthcoming publication where we study the formation of the \\SiIIblue\\ line."
    },
    "0512/astro-ph0512586_arXiv.txt": {
        "abstract": "Essentially all of our knowledge of the acceleration history of the Universe -- including the acceleration itself -- is predicated upon the validity of general relativity.  Without recourse to this assumption, we use SNeIa to analyze the expansion history and find (i) very strong ($5\\sigma$) evidence for a period of acceleration, (ii) strong evidence that the acceleration has not been constant, (iii) evidence for an earlier period of deceleration and (iv) only weak evidence that the Universe has not been decelerating since $z \\sim 0.3$. ",
        "introduction": "A still puzzling feature of our Universe is its accelerated expansion.  Measurements of type Ia supernovae (SNeIa) provide direct evidence \\cite{Riess:1998cb,Perlmutter:1998np,Astier:2005qq} and CMB anisotropy measurements provide important indirect evidence \\cite{Tegmark:2003ud}.  Within the context of general relativity, the acceleration can be explained if the present mix of matter and energy in the Universe includes a dark energy component that contributes about 2/3 of the critical density.  Dark energy is defined by its large negative pressure, $p\\sim-\\rho$, and nearly spatially uniform distribution \\cite{Turner:1998ex}.  While the energy of the quantum vacuum, mathematically equivalent to a cosmological constant, is the simplest explanation for the dark energy and cause of cosmic acceleration, all attempts to compute its numerical value have been unsuccessful, either leading to divergent results or numbers that are orders of magnitude too large; this is known as the cosmological constant problem \\cite{Carroll:2000fy}.  Because it is possible that the energy of the quantum vacuum is zero or too small to explain cosmic acceleration, theoretical physicists have explored a plethora of possibilities for the dark energy, from a very light scalar field \\cite{Ratra:1987rm,Wetterich:1987fm,Frieman:1995pm,Coble:1996te,Caldwell:1997ii} to the influence of extra dimensions \\cite{Deffayet:2001pu}, and have even considered the possibility that there is no need for dark energy and that it arises due to heretofore neglected ``ordinary'' physics within general relativity \\cite{Kolb:2005da}.  A fair summary of the present state of affairs is that there is little understanding of why the Universe is accelerating.  Given that state of affairs, it is not unreasonable to think more broadly, even entertaining the possibility that the explanation may well have to do with gravity theory itself \\cite{Carroll:2003wy,Capozziello:2003tk,Freese:2002sq,Arkani-Hamed:2002fu,Dvali:2003rk}.  After all, we are confident that Einstein's theory must be modified to make it compatible with quantum mechanics, and perhaps cosmic acceleration has something to do with that modification.  However, relaxing the assumption of general relativity means that we must abandon the Friedmann equation, which relates the expansion rate to the matter and energy densities.  We might then wonder what we really know about cosmic acceleration -- or if the Universe is {\\it really} accelerating since almost everything we have inferred about it depends upon the assumption of the existence of dark energy.  To be specific, dark energy is often described as a smooth energy component with constant equation-of-state ($w=p/\\rho\\,=$\\,const\\,$\\sim -1$), where quantum vacuum energy is the special case $w = -1$.  The acceleration history of the Universe is then simply fixed by the matter content ($\\Omega_M$) and the value of $w$. It begins with a period of decelerated expansion and then crosses over to accelerated expansion at redshift $$z_c = \\left[\\frac{\\Omega_M}{(3w+1)(\\Omega_M -1)}\\right]^{1/3w}-1\\sim 0.5\\,;$$ see \\rf{allqplot}.  Both $w$CDM with $w =-1.05 \\pm 0.14$ and $\\Omega_M =0.29\\pm0.03$ and $\\Lambda$CDM with $\\Omega_M =0.30 \\pm0.05$ are reasonable fits to the SNe data and a wealth of other cosmological data \\cite{Tegmark:2003ud}. (This in fact is the basic evidence for a period of accelerated expansion.)  The existence of an early period of decelerated expansion is also supported by the success of the gravitational instability theory of structure formation and of big-bang nucleosynthesis.  In particular, unless the equations governing the growth of small density perturbations are radically different, it is not possible to grow structure from small perturbations in an accelerating universe \\cite{Turner:1998ex}. In addition, it is difficult for a universe that accelerates during BBN to reproduce the observed light element abundances unless one artificially adjusts with the cosmic neutrino asymmetry \\cite{Carroll:2001bv}.  In this paper we relax the assumption of general relativity, but retain the weaker assumption of an isotropic and homogeneous space-time (which is supported by ample observational data) described by a metric theory of gravity (because there are few if any viable non-metric theory alternatives).  In this case, SNeIa still provide a direct kinematic probe of the expansion, and we show that we can directly infer the following:  (i) very strong ($5\\sigma$) evidence for a period of accelerated expansion; (ii) strong evidence that the acceleration has not been constant; (iii) evidence for an earlier period of deceleration; and (iv) only weak evidence that the Universe has not been decelerating since $z\\sim 0.3$.  \\begin{figure}[t] \\sfig{allqplot2.eps}{.5\\columnwidth} \\caption{{\\small Different parametrizations of $q(z)$.  The solid line is the best-fit 2-epoch model (\\re{q2epoch}); the long-dashed line is the best-fitting linear model; the short-dashed line is the best-fitting \\LCDM model ($\\Omega_M=0.31$).  The kinematic models suggest a later transition to acceleration than \\LCDM does. }} \\label{fig:allqplot} \\end{figure} ",
        "conclusions": "Understanding dark energy may ultimately require us to abandon or change the familiar equations of general relativity.  Hence it is crucial that we keep track of which observations and conclusions depend on the assumption of general relativity and which are more robust.  Without the Friedmann equation, it is still possible to describe the Universe with a Robertson-Walker metric and infer the expansion history via the dimming and cosmological redshifting of SNeIa.  We used this fact to learn about cosmic acceleration kinematically.  In particular, we have shown that the gold set of SNeIa from Riess \\etal provides the following information:  \\begin{itemize} \\item {\\bf Very strong evidence that the Universe once accelerated --}  Though various parametrizations can imply this, the strongest evidence comes from measuring the best-constrained principal component of $q(z)$.  The acceleration is likely to have been relatively recent in cosmic history. \\item {\\bf Strong evidence that $q$ was higher in the past --}  The constant $q$ model has a poor goodness-of-fit, especially when fit in mode space using only well-measured modes.  In addition, the linear $q(z)$ model indicates that the average $dq/dz$ is positive. \\item {\\bf Weak evidence that the Universe once decelerated --}  Both the two-epoch and two-mode models of $q(z)$ slightly favor past deceleration.  The linear model also predicts a transition to positive $q$.  Alas, it is difficult to show that this is not a model-dependent feature. \\item {\\bf Little or no evidence that the Universe is presently accelerating --}  It is manifestly difficult to constrain $q(z<0.1)$ with SNeIa.  Using a three-epoch model, we further showed that a long period (e.g., $z= 0 \\rightarrow 0.3$) of recent deceleration is consistent with data. \\item {\\bf \\LCDM and $w$CDM are acceptable fits to SNeIa data.} While there are tantilizing indications that the actual acceleration history may be different, none of the kinematic models studied here have revealed robust features about cosmic acceleration that differ from \\LCDM or $w$CDM.  Moreover, none of the kinematic models fit the data significantly better. \\end{itemize}"
    },
    "0512/astro-ph0512065_arXiv.txt": {
        "abstract": "{We present results of simulations of stellar collapse and explosions in spherical symmetry for progenitor stars in the 8--$10\\,M_\\odot$ range with an O-Ne-Mg core. The simulations were continued until nearly one second after core bounce and were performed with the \\textsc{Prometheus/Vertex} code with a variable Eddington factor solver for the neutrino transport, including a state-of-the-art treatment of neutrino-matter interactions. Particular effort was made to implement nuclear burning and electron capture rates with sufficient accuracy to ensure a smooth continuation, without transients, from the progenitor evolution to core collapse. Using two different nuclear equations of state (EoSs), a soft version of the Lattimer \\& Swesty EoS and the significantly stiffer Wolff \\& Hillebrandt EoS, we found no prompt explosions, but instead delayed explosions, powered by neutrino heating and the neutrino-driven baryonic wind which sets in about 200$\\,$ms after bounce. The models eject little nickel ($< 0.015M_\\odot$), explode with an energy of $\\ga0.1\\times 10^{51}\\,$erg, and leave behind neutron stars (NSs) with a baryonic mass near $1.36\\,M_\\odot$. Different from previous models of such explosions, the ejecta during the first second have a proton-to-baryon ratio of $Y_{\\rm{e}} \\ga 0.46$, which suggests a chemical composition that is not in conflict with galactic abundances. No low-entropy matter with $Y_{\\rm{e}} \\ll 0.5$ is ejected. This excludes such explosions as sites of a low-entropy r-process. The low explosion energy and nucleosynthetic implications are compatible with the observed properties of the Crab supernova, and the small nickel mass supports the possibility that our models explain some subluminous Type II-P supernovae. ",
        "introduction": "Recent observations of subluminous Type II-P supernovae (SNe) like 2005cs, 2003gd, 1999br and 1997D, have renewed attention to stars near the lower end of the mass range of core-collapse SN progenitors, i.e. to stars with about 8--$10\\,M_\\odot$, which develop O-Ne-Mg cores. A possible link between both has been suggested because of the low $^{56}$Ni and $^{16}$O ejecta masses and low progenitor luminosities (e.g., Chugai \\& Utrobin 2000, Hendry et al.\\ 2005). However, due to many uncertainties this connection is far from being clear (e.g., Pastorello et al.~2004, 2005; Hamuy 2003; Zampieri et al. 2003, and refs.\\ therein). Also the Crab Nebula's progenitor was proposed to be in this mass window (Gott et al.~1970, Arnett 1975, Woosley et al.~1980, Hillebrandt 1982). The observed composition of the Crab remnant (small C and O abundances, He overabundance) was interpreted as a strong indication that the Crab Nebula comes from a collapsing and exploding progenitor with an O-Ne-Mg core (Davidson et al.~1982, Nomoto et al.~1982, Nomoto 1983).  Moreover, these stars were considered as possible sites for a low entropy r-process (for example, Hillebrandt 1978, Wheeler et al.~1997, Sumiyoshi et al.~2001, Wanajo et al.~2003) based on the assumption that they explode by the prompt bounce-shock mechanism, which Hillebrandt et al.~(1984) found to work in a numerical simulation, taking Nomoto's O-Ne-Mg core model (Nomoto 1984, 1987). Such explosions are characterised by the direct propagation of the shock out of the core, the formation of a mass cut, and the continuous acceleration of the material outside of the mass cut to high velocities. They are expected to eject relatively large amounts of neutron-rich matter with low $Y_{\\rm{e}}$ ($\\sim$0.2) and low entropies ($\\sim\\,$10$k_{\\rm{B}}$ per nucleon). However, several groups could not confirm the viability of the prompt explosion mechanism (Burrows \\& Lattimer 1985, Baron et al.~1987, Mayle \\& Wilson 1988). Mayle \\& Wilson (1988) continued their simulations in the post-bounce phase for a longer time and obtained instead a so-called neutrino-driven, delayed explosion (Bethe \\& Wilson 1985) with a low production of $^{56}$Ni (approximately 0.002 $M_\\odot$) in agreement with subluminous Type II-P SNe as mentioned above, but with a vast overproduction of neutron-rich material (at least 0.02 $M_\\odot$ of ejecta with $Y_{\\rm e}\\la0.41$). The latter finding is inconsistent with the chemical composition of our galaxy, which allows for no more than 10$^{-3}$ $M_\\odot$ of material with $Y_{\\rm e}<0.42$ being ejected per SN (Hartmann et al.~1985). Moreover, the explosion energies of both studies, around $2\\times10^{51}$erg in Hillebrandt et al.'s (1984) model and between 0.6 and 1.2 $\\times10^{51}$erg in Mayle \\& Wilson's (1988) simulations, would be inconsistent with the long plateau phase of the above mentioned subluminous SNe, if their H-envelope masses were $\\la8M_\\odot$.  It was suggested that the reason for the discrepant results in the SN simulations of O-Ne-Mg cores (prompt explosions, delayed explosions, no explosions) could be explained by the different nuclear EoSs used by the groups (Fryer et al.~1999). Having in mind that the different approximations in the neutrino transport in previous calculations introduced additional uncertainties, we revisit this topic with a state-of-the-art neutrino transport treatment together with a careful description of weak interactions and including relevant nuclear burning reactions. We additionally make a comparison of collapse and post-bounce calculations with different nuclear EoSs. ",
        "conclusions": "Our 1D simulations of SN explosions from the collapse of O-Ne-Mg cores suggest that such SNe are powered by neutrino heating and by the neutrino-driven wind of the newly formed NS, similar to what Woosley \\& Baron (1992) found in case of the accretion-induced collapse of white dwarfs to NSs. Such events have a low explosion energy ($\\sim\\,$0.1$\\times 10^{51}\\,$erg) and produce little $^{56}$Ni ($\\la 10^{-2}\\,M_\\odot$). Most of the ejecta expand initially with velocities of 2--4$\\times 10^4\\,$km$\\,$s$^{-1}$, a small fraction has nearly $10^{5}\\,$km$\\,$s$^{-1}$. This is significantly faster than in SNe of more massive progenitors. Of course, sweeping up the matter of the stellar mantle and envelope, the shock will decelerate, and the ejecta velocities after shock breakout from the stellar surface will be correspondingly lower. During the first second of the explosion, the ejected matter has $0.46\\la Y_{\\mathrm{e}}\\la 0.53$ and modest entropies ($10\\la s/(k_{\\mathrm{B}}/{\\rm{by}}) \\la 40$). Such conditions exclude that r-process elements are formed in this matter during this early phase of the explosion of O-Ne-Mg cores: $Y_{\\mathrm{e}}$ is too large for the ``classical'' low-entropy r-process and $s$ is too low for high-entropy r-processing. The ejecta in our models, however, do not show the vast overproduction of some very neutron-rich, rare isotopes like $^{87}$Kr, which was made in low-$Y_{\\mathrm{e}}$ material ($Y_{\\mathrm{e}}\\la 0.44$) in previous simulations (Mayle \\& Wilson 1988), and which was interpreted as a severe constraint to the rate of such events. Our models yield considerably less energetic explosions than previous simulations and show other significant differences in the dynamics and explosion characteristics. These are probably mainly linked to the improved treatment of neutrino transport and neutrino-matter interactions.  The small explosion energy obtained in our simulations is more consistent than previous explosion models of O-Ne-Mg cores with the low present expansion velocities ($\\sim\\,$1500$\\,$km$\\,$s$^{-1}$) of the filaments of the Crab remnant of SN~1054 (Davidson \\& Fesen 1985), corresponding to a low kinetic energy of 0.6--1.5$\\times 10^{50}\\,$erg for an ejecta mass of $4.6\\pm 1.8\\,M_\\odot$ in ionized and neutral gas (Fesen et al.\\ 1997). The energy could be even somewhat larger if there were several solar masses of material in an undetected, extended halo.  While our simulations are spherically symmetric, we do not expect any qualitative changes in the multi-dimensional case, and probably only a modest increase of the explosion energy. Since very fast outflow develops on a relatively short timescale after core bounce, nonradial hydrodynamic instabilities are unlikely to have time to merge and grow to very large structures or global asymmetry before the anisotropic pattern freezes out in the accelerating expansion (Scheck et al.\\ 2006, in preparation). Therefore the recoil velocity of the NS due to anisotropic mass ejection should remain fairly small (see Scheck et al.\\ 2004), in agreement with speculations by Podsiadlowski et al.\\ (2004). Corresponding 2D simulations are in progress.  Our models have also an important bearing on the nuclear EoS constraints deduced by Podsiadlowski et al.\\ (2005) from the low-mass Pulsar~B of the double pulsar J0737-3039, which has a gravitational mass of $M_{\\mathrm{G}} = 1.249\\pm 0.001\\,M_\\odot$. Provided the progenitor model we use is valid, the mass loss of the collapsing O-Ne-Mg core during the explosion leaves the neutron star with a baryonic mass of $M_0 = 1.36\\pm 0.002\\,M_\\odot$. The error range accounts approximately for variations associated with the employed EoS and the wind ablation after our simulations are terminated. Our value implies a systematic left shift and reduction of the ``acceptance rectangle'' in Fig.~3 of Podsiadlowski et al.\\ (2005). Combined with the recent measurement of a pulsar of 2.1$\\pm 0.2\\,M_\\odot$ in PSR$\\,$J0751+1807 (Nice et al.\\ 2005), which is the largest well determined NS mass so far, this lends viability only to a limited number of NS EoSs which allow for a sufficiently large maximum mass and whose $M_{\\mathrm{G}}/M_0$-curves pass through the acceptance rectangle.  Based on our findings one might speculate that (B-) stars around $9\\,M_\\odot$ are also the progenitors of some of the subluminous Type II-P supernovae mentioned in the introduction. In fact, their peculiarities would be explained in a very natural way. The low peak luminosity and extended plateau phase could result from the combination of a low hydrogen envelope mass ($\\simeq\\,$6$\\,M_\\odot$) with low expansion velocities ($\\la\\,$3000 km/s). The small mass of radioactive $^{56}$Ni would explain the low tail-luminosity of these objects. An alternative interpretation of subluminous Type II-P supernovae is the explosion of rather massive stars with extended envelopes, but otherwise more ``normal'' explosion energies. This connection is supported by the long duration of the plateau phase of many of these events (cf.\\ Pastorello et al.\\ 2004, and refs.\\ therein). Provided the observations cover the full duration of the plateau, SN~1997D and 2003gd may still be viable cases for explosions of stars with main sequence masses around 9 $M_\\odot$ (Hendry et al.\\ 2005). There is a clear difference between such stars and more massive supernova progenitors. The former eject very little amounts (some $10^{-3}\\,M_\\odot\\,$?) of oxygen only, like SN~1997D (Chugai \\& Utrobin 2000), whereas the latter produce up to a solar mass or more. Therefore one might be able to distinguish between the two scenarios on the basis of observations by measuring the oxygen lines in the late nebular spectra.  {\\it Acknowledgements.} We are grateful to K.~Nomoto for providing us with the initial data and for encouraging discussions. We thank R.~Buras and M.~Rampp for their input to this work, A.~Marek especially for the W\\&H nuclear EoS, and P.~Mazzali, A.~Pastorello, D.~Sauer and V. Utrobin for very helpful discussions about SN observations. This work was supported by the Sonderforschungsbereich 375 ``Astroparticle Physics'' of the Deutsche Forschungsgemeinschaft and by the International Max Planck Research School (IMPRS). Supercomputer time at the Rechenzentrum Garching is acknowledged.   \\nocite{maywil88} \\nocite{wanom03} \\nocite{whecowhi98} \\nocite{sumter01} \\nocite{hilnom84} \\nocite{burla85} \\nocite{barcoka87} \\nocite{ramjan02} \\nocite{mnys80} \\nocite{miynom87} \\nocite{nom82} \\nocite{nom83} \\nocite{nom84} \\nocite{nom87} \\nocite{hainom93} \\nocite{frybeheco99} \\nocite{wooba92} \\nocite{betwil85} \\nocite{arn75} \\nocite{gott70} \\nocite{dav82} \\nocite{lan03} \\nocite{tak89} \\nocite{past05} \\nocite{hil78} \\nocite{hil82} \\nocite{woo80} \\nocite{latswe91} \\nocite{ham03} \\nocite{past04} \\nocite{chu00} \\nocite{hor02} \\nocite{hart85} \\nocite{mar05} \\nocite{bur05} \\nocite{jan99} \\nocite{fro04} \\nocite{hen05} \\nocite{zam03} \\nocite{fes97} \\nocite{dav85}  {\\small"
    },
    "0512/astro-ph0512253_arXiv.txt": {
        "abstract": "We have computed, based on 17 infrared radial velocities, the first set of orbital elements for the M giant in the symbiotic binary V2116 Ophiuchi.  The giant's companion is a neutron star, the bright X-ray source GX~1+4.  We rule out the previously proposed period of 304 days, and instead, find an orbital period of 1161 days, by far the longest of any known X-ray binary.  The orbit has a modest eccentricity of 0.10 with an orbital circularization time of $\\lesssim5 \\times 10^6$ years. The large mass function of the orbit significantly restricts the mass of the M giant.  Adopting a neutron-star mass of 1.35 M$_{\\sun}$, the {\\it maximum} mass of the M giant is 1.22 M$_{\\sun}$, making it the less massive star.  Spectrum synthesis analysis of several infrared spectral regions results in slightly subsolar abundances for most metals.  Carbon and nitrogen are in the expected ratio resulting from the red-giant first dredge-up phase.  The lack of $^{17}$O suggests that the M-giant has a mass less than 1.3 M$_{\\sun}$, consistent with our maximum mass.  The surface gravity and maximum mass of the M giant result in a radius of 103 R$_{\\sun}$, much smaller than its estimated Roche lobe radius.  Thus, the mass loss of the red giant is via a stellar wind.  These properties argue that the M giant is near the tip of the first-ascent giant-branch.  Although the M giant companion to the neutron star has a mass similar to the late-type star in low-mass X-ray binaries, its near-solar abundances and apparent runaway velocity are not fully consistent with the properties of this class of stars. Thus, in many ways this symbiotic and X-ray binary system is unique, and various scenarios for its possible evolution are discussed. ",
        "introduction": "The star \\object[V2116 Oph]{V2116 Ophiuchi} = GX 1+4 [$\\alpha$ = 17$^h$~32$^m$~02\\fs16 $\\delta$ = $-$24\\arcdeg~44\\arcmin~44\\farcs0 (2000), $V$ = 18.4 mag] is unique among currently known symbiotic stars because the companion to the late-type giant is a neutron star rather than a white dwarf or main sequence star.  \\citet{letal71} identified the object as a very bright, hard X-ray source, GX~1+4.  \\citet{d76} found GX~1+4 to be an X-ray variable with a period of two minutes, suggesting that it is a slow pulsar.  Early observations of its period changes led to the suggestion that the variations resulted from Doppler shifts associated with orbital motion \\citep{betal76}.  However, it was later determined that the spin up and spin down of the neutron star resulted from variable mass accretion \\citep{detal81, retal82, metal88}.  While the present paper is only indirectly concerned with the neutron star at the heart of GX~1+4, we note that it is an extraordinary object.  The rate of change of its rotation period is the greatest of any known pulsar, and the magnetic field, perhaps as large as B$_\\circ$\\,$\\sim$\\,3$\\times$10$^{13}$G, is among the largest measured in any astronomical object \\citep{c97}.  Confirmation of the GX~1+4 optical counterpart with an 18th visual magnitude star came from both a precise ROSAT position \\citep{pfs95} and the detection of optical pulsing and flickering from the candidate star \\citep{jetal97}.  \\citet{gf73} had previously shown that the suspected optical counterpart is an infrared source.  The system is situated in the direction of the galactic center (l = 1.94\\arcdeg, b = 4.79\\arcdeg), and is significantly reddened, with $A_v$\\,$\\sim$\\,5 mag \\citep{detal77, setal96, cr97}.  Distance estimates, depending on the evolutionary state of the red giant, range from three to 15 kpc \\citep{cr97}.  Based on optical spectra, \\citet{detal77} suggested that GX~1+4 = V2116~Oph\\footnote{ In this paper we will refer to the symbiotic system as V2116~Oph and reserve the name GX~1+4 for the X-ray source.  However, V2116~Oph is classified by the GCVS as a variable X-ray source and the names are fully synonymous.} \\citep{ketal81} is a symbiotic star.  They also noted that the emission-line strengths indicate the presence of a circumstellar envelope of radius $\\sim$860 R$_\\odot$ ($\\sim$6$\\times$10$^{8}$ km).  In addition, \\citet{detal77} concluded that if the orbital separation is the size of the circumstellar envelope, the orbital period would be years.  \\citet{setal96} observed the optical spectrum at a resolving power of $\\sim$800 and found strong emission lines of hydrogen, neutral helium, and neutral oxygen as well as O~III, Fe~VII and Fe~X forbidden emission.  Absorption bands of TiO and VO are conspicuous in the red, leading to an M5~III spectral type classification for the late-type star in the system \\citep{setal96, cr97}.  Because of the large amount of reddening and the energy distribution of the M giant, V2116~Oph is much brighter in the infrared, $K$\\,$\\sim$\\,8~mag, than at optical wavelengths.  \\citet{betal97} and \\citet{cvkl98} observed the near-infrared spectrum of V2116~Oph at resolving powers of up to 4000. In contrast to the complex optical spectrum, except for strong emission lines from H~I Brackett $\\gamma$ and He~I 10830 \\AA\\, the infrared spectrum is typical of a cool M giant.  \\citet{cvkl98} interpreted the He~I line as an outflow of $\\sim$300 km~s$^{-1}$ from the M giant. However, \\citet{ketal99} noted similar line widths for optical-region emission lines and argued that the broad emission lines arise in gas that is gravitationally bound to the neutron star.  Binary systems containing a red giant are required to have a long orbital period.  \\citet{detal81} noted that a period of at least 120 days is necessary for {\\it a}~sin~{\\it i} to exceed the stellar radius, which places V2116~Oph in a unique X-ray binary niche.  There have been several attempts to determine the orbital period by making the assumption that mass transfer from the M giant to the neutron star is enhanced at periastron.  This mass transfer will alter the rotation period of the neutron star, and hence the pulse rate, so a statistical analysis of the X-ray pulses should show peaks at the orbital period. From such analyses \\citet{cetal86} and \\citet{petal99} concluded that the orbital period is $\\sim$304 days.  In this paper we exploit the typical M-giant near-infrared spectrum first described by \\citet{betal97}.  Our high-resolution, infrared spectra of V2116~Oph span nearly five years, enabling us to determine the first spectroscopic orbit of the red giant.  We also present an evaluation of the surface abundances of the M giant.  The neutron star in the V2116~Oph binary system is the remnant of a supernova.  The survival of the binary is remarkable given that the supernova must have been catastrophic for both the star we now see as the M giant as well as the surviving remnant.  Our results enable us to further investigate the properties of the system and improve our understanding of the red giant member as well as the evolution of this neutron star binary.  \\section {OBSERVATIONS AND REDUCTIONS}  Our 17 spectroscopic observations of V2116~Oph were obtained with four telescopes at three different observatories (Table~1).  From 1999 July through 2001 March we obtained spectra at the Kitt Peak National Observatory (KPNO) with the Phoenix cryogenic echelle spectrograph at the f/15 focus of both the 2.1\\,m and the 4\\,m Mayall telescopes.   A complete description of the spectrograph can be found in \\citet{hetal98}.  The widest slit was used, which results in a resolving power of $\\sim$50,000.  The KPNO observations were centered at several wavelengths between 1.55~$\\mu$m and 1.66~$\\mu$m.  From 2001 May through 2002 August spectra were obtained in the southern hemisphere with the 1.88\\,m telescope and coud\\'e spectrograph system at the Mt. Stromlo Observatory (MSO).  The detector was an infrared camera, NICMASS, developed at the University of Massachusetts.  We obtained a 2-pixel resolving power of 44,000 at a wavelength of 1.63~$\\mu$m.  A more complete description of the experimental setup may be found in \\citet{jetal98} as well as in \\citet{fetal00}.  The detector and electronics were previously used for our survey of northern symbiotics, carried out with the coud\\'e feed telescope at KPNO.  The devastating bush fire of 2003 January destroyed both the 1.88\\,m telescope at MSO and our infrared NICMASS camera.  As a result, from 2003 April through 2004 April we obtained additional observations with the 8\\,m Gemini South telescope, El Pachon, Chile, and the Phoenix cryogenic echelle spectrograph.  The first and last spectra had spectral resolving power of $\\sim$70,000, while the other two had resolving power of $\\sim$50,000.  Although V2116~Oph was very challenging to find and observe at KPNO and MSO, observations with the Gemini South telescope were straightforward because the target was acquired with the help of a guide star.  At $K\\sim$~8 mag V2116~Oph is a relatively bright star for an 8\\,m telescope.  The spectra obtained with Phoenix on Gemini South are of excellent quality and were useful in abundance analysis as well as the determination of the orbit.  Standard observing and reduction techniques were used \\citep{j92}. Wavelength calibration posed a challenge, because the spectral coverage was far too small to include a sufficient number of ThAr emission lines for a dispersion solution.  Our approach was to utilize absorption lines in a K~III star to obtain a dispersion solution.  Several sets of lines were tried, including CO, Fe~I, and Ti~I.  These groups all gave consistent results.  Radial velocities of the program star were determined with the IRAF cross-correlation program FXCOR \\citep{f93}.  The reference star was $\\delta$~Oph, an M-giant IAU velocity standard, for which we adopted a radial velocity of $-$19.1 km~s$^{-1}$ from the work of \\citet{sbf90}. ",
        "conclusions": "The determination of a time series of infrared radial velocities of the M giant V2116~Oph have confirmed the association of this star with the neutron star/X-ray source GX 1+4.  The single-line spectroscopic binary orbit is now well determined and at 1161 days exceeds the orbital period of any other low-mass X-ray binary by a factor of nearly 50 \\citep{lph01}.  While the inclination is unknown, based on other X-ray binary systems the mass of the neutron star is probably $\\sim$ 1.35 M$_{\\sun}$.  The plane of the orbit must lie close to the line of sight, implying that eclipses of at least part of the mass transfer flow or accretion disk are possible.  The next predicted date of ``mid-eclipse'' is 2008 Mar 31.  The abundances of the M giant, as well as the mass function of the orbit combined with the adopted mass of the neutron star, indicate a mass of $\\sim$ 1.2 M$_{\\sun}$ for the M giant.  The abundances of the M giant also suggest that it is a first ascent giant, in agreement with arguments presented in the detailed study of the system by \\citet{cr97}.  Our radius of the M giant and the X-ray behavior of the system indicate that the giant does not fill its Roche lobe, a conclusion consistent with its identification as a first ascent giant.  The fact that the giant does not fill its Roche lobe, makes V2116~Oph very different from other low-mass X-ray binaries.  The evolutionary history of this and other neutron star binaries is an area of continued speculation.  Should the neutron star have originated with the supernova explosion of a $\\sim$10 M$_{\\sun}$ star, the neutron star is ancient, having co-existed with the progenitor of the current M giant for nearly 5 billion years.  However, the activity of the V2116~Oph neutron star strongly suggests that it was formed 10$^7$ years ago or less.  Since we know the age of the M giant is $\\sim$5 x 10$^9$ years, this implies that a massive star never existed in this system.  This strongly suggests that the neutron star resulted from accretion induced collapse.  The pre-collapse system consisted of a white dwarf and the first ascent giant nearly filling its Roche lobe.  If the accretion induced collapse scenario is correct, the V2116~Oph neutron star is a fairly recent remnant from a single degenerate type I supernova event.  To an observer prior to the V2116~Oph supernova, the white dwarf -- red giant system would have appeared as a symbiotic system.  There are few symbiotic systems known that contain high mass white dwarfs and none that also contain a companion close to filling its Roche lobe.  However, the fraction of symbiotic binaries that have been fully characterized is small.  Continued exploration of these systems may reveal a pre-accretion induced collapse system."
    },
    "0512/astro-ph0512409_arXiv.txt": {
        "abstract": "{} {To investigate the stability and angular momentum transport by the strato-rotational instability in the nonlinear regime. }{ The hydrodynamic compressible equations are solved in a cartesian box in which the outer cylinder is embedded. Gravity along the rotation axis leads to density stratification. No-slip boundary conditions are used in the radial direction, while free-slip conditions are used on the two ends of the cylinders. }{ The strato-rotational instability is confirmed and the Reynolds stress is shown to transport angular momentum away from the axis. However, the growth rate decreases with increasing Reynolds number. This, as well as the presence of boundaries renders this instability less relevant for astrophysical applications. }{} ",
        "introduction": "The magneto-rotational instability (MRI) is now commonly invoked in order to explain the origin of turbulence in accretion discs (Balbus \\& Hawley 1998). This is because there exists no purely hydrodynamic local instabilities in accretion discs in the sense that periodic shearing box calculations do not yield instability. Already when vertical shear is included, there is a linear instability, but for thin discs its growth rate is small compared with that of the MRI (Urpin \\& Brandenburg 1998; Arlt \\& Urpin 2004). Purely hydrodynamic instabilities are of interest for protostellar discs where the conductivity is usually so poor that magnetic effects are unimportant, and hence the MRI may be irrelevant. Another possibility for a hydrodynamic instability is the nonlinear shear instability (Chagelishvili et al.\\ 2003, Afshordi, Mukhopadhyay \\& Narayan 2005). The general difficulty with such models has been discussed by Balbus, Hawley \\& Stone (1996). However, a conclusive resolution of this problem using simulations remains difficult because of the large Reynolds numbers required.  In an attempt to address the possibility of hydrodynamic instabilities in accretion discs, one can resort to experiments. All these experiments have boundaries, so their relevance to accretion discs is questionable. Nevertheless, such experiments remain interesting in their own right. The possibility of a nonlinear instability of a Taylor-Couette flow experiments has recently been addressed by Richard \\& Zahn (1999). The issue was revived by recent findings of Molemaker, McWilliams \\& Yavneh (2001) who found the possibility of a {\\it linear} instability in Taylor-Couette flow with density stratification along the axis. Dubrulle et al.\\ (2005) confirmed this instability for a range of different radial boundaries conditions.  More recent work by Shalybkov \\& R\\\"udiger (2005), who considered regular Taylor-Couette flow in the presence of no-slip boundaries, gave a detailed stability criterion for the instability, which is henceforth referred to as the strato-rotational instability (SRI). Withjack \\& Chen (1974) observed nonaxisymmetric patterns in density-stratified Taylor-Couette flows close and  beyond the Rayleigh line.  Umurhan (2005) shows that, within the framework of the quasi-hydrostatic semi-geostrophic approximation, the instability only survives in the presence of no-slip radial boundaries. The purpose of the present paper is to demonstrate the existence of this instability using direct simulations and to calculate its saturation level in the nonlinear regime. In contrast to nonlinear instabilities, the SRI can be studied conveniently by simulations, because the critical Reynolds numbers are well below $10^3$. We are thus also able to compute sign and, in the nonlinear regime, also the amount of angular momentum transport due to the SRI. Let us emphasize here that the sign of angular momentum transport (inwards or outwards) is nontrivial. There are several examples in the literature (mostly for rotating convection) where the cross correlation \\EQ Q_{r\\phi}=\\overline{u'_ru'_\\phi} \\label{Q1} \\EN of the one-point correlation tensor \\EQ Q_{ij}=\\overline{u'_i(\\xx,t)  u'_j(\\xx,t)} \\label{Q2} \\EN turns out to be negative, i.e.\\ angular momentum is transported inwards by the rotating turbulence (see Ryu \\& Goodman 1992; Kley, Papaloizou \\& Lin 1993; Stone \\& Balbus 1996). The  primes in (\\ref{Q1}) and (\\ref{Q2}) denote the departures from the toroidally averaged flow, i.e.\\ $\\uu'=\\uu-\\meanuu$, where \\EQ \\meanuu={\\textstyle\\int_0^{2\\pi}}\\uu\\,\\dd\\phi/2\\pi. \\EN Here we shall find that the zonal flux of angular momentum,  (\\ref{Q1}), always proves to be positive  i.e.\\ the angular momentum  is transported outwards by the SRI. ",
        "conclusions": "In the present work we have explored the nonlinear saturation of the strato-rotational instability using a finite difference method. We verify the stability criterion of the Taylor-Couette flow in the presence of vertical stratification. In agreement with earlier results by Shalybkov \\& R\\\"udiger (2005), a nonaxisymmetric instability is found when the azimuthal velocity on the inner cylinder, $R_{\\rm in}\\Omega_{\\rm in}$, exceeds that on the outer cylinder, $R_{\\rm out}\\Omega_{\\rm out}$. This leads to instability in a regime that would be stable by the Solberg-H{\\o}iland criteria, which only apply to axisymmetric flows on in the tight-winding approximation (R\\\"udiger, Arlt \\& Shalybkov 2002).  Increasing $\\mbox{Re}$ sufficiently beyond the critical value for the instability, the growth rate decreases with increasing $\\mbox{Re}$. This is related to the fact that the flow is driven via the viscous force, which is obviously not relevant in astrophysics, but it is relevant to the experiments that are currently being designed to address these questions that are astrophysically motivated in other ways."
    },
    "0512/astro-ph0512123_arXiv.txt": {
        "abstract": "Primordial stars are likely to be very massive $\\geq30\\Msun$, form in isolation, and will likely leave black holes as remnants in the centers of their host dark matter halos in the mass range $10^{6}-10^{10}\\Ms$. Such early black holes, at redshifts z$\\gtsim10$, could be the seed black holes for the many supermassive black holes found in galaxies in the local universe. If they exist, their mergers with nearby supermassive black holes may be a prime signal for long wavelength gravitational wave detectors.  We simulate formation of black holes in the center of high redshift dark matter halos and explore implications of initial natal kick velocities conjectured by some formation models. The central concentration of early black holes in present day galaxies is reduced if they are born even with moderate kicks of tens of km/s. The modest kicks allow the black holes to leave their parent halo, which consequently leads to dynamical friction being less effective on the lower mass black holes as compared to those still embedded in their parent halos. Therefore, merger rates may be reduced by more than an order of magnitude.  Using analytical and illustrative cosmological N--body simulations we quantify the role of natal kicks of black holes formed from massive metal free stars on their merger rates with supermassive black holes in present day galaxies. Our results also apply to black holes ejected by the gravitational slingshot mechanism. ",
        "introduction": "It is firmly established that most galaxies have super-massive black holes (SMBH) at their centers. Their masses are in the range 10$^6 \\, \\Msun \\ltsim M \\ltsim 10^9 \\, \\Msun$ and it appears that there is a number of correlations (Kormendy $\\&$ Gebhardt 2001, Gerhard 2001) between their masses and properties of the galactic bulge hosting them: mass of the SMBH on one hand and mass (Laor 2001) or luminosity or velocity dispersion (Merritt $\\&$ Ferrarese 2001, Tremaine et al. 2002, Gebhardt et al. 2001) or light profile (Graham et al. 2001) of the galactic bulge on the other. These correlations point to the link between the formation of SMBHs and the evolution of their hosts. It also appears that SMBHs are linked to the properties of the host dark matter halos. If the SMBH precursors have been present from very early on, then their mergers, together with growth by accretion, could account for the abundance of the SMBHs today (Schneider et al. 2002).  Ab initio numerical simulations of the formation of the first luminous objects in the current structure formation models find metal free stars to form in isolation and may have masses $30\\,\\Msun \\leq m\\leq 300 \\,\\Msun$ (Abel et al. 2000, Abel et al. 2002, Bromm et al. 2002). In current models of structure formation, dark matter initially dominates and pregalactic objects form because of gravitational instability from small initial density perturbations. As they assemble via hierarchical merging, the metal-free primordial gas cools through rotational lines of hydrogen molecules. In the absence of metals, H$_2$ is the main coolant below $\\sim$ 10$^4$K, which is the temperature range typically encountered in collapsing Population III objects. Rotational transitions of H$_{2}$, occurring via electric quadrupole radiation, allow the gas to cool and sink to the center of the dark matter potential well. This leads to a top-heavy initial stellar mass function and to the production of very massive stars, unlike the modern stellar IMF which is declining rapidly with increasing mass.  Metal-free, high mass Population III stars die after 3 Myr, losing only a small fraction of their mass. If the fragmentation occurs during core collapse, two or more compact objects (black hole or neutron star) could be produced. This suggests that a large population of primordial massive black holes (MBH from now, with mass $\\leq 10^3 \\Msun$) could be an end product of such pregalactic star formation (Heger et al. 2003). Since they form in rare high-$\\sigma$ density peaks (Couchman $\\&$ Rees 1986, Madau $\\&$ Rees 2001), relic MBHs would be predicted to cluster in the cores of more massive halos (Abel et al. 2002) formed by subsequent mergers. It has been suggested (Volonteri et al. 2003, Islam et al. 2003) that mergers of massive halos and clustering of these MBHs should start at redshifts as large as z$\\sim$20. Growth of MBHs would proceed through accretion and coalescence which would lead to the formation of intermediate mass black holes (IMBH from now, with mass in range $10^3\\,\\Msun \\leq m\\leq 10^6 \\,\\Msun$). IMBHs can form sufficiently early on. In the most optimistic scenario (Haiman $\\&$ Loeb 2001), 10$^9\\Msun$ black hole could form at redshift z$\\sim$10. However, through the mergers of dark matter halos at high redshift, black holes would form binaries and their coalescence under gravitational radiation would give them a significant kick velocity (Favata et al. 2004, Merritt et al. 2004). Anisotropic emission of gravitational waves which carry away linear momentum causes center of mass recoil. The recoil velocities are expected to be most likely in the range 10-100 kms$^{-1}$, kicks of a few hundred kms$^{-1}$ are not unexpected, and the largest recoils should not be above 500 kms$^{-1}$ (Favata et al. 2004). The amplitude of the kick determines if the black hole will be ejected from its host dark matter halo or if it will fall back. Even the most massive dark matter halo at redshift z$\\geq$11 can not retain the black hole that received 150$\\kms$ kick while at the redhift of z$\\geq$8, kicks of 300$\\kms$ are sufficient to produce the same result (Merritt et al. 2004, Fig 3.). This is valid for black holes with masses $10^3\\,\\Msun \\leq m\\leq 10^8 \\,\\Msun$. Fully general relativistic simulations of radiation recoil from highly distorted Schwarzschild holes yield maximum kick velocities in excess of 400\\,km/s (Brandt $\\&$ Anninos 1999).  This should be sufficient to: eject IMBHs from globular clusters; displace IMBHs from the centers of dwarf galaxies; perhaps provide sufficient population of IMBHs for merging at the centers of spiral galaxies (Merritt et al. 2004). In fact, for Kerr black holes, recoil is significant even for holes of equal mass (Favata et al. 2004) and may be directed out of the plane of the orbit (Redmount $\\&$ Rees 1989). Kicks are also expected from the gravitational slingshot - the ejection of one or more black holes when three black holes interact (Xu et al. 1994, Valtonen et al. 2000, Rees 1988). Space velocities much greater than those of their progenitores are common in neutron stars (Lai et al. 2001, Colpi $\\&$ Wasserman 2002, Arzoumanian et al. 1997). Recent studies of pulsar proper motion implies characteristic velocities of the order of 200-500 kms$^{-1}$ (Wex et al. 2000), and as large as 100$\\kms$ for black holes (Colpi $\\&$ Wasserman 2002). The high space velocity of X-ray Nova Sco 1994 implicates black hole kicks (Brandt et al. 1995).  Likely values of the velocity of the merged objects at infinity lie in the range of 100 - 300 $\\kms$ (Davies et al. 2002).  Without exploring the nature of kicks, we are investigating the observational implications of other peoples conjectures on black holes getting impulsive kicks, whether on formation or through mergers. We are also exploring the consequences that kicks will have on the IMBHs merger rates and the formation of SMBHs that are indirectly observed in the nuclei of nearby luminous galaxies. LISA will test the Population III IMBHs formation models.  A powerful instrument for studying IMBHs will be LISA (Danzmann 2003). LISA can study much of the last year of inspiral, as well as the waves from the final collision and coalescence of IMBHS binaries, whenever the masses of the black holes are in the range $3\\times 10^4 \\, \\Msun \\leq M \\leq 10^8 \\, \\Msun$. It can also study the final coalescences with remarkable signal to noise ratios: S/N$\\geq$1000. For black holes with the masses in the range 100$\\Msun\\leq$M$\\leq10^4\\Msun$, out to cosmological distances, LISA can observe the last few years of inspiral, but not the final collisions. The equal-mass black hole binaries enter LISA's frequency band roughly 1000 years before their final coalescences, more or less independently of their masses, for the range 100$\\Msun\\leq$M$\\leq10^6\\Msun$. If the coalescence rate is one per year, LISA would see roughly 1000 additional IMBHs binaries that are slowly spiraling inward, with measurable inspiral rates. From the inspiral rates, the amplitudes of the two polarizations, and the waves' harmonic content, LISA can determine each such binary's luminosity distance, redshifted chirp mass (1+z)M$_c$, orbital inclination, and eccentricity. From the waves' modulation by LISA's orbital motion, LISA can learn the direction to the binary with an accuracy of order one degree (Thorne 1995, Cornish $\\&$ Levin 2002, Vecchio et al. 2004).  In addition to the mapping of spacetime around SMBH with inspiraling compact objects, BBO (Big Bang Observer) mission will have arcsecond-arcminute precision in identifying and localizing every merging black hole binary at any redshift, anywhere in the universe over the years to months before their actual merger. In combination with ground-based detectors at higher frequencies, it will measure the mass, angular momenta, and dynamic spacetime structure of these compact objects with unprecedented precision.  If these IMBHs would accrete from the interstellar medium they may also be found by deep X--ray observations (Agol $\\&$ Kamionkowski 2002) in the Milky Way. This observational signature will depend sensitively on the spatial distribution of these black holes in our galaxy. Similarly, the next generation of micro-lensing searches may be able to constrain high mass black holes from the longest events (Agol et al. 2002).   In our numerical simulations we use GADGET (GAlaxies with Dark matter and Gas intEracT), a code written by Volker Springel (Springel et al. 2001), for simulations of self-gravitating collisionless fluids evolution, with a tree method N-body approach, and a collisional gas using smoothed particle hydrodynamics.  In \\S 2 we present straightforward analytical arguments. In \\S 3 we describe the simulations setup, the codes used for analysis of the GADGET simulations, and present results for 3 different high resolution simulations with varying natal kick parameters for early black holes.  In \\S 4 we describe post-merger evolution through dynamical friction, before we conclude in \\S 5. ",
        "conclusions": "We have performed and analyzed high-resolution collisionless simulations of the evolution of structure in a $\\Lambda$CDM model. We have followed the formation and evolution of dark matter halos and by assuming that each halo is a host of one IMBH we have studied the formation of SMBH. We focused on two specific cases. First, IMBHs together with their host halos merge through dynamical friction. Second, when IMBHs are endowed with an initial kick, this leads to the ejection from their host halos in many cases. Analytically it is clear that the dynamical friction will act more efficiently on the host halos than on the much lower mass black holes formed within them. Our illustrative calculations highlight some of the expected differences in the density distribution of the final distributions of black holes which may be quite different, even in the presence of the modest kick velocities we have imposed.  In order for dynamical friction to work in numerical simulations, the density of the background - primary halo density - has to be more resolved. Larger resolution gives smaller softening lengths for particles and smaller softening lengths give more realistic dynamical friction. At the current resolution of our simulation, dynamical friction can not be efficiently realized in the model after z$\\sim$1 and r$\\ltsim$r$_{vir}/30$ because the background density of parent halos is not resolved well enough. Even at the current resolution, our simulations are able to account for effects of dynamical friction on the sub-halos at high z. Subsequent papers will track the dynamical evolution of the sub-halos and IMBH at late times, using higher resolution simulations and semi-analytic implementations of dynamical friction (Binney $\\&$ Tremaine 1987). A little over 4$\\%$ of IMBHs merge at the center in less than Hubble time. If this is the dominant way of creating SMBH then it is efficient even if the IMBHs are ejected from their parent halos, with merger numbers only reduced by a factor of two or so for modest kicks. It is also possible that kick velocities have been underestimated in our simulations. The value of the median kick of 150km/s may be increased to more then 1000km/s according to some authors (Favata et al. 2004, Merritt et al. 2004) which would also greatly enhance the effects introduced here, reducing the number of mergers too negligibly small. LISA observations should strongly constrain any natal kick on IMBHs formed from very massive Pop III stars in low mass proto-galactic sub-halos."
    },
    "0512/astro-ph0512315_arXiv.txt": {
        "abstract": "Goldreich-Sridhar model of incompressible turbulence provides an elegant approach to describing strong MHD turbulence. It relies on the fact that interacting Alfv\\'enic waves are independent and have random polarization. However, in case of strong interaction a spontaneous local assymetry can arise. We used direct numerical simulations to show that polarization alignment occurs and it grows larger at smaller scales. This effect would lead to a shallower spectrum and stronger anisotropy. Even small changes in these two properties will have important astrophysical consequencies, e.g. for the cosmic ray physics. ",
        "introduction": "Astrophysical plasmas is turbulent and magnetized. Magnetohydrodynamic turbulence affects many phenomena including the propagation and acceleration of cosmic rays (see Shlickeiser 2003). While foundations of the theory of anisotropic MHD turbulence model can be traced back to works in 80s (see Shebalin, Matthaeus \\& Montgomery 1983, Higdon 1984, Matthaeus \\& Brown 1988) a substantial recent progress was initiated by the pioneering study by Goldreich \\& Sridhar (1995, henceforth GS95). There a concept of balancing of linear and non-linear term in incompressible MHD equations was suggested, which results in a prediction between the wavenumbers parallel $\\|$ and perpendicular $\\bot$ to the local direction of magnetic field $k_{\\|}\\sim k_{\\perp}^{2/3}$. GS95 model also predicts Kolmogorov $E(k)\\sim k_{\\bot}^{-5/3}$ spectrum, which is consistent with both interstellar (see Armstrong, Rickett \\& Spangler 1995) and Solar wind (see Horbury 1999) data.  While the core concept of the GS95 model, namely, critical balance, was confirmed in numerical simulations (Cho \\& Vishniac 2000, Maron \\& Goldreich 2001, Cho, Lazarian \\& Vishniac 2002, henceforth  CLV02), the simulations revealed a difference in the spectral indexes obtained. For instance, Maron \\& Goldreich (2001) obtained the spectral index that is close to $-3/2$ in contrast to $-5/3$ in CLV02. Index of $-3/2$ for field-perpendicular energy was obtained in simulations of Muller \\& Grappin (2005). This stimulated theoretical efforts to understand the difference between the theory and simulations on one hand and between different sets of simulations on the other hand. In particular, Boldyrev (2005) proposed a model in which the interactions are weakened compared to the GS95 predictions as a result of 3D structure of eddies. An alternative point of view, namely, flattening of the spectrum as the result of intermittency was discussed in Maron \\& Goldreich (2001).  We note, that the problem of MHD turbulence spectrum is of great practical importance. For instance, the introduction of GS95 spectrum and its extension to compressible media (see Lithwick \\& Goldreich 2001, Cho \\& Lazarian 2002, 2003) resulted in a substnatial changes in understanding of cosmic ray propagation and acceleration (see Chandran 2000, Yan \\& Lazarian 2002, 2004, Cho \\& Lazarian 2006).  In what follows we address the issue of interaction efficiency using numerical simulations. We use both compressible and incompressible MHD runs. Contrary to some earlier claims, Cho \\& Lazarian (2002) demonstrated weak coupling of the compressible and Alfvenic parts of the spectra, which justifies comparing Alfvenic turbulence in incompressible and compressible media. ",
        "conclusions": "In our incompressible simulations magnetic field energy was typically slightly larger that kinetic energy (see Fig 1.) This difference is the so-called residual energy (see Muller, Grappin, 2005). Residual energy creates a polarization alignment, but of different type that we measure. In this case ${\\bf w}$ and ${\\bf u}$ are systematically {\\it antiparallel}. However, we are looking not for difference between p.d.f. of angles of 0 and 180 degrees, but for difference for 0 and 90 degrees. It would be a second-order correction and is generally pretty small, as long as kinetic and magnetic energies are close. We estimated that in our case this spurious effect is never larger than 2\\%.  Maron and Goldreich (2001) observed axial asymmetry or net polarization of Alfv\\'en waves. They claimed that this effect is strongest in decaying turbulence, but in forced turbulence it is bound. They did not, however, connected this effect to the flatter than $-5/3$ spectra observed in both decaying and forced runs. They speculated that spectra could be flatter due to intermittency.  Boldyrev (2005) considers weakening of interaction that comes from three-dimensional structure of the eddy. As the model he uses same Kolmogorov-type arguments, so his formulae can be reproduced by ours by taking $\\alpha=\\alpha'/(3+\\alpha')$, where $\\alpha'$ is Boldyrev's $\\alpha$. However, in our simulations we saw that purely geometrical arguments can not fully describe weakening and the real weakening is due to correlation between angular alignment and the wave strength.  Somewhat unexpected result is that compressible cubes show more of a polarization alignment, even though they naturally have a smaller inertial range due to large numerical diffusivity and viscosity. So far within our approach $\\alpha$ is not constrained and we do not claim that weakening will nessesarily behave as a power-law. We leave this two issues to further study.  While the magnitude of the changes of the spectral index as well as anisotropy that we observe is not large, the potential implications of this can be very substantial. Turbulence in interstellar medium is injected on the scale of dozens of parsecs (see Lazarian \\& Pogosyan 2000, Farmer \\& Goldreich 2004), and on the megaparsec scales in clusters of galaxies (see Cassano \\& Brunetti 2005). The scale at which Alfven turbulence dissipates can be of the order of thermal proton Larmor radius, i.e. thousands of kilometers. As the result of such a humongous scale separation any changes in the spectra and scale-dependent anisotropy will have important consequencies for the turbulent energy available at a sufficiently small scale. Cosmic ray acceleration and propagation are the processes that are directly affected."
    },
    "0512/astro-ph0512137_arXiv.txt": {
        "abstract": "Hydrodynamical simulations of semi-detached, polytropic binary stars are presented in an effort to study the onset and stability of dynamical mass transfer events. Initial, synchronously rotating equilibrium models are constructed using a self-consistent-field technique and then evolved with an Eulerian hydrodynamics code in a fully self-consistent manner. We describe code improvements introduced over the past few years that permit us to follow dynamical mass-transfer events through more than 30 orbits. Mass-transfer evolutions are presented for two different initial configurations: A dynamically unstable binary with initial mass ratio (donor/accretor) $q_0 = 1.3$ that leads to a complete merger in $\\sim 10$ orbits; and a double-degenerate binary with initial mass ratio $q_0 = 0.5$ that, after some initial unstable growth of mass transfer, tends to separate as the mass-transfer rate levels off. ",
        "introduction": "Most types of binary stars observed today in a long-term interacting phase of their evolution --- such as cataclysmic variables, X-ray binaries and contact binaries --- and other types that are no longer actively transferring mass, are survivors of a dynamical phase of evolution in a more violent past. Such rapid evolutionary phases accompany the onset of mass transfer that occurs when angular momentum losses or stellar evolutionary processes bring one component of a detached system into contact with its Roche lobe. It is important to investigate the onset and stability of the ensuing mass transfer in order to understand the origin and the evolutionary links between the various types of binary systems that we observe today. The stability of a given binary system upon contact depends on the system's mass ratio, the structure of the donor star, and how mass and angular momentum is redistributed in the binary during the ensuing mass-transfer event. By making reasonable assumptions that encapsulate the essential physics while simplifying the governing equations, analytical treatments can be devised that provide insight into not only the stability, but the long-term evolutionary behavior of mass-transferring, semi-detached binaries \\citep{SoPh97, HaWe99, GPF}. But because rapid mass-transfer events, in particular, are intrinsically complex --- involving, for example, nonlinear supersonic flows and a gravitational field that changes in a nontrivial way in response to the dynamical exchange of matter between the stars --- the rapid phases through which many binaries evolve are not likely to be well understood until they are modeled hydrodynamically with the inclusion of as much physics as possible. In this paper, we demonstrate how large-scale numerical, hydrodynamic simulations can be employed to better understand the stability and the long-term evolutionary behavior of binaries as they undergo rapid phases of mass-transfer. Note that in this paper the term ``dynamical\" is used to mean processes that result in significant changes in the binary parameters and mass-transfer rates in a few orbital periods.  Among several possible mechanisms driving binary evolution, close binaries inevitably lose orbital energy and angular momentum by the emission of gravitational waves. Double degenerate white dwarfs (DWDs) are of particular interest since their expected numbers and rates of emission of gravitational waves are such that they are likely to constitute an important source of background noise for the LISA mission \\citep{HiBe00}. Details of the time-evolutionary behavior of the gravitational-wave signal that is emitted from any given system depend on the details of mass and angular momentum flow in the binary. Reliable estimates for the rates of emission and the time-frequency characteristics of the emitted signal (the ``template\") depend on having a good understanding of the above processes. While some progress can be made with analytical prescriptions, theoretically constructed templates are not likely to be useful for the extraction from the background of signals from individual sources until detailed hydrodynamical models can accurately simulate dynamical mass-transfer events in DWDs. Furthermore, it has been realized recently that many of the theoretically predicted and abundant detached DWDs that are driven to contact in a Hubble time, would start mass transfer in a direct impact mode \\citep{Maet04}. In this mode, the accretion stream hits the accretor's surface and dissipates without forming an accretion disk. Therefore the orbital angular momentum carried by the stream serves to spin up the accretor and is generally not returned by tides to the orbit. One would like to study these processes without the approximations usually adopted for mathematical convenience. Since relatively little is known about these direct impact binaries, it is also of interest to investigate the structure and properties of these flows and explore possible observational signatures.  The above considerations have motivated us to undertake a long-term program that utilizes detailed numerical hydrodynamic simulations to model dynamical phases of mass transfer that can arise when binaries of different types reach a semi-detached state. We have developed a hydrodynamical code that is capable of following such binaries through more than 30 orbits while conserving mass and angular momentum to a high level of precision. As a first step, we studied the evolution of detached polytropic binaries with components of equal and unequal mass (Motl, Tohline \\& Frank 2002, hereafter MTF). In the present paper we describe some important improvements that have been made to the code described by MTF and report on our first results of simulations that follow dynamical mass-transfer events in a fully self-consistent manner. We follow the evolution of two different polytropic binary systems whose properties were selected to illustrate the capabilities of our code and to provide some overlap with earlier, related studies. In particular, one of our models has been chosen specifically for comparison with a smooth particle hydrodynamics (SPH) simulation that was presented a decade ago by \\citet{RS95}. To our knowledge, this work has not been repeated with an Eulerian code in the intervening years.  Much of the related numerical simulation work that has preceded ours (see \\S\\ref{PreviousSimulations} for elaboration) has concentrated on stars having relatively stiff equations of state and on systems that encounter and evolve through a tidal instability.  The hydrodynamics of the coalescence, including brief phases of mass-transfer, has been modeled when necessary as a prelude to an almost inevitable merger. Our focus is, instead, on binaries whose components obey a relatively soft equation of state.  Our objective is to investigate in more detail the hydrodynamics of mass transfer, the structures arising from this transfer, and the role mass transfer plays during the dynamical phase, sometimes leading the binary closer to a tidal instability and merger, but sometimes saving it from merger and leading it to a slower evolutionary phase. ",
        "conclusions": "\\label{Conclusions}  This is the second in a series of papers that describe results of direct simulations of the dynamical evolution of unequal-mass binaries using a three-dimensional, finite-difference hydrodynamics technique.  Our long-term goal is to gain a better understanding of the origin and survival of various classes of binaries --- including double white dwarfs, contact binaries and direct-impact accretors --- by accurately simulating hydrodynamical flows that arise when binaries of various kinds become semi-detached. In contrast to this, much of the related analytic and numerical work that has been carried out over the past decade has placed an emphasis on the onset of the tidal instability and on the coalescence and merger of nearly identical neutron stars. In these earlier studies, mass transfer events were modeled when necessary as a prelude to the almost inevitable merger.  In fact, the inevitability of mergers seems to have been accepted also for double white dwarf binaries.  While this may well be the case with massive white dwarfs, the results presented here already suggest that low-mass WD binaries may escape merger and survive as AM CVn systems \\citep{TY96, NPVY01}.  In the present paper, after describing improvements that have been made in the hydrodynamics code that was developed by and described in MTF, we have taken the first steps toward elucidating the role that is played by mass transfer in the outcome of the dynamical phase of evolution following first contact. We have been especially concerned with those binary systems in which the mass ratios and equations of state are such that Roche lobe contact occurs before a tidal instability, so we have focused on binaries with $n=3/2$ ($\\gamma=5/3$) polytropic components \\citep{RS95, UE1, UE2}. We have presented two evolutions of a system having an initial mass ratio $q_0 = 1.323$ that was dynamically unstable to mass transfer. In one evolution, mass transfer was initiated by removing angular momentum from the binary (Q1.3-D) while in the other evolution we expanded the donor star causing it to overflow its Roche lobe (Q1.3-E). While there were subtle differences in the structure of the binary and the depth of contact by the time driving was terminated, in both cases the binary was dynamically unstable and a merger was the final result. In fact, we demonstrated that if one looks at the last two or three orbital periods before the merger, the behavior of both evolutions is consistent with being the same. We therefore conclude that, for unstable systems, it does not matter how the system gets into contact; once the instability kicks in, the behavior is the same. This result also has served as a convergence test for our hydrodynamics code since two evolutions with identical initial states but that were brought into contact in different ways and therefore had different detailed histories, ended up looking indistinguishable during the final stages leading to a merger.  At the very end of the ``mass-transfer'' phase of both the Q1.3-D and Q1.3-E evolutions, the binary transferred orbital angular momentum to the spin of both components catastrophically, as the stars plunged towards one another and eventually merged.  Overall the behavior we observed during the final few orbits agrees with the predictions of the tidal instability \\citep{LRS1, LRS2, LRS3, LRS4, LRS5}. This suggests that one can recognize two epochs during the evolution of a $q_0 > q_{\\rm stable}$ binary system that undergoes an episode of mass-transfer: Over an extended period of time, the rate of mass transfer grows while the orbital angular momentum changes very slowly. This instability is driven by an ever increasing depth of contact and thus can be considered a mass-transfer instability. Once the separation has been reduced sufficiently, the second instability sets in, during which the orbital angular momentum drops rapidly while the spin angular momenta increase driven by tides.   We have also presented detailed results from three mass-transfer evolutions (simulations Q0.5-Da, Db, and Dc) of a polytropic binary with an initial mass ratio $q_0=0.5$ and in which the two components had the same specific entropy. Thus this initial state could represent a DWD system with components having the same composition. It also corresponds to one system that was simulated by \\citet{RS95} using an SPH technique. In each of our simulations, the less massive star was driven into contact with its Roche lobe via the slow removal of orbital angular momentum, but in simulations Q0.5-Da and Q0.5-Db this artificial driving was turned off shortly after the mass-transfer event was initiated (after 2.7 and 5.3 initial orbital periods, respectively). Our results differ in some important ways from the results reported by RS95.  Most significantly, in the two simulations that were evolved for an extended period of time in the absence of systemic angular momentum losses, we do not observe a merger nor a tidal disruption; in evolution Q0.5-Da (Db) our binary survives for more than 30 (14) orbital periods and at the end it is separating while the mass-transfer rate has leveled off. Via these extended simulations, we have demonstrated that our numerical tools permit us to accurately model accretion flows in dynamically evolving mass-transfer systems.  We can, for example, analyze how angular momentum is exchanged between the orbit and the spin of the two stars during a phase of direct-impact accretion, and we can analyze how the structure of the accretor dynamically readjusts as relatively high specific angular momentum material is deposited onto its equatorial region. Studies of this type should assist in the determination of which simplifying assumptions are justified -- as well as which are not -- in models that attempt to describe extended phases of mass-transfer evolutions semi-analytically \\citep{WebIben, Maet04}. By conducting a variety of related simulations, we ultimately hope to be able to determine what the critical mass ratio $q_\\mathrm{stable}$ is that defines which binary systems are stable or unstable against mass-transfer.  We speculate that the outcomes of our Q0.5-Da and Q0.5-Db simulations were different from the Q0.5-RS simulation because the evolution presented in \\citet{RS95} was started from a significantly deeper initial contact, and thus transferred a larger fraction of the mass before separating. In RS95, the maximum density of the donor appears to be moving away and yet the separation calculated as the distance between the centers of mass of the material in their respective Roche lobes seems to be decreasing. Therefore, our results could be described as reproducing some of the initial features of the SPH simulation by \\citet{RS95} at a slower pace. Evidently, if the mass ratio of the binary is such that initially the system is unstable to mass transfer, the final fate of the binary will depend on whether during the initial phase of mass transfer the separation decreases sufficiently for the tidal instability to take over. In our Q0.5-Dc simulation, in which driving remained on throughout the evolution, it appears as though the system encounters the tidal instability; the last 2-3 orbits of this evolution resemble fairly closely the evolutionary behavior of the published Q0.5-RS simulation.  However, in our Q0.5-Da and Q0.5-Db simulations, once the driving is cut off the tidal effects only succeed in delaying the tendency of the binary to separate as mass transfer proceeds by temporarily storing some orbital angular momentum in the spin of the donor. Once the separation begins to increase, tides become ineffective and the binary avoids the merger.  Given an equation of state for the binary components, the mode of mass transfer, and the expected mass loss, if any, it is possible to make predictions about the evolutionary outcome after contact. However, our simulations suggest that the eventual fate of such a binary depends not only on the initial mass ratio $q_0$, but also on how far $q_0$ is above the appropriate $q_{\\rm stable}$, and even on the rate of driving unless this is very slow. In other words, the detailed outcome depends on the non-linear development of the mass-transfer and tidal instabilities. Only when $q_0$ is well above $q_{\\rm stable}$ will the evolution of mass transfer proceed rapidly enough to resemble qualitatively the predictions of the analytic solution derived in \\citet{WebIben}. When $q_0$ is only slightly above $q_{\\rm stable}$, non-linear effects come into play that make it possible for the system to survive the mass-transfer instability and avoid merger. We shall discuss these questions further in two forthcoming papers \\citep{GPF, MDTF}. As has already been mentioned, the tools we have developed will also enable us in the future to investigate in more detail the hydrodynamics of mass transfer and the structures arising from this transfer, transient flows, oscillations, mixing and convection. We see already some of these features in our simulations and we think they warrant further investigation."
    },
    "0512/astro-ph0512247_arXiv.txt": {
        "abstract": "{Jeans showed analytically that, in an infinite uniform-density isothermal gas, plane-wave perturbations collapse to dense sheets if their wavelength, $\\lambda$, satisfies $\\lambda > \\lambda_{_{\\rm JEANS}} = \\left(\\pi a^2 / G \\rho_{_0}\\right)^{1/2}$ (where $a$ is the isothermal sound speed and $\\rho_{_0}$ is the unperturbed density); in contrast, perturbations with smaller $\\lambda$ oscillate about the uniform density state. Here we show that Smoothed Particle Hydrodynamics reproduces these results well, even when the diameters of the SPH particles are twice the wavelength of the perturbation. Our simulations are performed in 3-D with initially settled (i.e. non-crystalline) distributions of particles. Therefore there exists the seed noise for artificial fragmentation, but it does not occur.  We conclude that, although there may be -- as with any numerical scheme -- `skeletons in the SPH cupboard', a propensity to fragment artificially is evidently not one of them. ",
        "introduction": "\\label{SEC:INTRO}  Stars form through the collapse and fragmentation of molecular clouds. This is a highly chaotic and non-linear process; it involves many different physical effects, of which the dominant one is arguably gravitational fragmentation; and it involves a very large dynamic range of physical scales and complex geometries. Because the process is chaotic and non-linear, numerical simulations have a central role to play in understanding the interplay between the different physical effects. Because gravitational fragmentation is a critical effect, it is essential that numerical schemes are able to capture this effect properly, i.e. that true gravitational fragmentation is not suppressed by inadequate resolution, and that artificial fragmentation does not occur. Because the process involves a large dynamic range of physical scales and complex geometries, Smoothed Particle Hydrodynamics has been used extensively to simulate star formation (e.g. in 2004 alone, Bonnell, Vine \\& Bate, 2004; Clark \\& Bonnell, 2004; Delgado-Donate, Clarke \\& Bate, 2004a; Delgado-Donate et al., 2004b; Goodwin, Whitworth \\& Ward-Thompson, 2004a,b,c; Hennebelle et al., 2004; Hosking \\& Whitworth, 2004a,b; Jappsen \\& Klessen, 2004; Li, Mac Low \\& Klessen, 2004; Rice et al., 2004; Kurosawa et al., 2004; Price \\& Monaghan, 2004a,b; Schmeja \\& Klessen, 2004; Whitehouse \\& Bate, 2004).  In this paper we present a new demonstration of the ability of Smoothed Particle Hydrodynamics to simulate gravitational fragmentation properly, even at very poor resolution, using the plane-wave analysis initially performed by Jeans (1929). In Section \\ref{SEC:SPH}, we briefly describe SPH in general, and the implementation we use in particular. In Section \\ref{SEC:RESO}, we define -- from first principles -- the resolution required for simulating gravitational fragmentation, i.e. the so-called Jeans Condition. In Section \\ref{SEC:TEST}, we describe the Jeans Test, and in Section \\ref{SEC:ICS}, we explain how the initial conditions for the test are set up. In Section \\ref{SEC:RESU}, we present the results of the test, emphasising how poor the resolution can be made before the results are significantly corrupted. In Section \\ref{SEC:SGCO}, we derive and demonstrate a correction term for use when the Jeans condition is not satisfied or only weakly so. In Section \\ref{SEC:CONC}, we summarise our main conclusions. ",
        "conclusions": "\\label{SEC:CONC}  The conclusions are very simple. SPH using the standard M4 kernel and kernel-softened gravity (i.e. the standard options) only captures fragmentation which is (a) genuine, and (b) resolved. It does not suffer from artificial fragmentation."
    },
    "0512/astro-ph0512071_arXiv.txt": {
        "abstract": "Adopting the framework of the halo occupation distribution (HOD), we investigate the ability of galaxy clustering measurements to simultaneously constrain cosmological parameters and galaxy bias. Starting with a fiducial cosmological model and galaxy HOD, we calculate spatial clustering observables on a range of length and mass scales, dynamical clustering observables that depend on galaxy peculiar velocities, and the galaxy-matter cross-correlation measurable by weak lensing. We then change one or more cosmological parameters and use $\\chi^2$-minimization to find the galaxy HOD that best reproduces the original clustering. Our parameterization of the HOD incorporates a flexible relation between galaxy occupation numbers and halo mass and allows spatial and velocity bias of galaxies within dark matter halos. Despite this flexibility, we find that changes to the HOD cannot mask substantial changes to the matter density $\\Omega_m$, the matter clustering amplitude $\\sigma_8$, or the shape parameter $\\Gamma$ of the linear matter power spectrum --- cosmology and bias are not degenerate. With the conservative assumption of 10\\% fractional errors, the set of observables considered here can provide $\\sim 10\\%$ ($1\\sigma$) constraints on $\\sigma_8$, $\\Omega_m$, and $\\Gamma$, using galaxy clustering data {\\it alone}. The combination $\\sigma_8\\Omega_m^{0.75}$ is constrained to $\\sim 5\\%$. In combination with traditional methods that focus on large-scale structure in the ``perturbative'' regime, HOD modeling can greatly amplify the cosmological power of galaxy redshift surveys by taking advantage of high-precision clustering measurements at small and intermediate scales (from sub-Mpc to $\\sim 20\\hmpc$). At the same time, the inferred constraints on the galaxy HOD provide valuable tests of galaxy formation theory. ",
        "introduction": "From the 1970s through the early 1990s, studies of galaxy clustering drove much of the progress in cosmology.  Measurements of steadily improving dynamic range and precision demonstrated good agreement with the predictions of a cosmological model incorporating scale-invariant, Gaussian primeval fluctuations modulated by the transfer function expected in a universe dominated by cold dark matter (CDM) with $\\Omega_m h \\sim 0.2$ (where $\\Omega_m$ is the matter density parameter and $h \\equiv H_0/100\\;\\hubunits$). The advent of multi-fiber galaxy redshift surveys and improved photometric input catalogs has dramatically improved the precision of clustering measurements over the last decade, beginning with the Las Campanas Redshift Survey (LCRS; \\citealt{Shectman96}) and continuing with the Two-Degree Field Galaxy Redshift Survey (2dFGRS; \\citealt{Colless01}) and the Sloan Digital Sky Survey (SDSS; \\citealt{York00,Abazajian04}).  In parallel, numerical simulations and numerically tested analytic approximations have turned the task of calculating non-linear dark matter clustering from specified initial conditions into an essentially solved problem.  The principal obstacle to drawing cosmological inferences from galaxy clustering measurements is now the uncertainty in the relation between the distribution of observable galaxies and the underlying distribution of dark matter, the problem known as galaxy bias.  Much of the cosmological progress in the last decade has been driven by observations that circumvent this complication, such as cosmic microwave background (CMB) anisotropies, weak gravitational lensing, the Ly$\\alpha$ forest, and the Type Ia supernova diagram.  These new observables favor an inflationary, low-$\\Omega_m$, CDM-dominated model similar to that originally suggested by galaxy clustering, which in turn implies that the galaxies that dominate typical optically selected galaxy surveys must be approximately unbiased, in the sense that the rms galaxy count fluctuations are similar to the rms dark matter density fluctuations on large scales (see, e.g., \\citealt{Lahav02}). However, observed galaxy clustering varies systematically with galaxy luminosity, color, and spectral or morphological type (\\citealt{Norberg02,Zehavi05}, and numerous references therein), and reproducing the observed galaxy correlation function in an inflationary CDM model requires that the bias of the correlation function vary with separation on scales below a few megaparsecs (\\citealt{Jenkins98,Zehavi04}).  The advances in observational cosmology over the last few years have also raised the stakes for galaxy clustering studies.  We are no longer interested in, for example, distinguishing $\\Omega_m \\sim 0.3$ from $\\Omega_m=1$; instead, we want to constrain $\\Omega_m$ at the few percent level to increase the power of tests for the nature of dark energy.  Despite improvements in semi-analytic and numerical modeling of galaxy formation, it is not clear that these methods predict galaxy bias robustly enough for this kind of precision cosmology.  Faced with these challenges, most cosmological applications of the 2dFGRS and the SDSS have focused on the linear or near-linear regime, where generic arguments suggest that the effects of galaxy bias should be relatively simple. These ``perturbative'' analyses of large-scale structure play a significant role in the current web of cosmological constraints (e.g., \\citealt{Percival02,Spergel03,Tegmark04b,Cole05,Seljak05a,Tegmark06}), but they are restricted to large scales where even these enormous surveys have limited statistical precision.  In this paper we argue that recent developments in the theoretical description of galaxy bias allow a more aggressive approach to inferring cosmological constraints from galaxy clustering measurements. We work in the framework of the halo occupation distribution (HOD), which characterizes galaxy bias in terms of the probability distribution $P(N|M)$ that a dark matter halo of virial mass $M$ contains $N$ galaxies of a specified type, together with prescriptions for the spatial and velocity bias of galaxies within dark matter halos \\citep{Ma00,Peacock00,Seljak00,Scoccimarro01,Berlind02,Cooray02}. Here the term ``halo'' refers to a bound dark matter structure of typical overdensity $\\rho/\\bar{\\rho} \\sim 200$, in approximate dynamical equilibrium, which may be the individual halo of a single bright galaxy or the common halo of a galaxy group or cluster.\\footnote{We have in mind the kinds of structures identified in $N$-body simulations by a friends-of-friends algorithm with linking length $l \\sim 0.15-0.2\\bar{n}^{-1/3}$, but the precise definition of halo does not matter provided that one is consistent throughout all calculations.} The flow diagram in Figure~\\ref{fig:flowchart}, adapted from \\cite{Weinberg02}, sketches the interplay between the ``cosmological model'' and the ``physics of galaxy formation'' in determining observable galaxy clustering, which we take to include both the traditional statistics measured from redshift surveys and the galaxy-matter correlations measured by galaxy-galaxy lensing (e.g., \\citealt{Fischer00,Hoekstra01,Sheldon04}). On the left side, the cosmological model, which specifies the initial conditions and the energy and matter contents of the universe, determines the mass function, spatial correlations, and velocity correlations of the dark halo population. The intervening box indicates that the only features of the cosmological model that really matter in this context are $\\Omega_m$ and the amplitude and shape of the linear matter power spectrum $P(k)$, here represented by $\\sigma_8$ (the rms linear matter fluctuation in $8\\hmpc$ spheres), the inflationary spectral index $n_s$, and the transfer function shape parameter $\\Gamma$, which itself depends on the values of $\\Omega_m$, $h$, and the baryon density \\citep{Bardeen86,Hu96}. Other features of the cosmological model, such as the energy density and equation of state of the vacuum component, may have an important impact on other cosmological observables or on the {\\it history} of matter clustering, but they have virtually no effect on the halo population at $z=0$, if the shape of $P(k)$ and the present day value of $\\sigma_8$ are held fixed (see \\citealt{Zheng02}).  \\begin{figure} \\epsscale{1.2} \\plotone{f1.ps} \\epsscale{1.0} \\caption[]{\\label{fig:flowchart} Interplay between the cosmological model and galaxy formation physics in determining observable galaxy clustering.  The cosmological model determines the mass function and clustering of the dark halo population. Galaxy formation physics, operating within this cosmological model, determines the HOD of different galaxy classes. The clustering of any given class of galaxies can be predicted from the halo population and the HOD.  In this paper we investigate how well one can use observations of clustering and galaxy-mass correlations to infer cosmological and HOD parameters simultaneously. } \\end{figure}  On the right side of Figure~\\ref{fig:flowchart}, the box titled ``galaxy formation physics'' represents the additional processes --- such as shock heating, radiative cooling, star formation, feedback, and mergers --- that are essential to producing galaxies and determining their masses, luminosities, diameters, colors, and morphologies. These physical processes operate in the background provided by the evolving halo population, so together with the cosmological model they determine the HODs of different classes of galaxies. The halo population and the HOD together determine galaxy clustering and galaxy-mass correlations.  One can, of course, use hydrodynamic simulations or semi-analytic models to predict galaxy clustering statistics directly, without computing the HOD as an intermediate step (e.g., \\citealt{Kauffmann99,Benson00,Cen00,Pearce01,Weinberg04}). However, these predictions reflect the combination of the cosmological model and the galaxy formation theory, and unless one has complete confidence in the latter, one cannot draw secure conclusions about the former. We would therefore like to know how well cosmological parameters can be constrained {\\it without} relying on a detailed theory of galaxy formation, by using the data themselves to determine the relation between galaxies and dark matter.  For this purpose, the HOD formulation of bias has two key strengths, both emphasized by \\cite{Berlind02}.  The first is the division of labor implied by Figure~\\ref{fig:flowchart}: galaxy formation physics influences the HOD, but the properties of the halo population defined at $\\rho/\\bar{\\rho} \\sim 200$ are determined almost entirely by the much simpler physics of gravitational clustering, which can be modeled accurately using $N$-body simulations or numerically tested analytic approximations.  The second is completeness: given a cosmological model and a specified HOD, one can calculate any galaxy clustering statistic on any scale, by populating $N$-body halos or by drawing on a steadily expanding array of analytic techniques.  The HODs of different galaxy classes therefore encode, statistically, all the aspects of galaxy formation physics that are relevant to predictions of galaxy clustering. (We discuss an important caveat to this statement below.) HOD modeling enables one to take full advantage of galaxy clustering data, with no restriction to large scales or particular statistics.  The goal of our empirical approach is to reverse the causal arrows in Figure~\\ref{fig:flowchart}, working backwards from the data to properties of the HOD and the halo population, and from there to conclusions about galaxy formation physics and fundamental aspects of cosmology. \\cite{Berlind02} considered a fixed cosmological model and showed how changes to the HOD affect many of the traditional statistical measures of galaxy clustering, such as correlation functions, the group multiplicity function, and redshift-space distortions.  They argued that the complementary information from different statistics and different spatial scales should permit an accurate empirical reconstruction of the HOD for a specified cosmology, yielding physically informative tests of galaxy formation theories.  \\cite{Zheng02} showed that $\\Omega_m$, $\\sigma_8$, and the power spectrum shape have non-degenerate effects on the {\\it halo} population, and they argued that changes to the galaxy HOD could not fully mask the effects of cosmological parameter changes. Here we complete this theoretical program by showing how galaxy clustering measurements can constrain simultaneous changes to the cosmological model and the galaxy HOD.  Our results add quantitative teeth to the qualitative arguments and speculations in \\cite{Berlind02}, \\cite{Zheng02}, and \\cite{Weinberg02}.  We define our cosmological parameter space by the values of $\\Omega_m$, $\\sigma_8$, and $\\Gamma$.  We concentrate mainly on $\\Omega_m$ and $\\sigma_8$ because $\\Gamma$ (or, more generally, the shape of the linear power spectrum) can be well constrained by combining the large-scale galaxy power spectrum with other observables like CMB anisotropy and the Ly$\\alpha$ forest. We first compute predicted values of a number of galaxy clustering observables, assuming a central model with observationally motivated choices of cosmological parameters and the galaxy HOD.  We then change one or more cosmological parameters, and we test how well this new cosmological model can reproduce the original clustering ``data'' given complete freedom to vary the HOD within a very flexible parameterization. To gain insight, we first explore several interesting axes within the $(\\Omega_m,\\sigma_8,\\Gamma)$ parameter space, varying parameters individually or in physically motivated combinations. We arrive at definite numbers by assuming that each of our 30 observables, some of which represent the same clustering statistic measured at multiple scales, can be measured with 10\\% uncorrelated fractional uncertainty from a survey like the SDSS.  This assumption seems roughly plausible, but the ultimate strength of cosmological conclusions will depend on the precision and dynamic range of the measurements. Forecasting this precision for a survey like the SDSS, including the covariance of errors among different observables, is itself a major theoretical task, which we will not undertake here. For our theoretical investigation, we further restrict ourselves to observables for which we have reasonable analytic approximations, and we suspect that our quantitative conclusions will in the end prove overly pessimistic because we must omit some observables that contain significant additional information.  As discussed by \\cite{Berlind02}, the completeness of the HOD as a description of bias rests on the assumption that the galaxy content of a halo of virial mass $M$ is statistically independent of the halo's larger scale environment. \\cite{Berlind03} show that this assumption accurately describes the galaxy population in Weinberg et al.'s (\\citeyear{Weinberg04}) hydrodynamic cosmological simulation. It has fundamental theoretical roots in the excursion set model of \\cite{Bond91}, which predicts that the statistical features of a halo's assembly history depend only on its present mass; indeed, all semi-analytic galaxy formation models that use statistically generated merger trees make this assumption implicitly.  Since the halo mass function itself varies with environment, with high-mass halos absent in low-density regions, models that tie galaxy populations to halo masses still predict strong correlations between galaxy properties and the large-scale environment, in good agreement with observations (e.g., \\citealt{Benson01,Berlind05,Zehavi05}). However, while \\cite{Lemson99} show that $N$-body halos of mass $M \\ga 10^{13}M_\\odot$ have similar properties and formation histories in different environments, recent studies show a substantial correlation of halo formation redshift with large-scale overdensity for lower mass halos ($M \\la 10^{12.5} M_\\odot$), contradicting the simplest form of the \\cite{Bond91} model (\\citealt{Gao05,Harker06}; see also \\citealt{Sheth04}). As discussed below in \\S\\ref{sec:environment}, we expect the impact of such dependence on the clustering statistics of mass- or luminosity-thresholded galaxy samples to be small, but given the level of precision we eventually hope to reach, it will probably be necessary to allow for it when fitting the data.  Strategies for incorporating environmental variations into HOD modeling will require guidance from the next generation of hydrodynamic and semi-analytic models of galaxy clustering, and we reserve this task for future investigation.  Ultimately, the assumption of an environment-independent HOD, or of any particular parameterization of environmental dependence, must be tested empirically by checking that the adopted model can consistently explain the variations of galaxy clustering with large-scale environment.  Many current efforts to constrain cosmological parameters with galaxy clustering data make (explicitly or implicitly) the much stronger assumption that galaxy bias can be adequately described by a linear model, $\\delta_g=b\\delta_m$, on the relevant scales.  Here $\\delta_g$ and $\\delta_m$ are the galaxy and dark matter density contrast, respectively, and the linear bias factor $b$ may depend on galaxy type but is assumed to be independent of scale. In this case, $P_g(k)=b^2 P_m(k)$, and the shape of the galaxy power spectrum provides powerful cosmological constraints in combination with CMB data even if $b$ is unknown \\citep{Percival02,Tegmark04b,Tegmark06}. Large-scale redshift-space distortions measure $\\beta \\equiv \\Omega_m^{0.6}/b$ \\citep{Kaiser87,Hamilton98,Hawkins03}, which in combination with clustering measurements constrains $\\sigma_8\\Omega_m^{0.6} = \\Omega_m^{0.6}\\sigma_{8,g}/b$, where $\\sigma_{8,g}$ is the measured rms fluctuation of the {\\it galaxy} density contrast at $8\\hmpc$.  The mass function or X-ray temperature function of galaxy clusters constrains a similar combination of $\\sigma_8$ and $\\Omega_m$ \\citep{White93} by a route that is independent of galaxy bias.  At low redshift, the combination of galaxy-galaxy lensing with the galaxy-galaxy correlation function constrains a somewhat different combination of these parameters, $\\Omega_m/b \\propto \\sigma_8\\Omega_m$, assuming a linear bias model with $\\xigg(r) = b\\xigm(r) = b^2\\ximm(r)$. The triangle shape dependence of the reduced galaxy bispectrum can constrain $b$ directly \\citep{Fry94,Verde02}, in this case assuming a quadratic bias model $\\delta_g = b\\delta_m + b_2\\delta_m^2 + {\\rm const.}$ to relate the galaxy and mass density fields at second order \\citep{Fry94,Juszkiewicz95}.  General theoretical arguments suggest that linear bias should be a good approximation for the power spectrum on sufficiently large scales, provided that the efficiency of galaxy formation is determined by the local ($r < $ few Mpc) environment \\citep{Coles93,Fry93,Weinberg95,Mann98,Scherrer98,Narayanan00}. Numerical experiments using such ``local'' bias models also show that the ``$b$'' multiplying the power spectrum amplitude should be similar to the ``$b$'' affecting redshift-space distortions \\citep{Berlind01}. Thus, the combination of these perturbative galaxy clustering analyses can in principle yield separate constraints on $\\Gamma$ (measured directly from the power spectrum shape), $\\Omega_m$, and $\\sigma_8$, with the degeneracy of $\\Omega_m$ and $\\sigma_8$ broken either by the bispectrum analysis or by the different $\\Omega_m$ dependence of the galaxy-galaxy lensing constraint and the $\\beta$ or cluster normalization constraints.  The applications of these techniques to 2dFGRS and SDSS data show an impressive degree of internal consistency and good agreement with external cosmological constraints (\\citealt{Percival02,Verde02,Spergel03,Tegmark04b,Cole05, Seljak05a,Tegmark06}).  However, reliance on the linear or quadratic bias approximations restricts these analyses to large scales, and it is not clear that these approximations hold consistently at the level of accuracy desired for further improvements (better than 10\\%, say), since scatter in the relation between galaxy and mass densities can have different effects on different statistics and at different scales \\citep{Pen98,Dekel99}. The HOD approach to galaxy clustering analysis substitutes a much more general model of galaxy bias, and it makes use of high-precision measurements from small and intermediate scales (from sub-Mpc to $\\sim 20\\hmpc$) in addition to the lower precision measurements in the perturbative regime. Furthermore, regardless of the strength of the cosmological constraints, the HOD parameters derived from the data themselves provide detailed tests of galaxy formation models.  The program of HOD-based interpretation of observed galaxy clustering is, in fact, well underway. \\cite{Jing98a} and \\cite{Jing02} used HOD-type bias models to interpret the correlation functions and pairwise velocity dispersions measured from the LCRS and the Point-Source Catalog Redshift Survey (PSCz; \\citealt{Saunders00}).  \\cite{Peacock00}, \\cite{Marinoni02}, and \\cite{Kochanek03} used the group multiplicity function to constrain the galaxy occupations of high-mass halos. \\cite{Guzik02}, \\cite{Seljak05a}, and \\cite{Mandelbaum06} applied HOD modeling to galaxy-galaxy lensing measurements from the SDSS.  \\cite{Zehavi04}, analyzing a volume-limited sample of bright ($M_r<-21$) SDSS galaxies, showed that HOD models naturally explain the observed deviation from a power-law correlation function. \\cite{Zehavi05} measured the luminosity and color dependence of the SDSS galaxy correlation function and used it to infer the luminosity and color dependence of the galaxy HOD, finding results in good qualitative agreement with theoretical predictions \\citep{Zheng05}. \\cite{Magliocchetti03} used a similar approach to investigate the halo occupations of early- and late-type galaxies in the 2dFGRS, while \\cite{Collister05} derived HODs of red and blue 2dFGRS galaxies from the group catalog of \\cite{Eke04}, testing the consistency of their result by comparing predicted and observed correlation functions. \\cite{Abazajian05} used the SDSS measurements for $M_r<-21$ galaxies in conjunction with CMB anisotropy data to infer simultaneous constraints on HOD and cosmological parameters. \\cite{Tinker05} used HOD modeling of Zehavi et al.'s (\\citeyear{Zehavi05}) clustering measurements to predict cluster mass-to-light ($M/L$) ratios, and they inferred constraints on $\\sigma_8\\Omega_m^{0.6}$ by comparing to published $M/L$ measurements. HOD models have been applied to the interpretation of high-redshift clustering by \\cite{Bullock02}, \\cite{Moustakas02}, \\cite{Yan03}, \\cite{Zheng04}, \\cite{Ouchi05}, \\cite{Lee06}, \\cite{Coil06}, and \\cite{Cooray06}. Finally, \\cite{Bosch03a} have initiated a comprehensive program similar to the one described here, based on the closely related conditional luminosity function (CLF) formalism (see also \\citealt{Bosch03b,Bosch04,Bosch05,Bosch06,Yang03,Yang04,Yang05,Mo04, Wang04}). We discuss the similarities and differences between the HOD and CLF approaches in \\S~\\ref{sec:discussion}.  In the next section we define our class of cosmological models, list the analytic approximations we use for properties of the halo population, and describe our flexible parameterization of the galaxy HOD. In \\S~\\ref{sec:observables} we describe the methods and approximations that we adopt for analytic calculation of galaxy clustering observables. Section~\\ref{sec:results} presents our main results, showing the ability of complementary galaxy clustering measurements to constrain the galaxy HOD and cosmological parameters simultaneously. Section~\\ref{sec:environment} discusses the issue of environmental variations of the HOD. Section~\\ref{sec:discussion} summarizes our findings and discusses the overall prospects for application of this approach. A reader familiar with the field who wants an overview of our main results can read the first paragraph of \\S\\ref{sec:cosmology}, skim \\S~\\ref{sec:hodpar} to understand our HOD parameterization, then skip to \\S~\\ref{sec:results}, paying particular attention to \\S~\\ref{sec:omegam}, \\S~\\ref{sec:general}, and \\S~\\ref{sec:influence} and Figures~\\ref{fig:chi2_omega}, \\ref{fig:chgerr}, and \\ref{fig:matrix}. ",
        "conclusions": "\\label{sec:discussion}  Our analysis shows that HOD modeling can substantially increase the cosmological power of galaxy clustering measurements, by breaking degeneracies between the clustering of dark matter and the bias of galaxies with respect to mass.  Changing the shape or amplitude of the matter power spectrum or the value of $\\Omega_m$ alters the mass function, spatial clustering, and velocity statistics of the dark halo population in well-understood ways \\citep{Zheng02}.  Our experiments here, which verify the qualitative arguments of \\cite{Berlind02} and \\cite{Zheng02}, show that changes to the galaxy HOD cannot mask these changes in the underlying dark halo population. With our highly flexible parameterization of the HOD, the set of observables considered here yields $1\\sigma$ uncertainties of $\\sim 10\\%$ in $\\sigma_8$, $\\Omega_m$, and $\\Gamma$ and $\\sim 5\\%$ uncertainty in the combination $\\sigma_8 \\Omega_m^q$ with $q \\sim 0.75$. We expect these forecasts to be conservative, as we have not included observables for which we did not have ready analytic approximations, and our assumption of 10\\% measurement errors is pessimistic in at least some cases.  The physical origin of these cosmological constraints is straightforward to understand for simple changes in $\\Omega_m$, $\\sigma_8$, or $\\sigma_8 \\Omega_m^{0.5}$, as discussed in \\S\\S\\ref{sec:omegam}-\\ref{sec:clnorm}. The general theme of these discussions is that, for a given cosmological model, the spatial clustering of galaxies largely determines the number of galaxies in halos of a given spatial abundance. Dynamically sensitive statistics then reveal the halo mass scale, which depends on $\\sigma_8$ and $\\Omega_m$.  We allow an arbitrary bias $\\alpha_v$ between the galaxy and dark matter velocity dispersions within halos, but this freedom does not eliminate the constraining power of dynamical observables because the space velocities of the halos themselves do not change.  The main parameter degeneracy is approximately $\\sigma_8 \\Omega_m^{0.75}$ because fixing this combination roughly fixes the halo velocity scale and the abundance of halos at the mass scale of rich galaxy groups. However, the changing shape of the mass function, the differing sensitivities of different velocity measures, and the different $(\\sigma_8,\\Omega_m)$ dependence of galaxy-galaxy lensing all serve to break this degeneracy.  Figures~\\ref{fig:chgerr} and~\\ref{fig:matrix} demonstrate that the cosmological constraints emerge from the full web of clustering observables and are not dominated by one or two statistics on their own.  Constraints on the galaxy HOD will themselves provide valuable tests of galaxy formation models.  The cutoff regime of $\\Navg$ is difficult to pin down with the clustering statistics considered here, but for a known cosmological model the relation between average satellite number and halo mass is well determined, and the satellite distribution width $\\omega$, concentration $c_g$, and velocity bias $\\alpha_v$ are measured to $\\sim 10\\%$, 30\\%, and 3\\%, respectively (Fig.~\\ref{fig:mcmc_hodpar}). All of these quantities depend in detail on the physics that governs the evolution of satellites in larger halos (see, e.g., \\citealt{Taylor04,Taylor05a,Taylor05b,Zentner05}), while the relative mass of halos that host central and satellite galaxies depends on the efficiency with which halos feed baryonic mass to their central objects (see, e.g, discussions by \\citealt{Berlind03,Zheng05}). These galaxy formation constraints will be especially powerful when derived as a function of luminosity, stellar mass, or other observables.  The two main assumptions built into our modeling are the central-satellite parameterization and environment independence of the HOD.  The central-satellite distinction appears well rooted in galaxy formation physics, and it allows us to represent the range of plausible galaxy HODs more completely with a moderate number of parameters.  However, we have confirmed with other tests that if we model galaxy bias with a flexible HOD parameterization that does not impose a central-satellite distinction, but instead introduces a characteristic mass for the narrow-to-wide transition of $\\PNM$ as in \\cite{Scranton03}, then we reach almost identical conclusions about the cosmological constraining power of the clustering observables considered in this paper.  Recent numerical results imply an environmental dependence of halo formation times that opens the door to environmental variation of the HOD, especially in the single-galaxy regime.  As discussed in \\S\\ref{sec:environment}, we expect the quantitative impact of such dependence to be small, but potentially significant at the high level of precision we ultimately hope to attain.  Investigation of environmental dependence effects and methods of allowing for them in HOD modeling are a high priority for future work.  The cosmological modeling approach advocated here is closely related to the CLF method introduced by \\cite{Yang03} and \\cite{Bosch03a}, who use clustering data and the global galaxy luminosity function to constrain the dependence of the luminosity function on halo mass. In principle, the CLF and HOD methods are equivalent --- they are merely differential and integral forms of one another. One can derive the CLF from a series of HOD fits to galaxy samples with different luminosity thresholds \\citep{Zehavi05,Tinker05}. Conversely, one can integrate the CLF to infer $\\Navg$ for galaxies above a luminosity threshold \\citep{Yang03,Bosch03b,Bosch05}. The principal virtue of our HOD-based approach is that by focusing on a single, well-defined class of galaxies, we can parameterize the HOD in a way that seems likely to capture the predictions of any reasonable galaxy formation model. This kind of comprehensive parameterization is more difficult to achieve for the full CLF, and most analyses to date have assumed, for example, that the CLF has a Schechter form in halos of fixed mass.  Nonetheless, it is valuable to pursue both HOD and CLF approaches and test for consistency of conclusions.  HOD modeling complements rather than replaces the perturbative approach based on large-scale measures that can be modeled with linear or quadratic bias.  HOD modeling is more complex, but it can take advantage of high-precision clustering measurements on small and intermediate scales.  HOD modeling can also amplify the power of the perturbative approach, extending its reach further into the non-linear regime and checking its range of validity at a desired level of precision.  For example, \\cite{Tinker06} show that an HOD-based approach to redshift-space distortions can improve recovery of the perturbative parameter $\\beta$ that controls large-scale flows.  \\cite{Yoo06} show that the linear bias model provides an accurate description of galaxy-galaxy lensing for $r \\geq 2\\hmpc$, and they show how to accurately model this phenomenon on smaller scales (see also \\citealt{Guzik02,Mandelbaum05}). J. Yoo et al. (2007, in preparation) show that the scale-dependent bias factors derived by fitting the projected galaxy correlation function can extend recovery of the shape of the linear matter power spectrum into the mildly non-linear regime.  CLF and HOD analyses of the 2dFGRS and SDSS redshift surveys have already produced a number of interesting results, even though they have considered only a fraction of the potential clustering observables.  These results confirm, at least qualitatively, many of the basic predictions of current galaxy formation models, including the general form of the mean occupation function, the dependence of this function on luminosity, the existence of a minimum mass-to-light ratio in the halos of $\\sim L_*$ galaxies where galaxy formation is most efficient, the large gap between the minimum halo mass for central and satellite galaxies above a luminosity threshold, the sub-Poisson fluctuations of $\\PNM$ that are a consequence of this gap, and the strong preference of galaxies with older stellar populations for higher mass halos \\citep{Bosch03a,Bosch03b,Magliocchetti03,Tinker05,Yang05,Zehavi04,Zehavi05, Collister05}. In combination with CMB data, the cosmological constraints from HOD modeling of the SDSS projected galaxy correlation function are almost as tight as those from the large-scale galaxy power spectrum, and the two analyses are consistent within their statistical uncertainties \\citep{Abazajian05}. For the most part, the cosmological inferences from CLF/HOD modeling of galaxy clustering agree with those from other methods, but matching the mass-to-light ratios of galaxy clusters simultaneously with other clustering data appears to require values of $\\Omega_m$ and/or $\\sigma_8$ that are substantially lower than the commonly adopted values of 0.3 and 0.9 \\citep{Bosch03a,Tinker05}. If this conclusion is correct, then the evidence for it should become much stronger as more clustering observables are brought into play and the SDSS data set itself moves to completion.\\footnote{Since the original submission of this paper, the analysis of the three-year {\\it WMAP} data set \\citep{Spergel06} has provided strong support for this shift in cosmological parameter values, in excellent agreement with the results of the HOD and CLF modeling and the related ``empirical model'' method of \\citet{Vale06}.}  We have focused in this paper on the cosmological parameter constraints that can be derived from galaxy clustering alone, using external data only to guide the choice of power spectrum shape and motivate the assumption of Gaussian initial conditions.  As with perturbative analyses of large-scale structure, the long-term interest lies in combining these constraints with those from the CMB, the Ly$\\alpha$ forest, Type Ia supernovae, and other cosmological observables.  The complementary sensitivities of these observables lead to much tighter parameter constraints.  More importantly, conflicts among them could point the way to physics beyond the simplest versions of $\\Lambda$CDM, such as evolving dark energy, a gravitational wave contribution to CMB anisotropy, departures from scale invariance in the primordial power spectrum, non-zero space curvature, cosmologically significant neutrino masses, and so forth. By sharpening the constraints from large-scale structure in the new generation of galaxy redshift surveys, HOD modeling can play a critical role in efforts to test the standard cosmological model and, perhaps, discover its breaking points."
    },
    "0512/astro-ph0512592_arXiv.txt": {
        "abstract": "{}{ Nonlinear behaviour of galactic dynamos is studied, allowing for magnetic helicity removal by the galactic fountain flow.}{ A suitable advection speed is estimated, and a one-dimensional mean-field dynamo model with dynamic $\\alpha$-effect is explored. }{ It is shown that the galactic fountain flow is efficient in removing magnetic helicity from galactic discs. This alleviates the constraint on the galactic mean-field dynamo resulting from magnetic helicity conservation and thereby allows the mean magnetic field to saturate at a strength comparable to equipartition with the turbulent kinetic energy. }{} ",
        "introduction": "The major controversy in mean-field dynamo theory is related to its nonlinear form relevant when the initial exponential growth of the large-scale magnetic field saturates and the field reaches statistical equilibrium. The core of the problem is the effect  of the small-scale (turbulent) magnetic field on the evolution of the large-scale (mean) magnetic field. It has been argued (Vainshtein \\& Cattaneo \\cite{VC92}, Cattaneo \\& Hughes \\cite{CH96}) that the Lorentz force due to the rapidly growing small-scale magnetic field $\\vec b$ can make the large-scale dynamo action inefficient so that it produces only a negligible mean field $\\meanBB$, with $\\meanBB^2\\simeq \\Rm^{-1}\\overline{\\bb^2}$, where $\\Rm\\ (\\gg1)$ is the magnetic Reynolds number. Here and below, overbars denote averaged quantities. In particular, the $\\alpha$-effect (a key ingredient of the mean-field dynamo) is catastrophically quenched well before the large-scale magnetic field can be amplified to the strength observed in astrophysical objects, $\\meanB\\simeq b$.  The suppression of the $\\alpha$-effect can be a consequence of the conservation of magnetic helicity in a medium of high electric conductivity (see Brandenburg \\& Subramanian \\cite{BS05a} for a review). In a closed system, magnetic helicity can only evolve on the Ohmic time scale which is proportional to $\\Rm$; in galaxies, this time scale by far exceeds the Hubble time. Since the large-scale magnetic field necessarily has non-zero helicity in each hemisphere through the mutual linkage of poloidal and toroidal fields, the dynamo also has to produce small-scale helical magnetic fields with the opposite sign of magnetic helicity. Unless the small-scale magnetic field can be transported out of the system, it quenches the $\\alpha$-effect together with the mean-field dynamo.  Blackman \\& Field (\\cite{BF00}) first suggested that the losses of the small-scale magnetic helicity through the boundaries of the dynamo region can be essential for mean-field dynamo action. Such a helicity flux can result from the anisotropy of the turbulence combined with large-scale velocity shear (Vishniac \\& Cho \\cite{VC01}, Subramanian \\& Brandenburg \\cite{SB04}) or the non-uniformity of the $\\alpha$-effect (Kleeorin et al.\\ \\cite{KMRS00}). The effect of the Vishniac-Cho flux has already been confirmed in simulations, allowing the production of significant fields, $\\meanB\\simeq b$ (Brandenburg \\cite{B05}).  Here we suggest another simple mechanism where the advection of small-scale magnetic fields (together with the associated magnetic helicity) away from the dynamo region allows healthy mean-field dynamo action. This effect naturally arises in spiral galaxies where magnetic field is generated in the multi-phase interstellar medium. The mean magnetic field is apparently produced by the motions of the warm gas (\\S4.3 in Beck et al.\\ \\cite{BBMSS96}), which is in a state of (statistical) hydrostatic equilibrium with a scale height of $h\\simeq0.5\\kpc$ (e.g.\\ Korpi et al.\\ \\cite{K99}).  However, some of the gas is heated by supernova explosions producing a hot phase whose isothermal scale height is 3\\,kpc. The hot gas leaves the galactic disc, dragging along the small-scale part of the interstellar magnetic field. Thus, the disc-halo connection in spiral galaxies represents a mechanism of transport of small-scale magnetic fields and small-scale magnetic helicity from the dynamo active disc to the galactic halo. As we show here, this helps to alleviate the catastrophic $\\alpha$-quenching under realistic parameters of the interstellar medium. ",
        "conclusions": "The vertical advection of magnetic helicity by galactic fountain flow resolves straightforwardly the controversy of nonlinear mean-field galactic dynamos. For $C_U=0$, the mean magnetic field does initially reach a level consistent with Eq.~(\\ref{steady}), but then rapidly decays to negligible values (see Brandenburg \\& Subramanian \\cite{BS05c}). The essential role of the advection is to provide the system with an opportunity to reach and then maintain the steady state (\\ref{steady}) within the galactic lifetime.  These conclusions follow from Fig.~\\ref{axsinz}, where moderate advection drastically changes the mean field levels achievable at $t\\la10^{10}\\yr$ and prevents catastrophic quenching of the dynamo. Excessive advection, however, hinders the dynamo as it removes the mean field from the dynamo active region.  The steady-state strength of the mean magnetic field obtained in our model is of order $(0.1$--$0.2)\\,B_\\mathrm{eq}\\simeq0.5$--$1\\mkG$, which is a factor of several weaker than what is observed; we made no attempt here to refine the model. What is important, the mechanism suggested here resolves the problem of catastrophic quenching of the dynamo.  The applicability of the vacuum boundary conditions (\\ref{bconds}) to a system with gas outflow from the disc can be questionable. Analysis of dynamo models with outflow (Bardou et al.\\ \\cite{BR01}) suggests that reasonable changes to the boundary conditions do not affect the dynamo too strongly, but this aspect of our model should be further explored.  We have neglected the intrinsic difference of the behaviours of the mean and turbulent magnetic fields near the disc surface. Since the horizontal size of the hot cavities is larger than the scale of the turbulent magnetic field but smaller than the scale of the mean magnetic field, the Lorentz force can resist the advection of the mean field more efficiently than that of the turbulent field. Furthermore, large-scale magnetic field loops drawn out by the fountain flow can be detached from the parent magnetic lines by reconnection, so that the flow will carry mostly small-scale fields. Therefore, a more plausible (albeit less conservative) model would include advection of the small-scale (but not the large-scale) magnetic field. In such a model, the effect discussed here can be even better pronounced.  Brandenburg et al.\\ (\\cite{BMS95}) argue that the fountain flow can transport the large-scale magnetic field into the halo by topological pumping if the hot gas forms a percolating cluster in the disc and if the turbulent magnetic Reynolds number $C_U$ in the fountain flow exceeds $\\simeq20$; our estimate given below is $C_U\\la 1$, and we expect that the small-scale magnetic fields are removed from the disc more efficiently than the mean field.  The idea that advection of small-scale magnetic fields can help the galactic dynamo may be more robust than our particular model of dynamo quenching that involves the magnetic $\\alpha$-effect. For example, if the dynamo coefficients are quenched due to the suppression of Lagrangian chaos by the small-scale magnetic fields (Kim \\cite{Kim99}), their advection out of the galaxy will still allow the dynamo to operate efficiently."
    },
    "0512/astro-ph0512521_arXiv.txt": {
        "abstract": "The Indo-US coud\\'{e} feed stellar spectral library (CFLIB) made available to the astronomical community recently by Valdes et al. (2004) contains spectra of 1273 stars in the spectral region 3460 to 9464 \\AA~ at a high resolution of 1 \\AA~ FWHM and a wide range of spectral types. Cross-checking the reliability of this database is an important and desirable exercise since a number of stars in this database have no known spectral types and a considerable fraction of stars has not so complete coverage in the full wavelength region of 3460-9464 \\AA~ resulting in gaps ranging from a few \\AA~ to several tens of \\AA~. In this paper, we use an automated classification scheme based on Artificial Neural Networks (ANN) to classify all 1273 stars in the database. In addition, principal component analysis (PCA) is carried out to reduce the dimensionality of the data set before the spectra are classified by the ANN. Most importantly, we have successfully demonstrated employment of a variation of the PCA technique to restore the missing data in a sample of 300 stars out of the CFLIB. ",
        "introduction": "Automated schemes of data validation and analysis have assumed added significance recently as larger databases are increasingly becoming available in almost all areas of observational astronomy. With the advent of bigger CCD detectors in spectroscopy, the need for having large libraries of stellar spectra at high spectral resolution is also getting fulfilled. Jacoby, Hunter, \\& Christian (1984, hereafter JHC) made 158 spectra available in the range of 3510-7427\\AA~ at 4.5\\AA~ resolution. Prugniel \\& Soubiran (2001) published a library of 708 stars using the ELODIE eschelle spectrograph at the Observatoire de Haute-Province that covers a wavelength band of 4100-6800 \\AA~ at a resolution of R=42,000. Cenarro et al. (2001) have provided a database of 706 stellar spectra in the wavelength region 8350-9020 \\AA~ at 1.5 \\AA~ resolution and Le Borgne et al. (2003, hereafter STELIB) with 247 spectra in the range of 3200-9500\\AA~ at 3\\AA~ resolution. Moultaka et.al. (2004, hereafter ELODIE) compiled 1959 spectra in the range of 4000-6800\\AA~ at a resolution of 0.55\\AA~.  More recently, Valdes et al. (2004) observed more than 1200 stars with emphasis on broad wavelength coverage (3400-9500 \\AA) at a resolution of $\\sim 1$ \\AA~ FWHM at an original dispersion of 0.44 \\AA~ per pixel. Their coud\\'{e} feed stellar spectral library (CFLIB) provides a resolution sufficient to resolve numerous diagnostic spectral features that can be used in the automated parameterization of spectra.    Neural networks are a form of multiprocessor computing system, with simple processing elements with a high degree of inter-connection, simple scalar messaging and adaptive interaction between elements. In a supervised back propagation algorithm, the network topology is constrained to be feed-forward, i.e., connections are generally allowed from the input layer to the first (and mostly only) hidden layer; from the first hidden layer to the second,..., and from the last hidden layer to the output layer. The hidden layer learns to recode (or to provide a representation for) the inputs. More than one hidden layer can be used. The architecture is more powerful than single-layer networks: it can be shown that any mapping can be learned, given two hidden layers (of units).  Automated schemes like the Artificial Neural Networks (ANN) have been used in Astronomy for a number of data analysis tasks like scheduling observations (Johnson and Adorf, 1992), adaptive optics (Angel, Wizinowich and Lloyd-Hart, 1990), stellar spectral classification (Gulati et al., 1994; von Hippel et al. 1994) and star-galaxy separation studies (Odewahn et al., 1992). In addition Gulati, Gupta and Rao (1997) extended the ANN analysis to compare synthetic and observed spectra of G and K dwarfs. Gulati, Gupta and Singh (1997) estimated interstellar extinction E(B-V) using ANN from low-dispersion ultraviolet spectra for O and B stars. Bailer-Jones, Gupta and Singh (2002) provide a review of the ANN applications in astronomical spectroscopy.  Another powerful statistical tool for data analysis is the Principal component analysis (PCA). It involves a mathematical procedure that transforms a number of (possibly) correlated variables into a (smaller) number of uncorrelated variables called principal components. The first principal component accounts for as much of the variability in the data as possible, and each succeeding component accounts for as much of the remaining variability as possible. Objectives of principal component analysis are to discover or to reduce the dimensionality of a data set and to identify new meaningful underlying variables. The technique has been used widely for a number of applications in Astronomy, viz., for stellar classification by Murtagh and Heck (1987), Storrie-Lombardi et al. (1994) and Singh et al. (1998), by Francis et al. (1992) for QSO spectra, and for galaxy spectra by Sodr\\'{e} and Cuevas (1994), Connolly et al. (1995), Lahav et al. (1996) and Folkes, Lahav and Maddox (1996).  Another important application was developed by Unno and Yuasa (1992, 2000) for supplementing missing observational data using a generalized PCA technique. Subsequently, Yuasa, Unno and Magono (1999) made use of this technique to determine distances of 183 mass-losing red giants.   The primary aim of this paper is to perform validity checks on the CFLIB by running the ANN code on various inter library sets like JHC, ELODIE and STELIB. In the next section we describe the ANN analysis. In Section 3, we demonstrate the possibility of using PCA for filling the gaps in the spectra of CFLIB. In Section 4, we summarize important conclusions of the study.  \\begin{table} \\caption{ANN training and test cases. Two hidden layers were used for all the cases} \\begin{center} \\begin{tabular}{cccccc} \\hline \\hline Case & Training  & Testing & $\\lambda$ Region & Resolution & Class. error\\\\ & (Library, & (Library, & (used) & &\\\\ & No. of Spectra)& No. of Spectra) & & &\\\\ \\hline \\hline A1& JHC, 158 & CFLIB, 1273& 4100-5500\\AA~ & 4.5\\AA~ & 674.3 \\\\ A2& JHC, 158 & ELODIE, 1959& 4100-5500\\AA~ & 4.5\\AA~ & 514.8 \\\\ A3& JHC, 158 & STELIB, 247& 3600-7400\\AA~ & 4.5\\AA~ & 861.2 \\\\ B1& ELODIE, 174 & ELODIE, 1959 & 4000-5500\\AA~ & 1\\AA~ & 496.4 \\\\ B2& ELODIE, 174 & ELODIE, 1959 & 4000-6800\\AA~ & 1\\AA~ & 576.0 \\\\ B3& ELODIE, 174 & CFLIB, 1272 & 4000-5500\\AA~ & 1\\AA~ & 742.7 \\\\ B4& ELODIE, 174 & CFLIB, 1273 & 4000-6800\\AA~ & 1\\AA~ & 848.5 \\\\ C1& STELIB, 247 & CFLIB, 1273 & 4000-6800\\AA~ & 3\\AA~ & 643.9\\\\ C2& STELIB, 247 & CFLIB, 1273 & 3500-9400\\AA~ & 3\\AA~  & 670.6\\\\ C3& STELIB, 247 & ELODIE, 1959 & 4000-6800\\AA~ & 3\\AA~ & 501.8\\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{figure} \\begin{center} \\FigureFile(120mm,120mm){fig1.eps} \\end{center} \\caption{Classification scatter plots for Cases A1-A3} \\end{figure} ",
        "conclusions": "We have performed an extensive analysis based on artificial neural networks to classify stars in an automated manner in the Indo-US CFLIB using three databases viz., JHC, STELIB and ELODIE. The main aim of this exercise was two-fold. One was to perform the reliability checks on CFLIB to see how the gaps in the library affect the classification accuracy. We find that the despite the presence of gaps, we have achieved classification accuracy of less than one main class. The second aim was to evolve and test automated procedures of classifying stellar spectra. This was achieved by trying our ANN scheme on ten different cases of training and testing on different pairs of libraries. The schemes are numerically intensive, with the ANN training stage requiring several hours of CPU time on the fastest of workstations. A PCA analysis was employed successfully to reduce the dimensionality of the data set and hence faster training without appreciable loss in classification accuracy.  Both STELIB and ELODIE libraries have variable spans of wavelength gaps where the fluxes are filled with zeros (similar to CFLIB). Such gaps lead to classification errors in the PCA and ANN schemes. However, we have carried out some preliminary analysis with the basic $\\chi^2$ minimization scheme on these libraries wherein the gap portions were omitted in both train and test sets. This lead to a remarkable improvements in the classification accuracy.  We have also employed a generalized principal component analysis to first create and then fill the gaps in a sample of 300 stars out of the CFLIB in the blue region. At present, we have used a simplified system to reconstruct flux values at one wavelength bin at a time for these 300 stars. We hope to exploit the full potential of the scheme and attempt to fill larger gaps in stellar spectra in a subsequent study.  MY and HPS are grateful to JSPS (Japan Society for Promotion of Science) and DST (Department of Science \\& Technology, India) for financial support for exchange visits which made this work possible. MY would like to thank Emeritus Professor W. Unno of the University of Tokyo for helpful discussions. The research has made use of the SIMBAD database, operated by CDS, Strasbourg, France and the INDO-US CFLIB managed by NOAO, Tucson, AZ, USA.     \\appendix"
    },
    "0512/astro-ph0512651_arXiv.txt": {
        "abstract": "It is the common consensus that the expansion of a universe always slows down if the gravity provided by the energy sources therein is attractive and accordingly one needs to invoke dark energy as a source of anti-gravity for understanding the cosmic acceleration. To examine this point we find counter-examples for a spherically symmetric dust fluid described by the Lemaitre-Tolman-Bondi solution without singularity. Thus, the validity of this naive consensus is indeed doubtful and the effects of inhomogeneities should be restudied. These counter-intuitive examples open a new perspective on the understanding of the evolution of our universe. ",
        "introduction": "} In 1998, via the distance measurements of type Ia supernovae, it was discovered that the expansion of the universe is accelerating \\cite{Perlmutter:1999np,Riess:1998cb}. The accelerating expansion of the present universe was reinforced recently by the updated supernova data \\cite{Tonry:2003zg,Knop:2003iy,Barris:2003dq,Riess:2004nr,Riess:2006fw} and the WMAP measurement \\cite{WMAP:2008} of cosmic microwave background (CMB). Regarding the cosmic evolution, it is the common consensus that normal matter (such as protons, neutrons, electrons, etc.) can only provide attractive gravity and therefore should always slow down the cosmic expansion, i.e., \\begin{equation} \\textrm{Normal Matter} \\, \\Rightarrow \\, \\textrm{Attractive Gravity} \\, \\Rightarrow \\, \\textrm{Deceleration} \\, . % \\label{consensus} \\end{equation} Thus, to explain this surprising, mysterious phenomenon of the accelerating expansion, many people rely on exotic energy sources, as generally called ``dark energy'', which provide significant negative pressure and accordingly anti-gravity (repulsive gravity).  The above conclusion about the existence of the cosmic acceleration and the necessity of introducing dark energy is based on a simplified cosmological model, the Friedmann-Lemaitre-Robertson-Walker (FLRW) model, and indeed could be model-dependent. In the FLRW model the universe is assumed to be homogeneous and isotropic (i.e.\\ the cosmological principle) and accordingly the Robertson-Walker (RW) metric is invoked in the Einstein equations that describe % the evolution of the universe. Meanwhile the energy-momentum tensor in the right-hand side of the Einstein equations is regarded to truly reflect the real energy distribution (averaged in space) of the universe. Nevertheless, so far there is no convincing proof validating this simplification.  Apparently, our universe at present is not homogeneous at small scales. The cosmological principle is roughly realized only at very large scales. To take advantage of the cosmological principle and invoke the RW metric, it is necessary to perform spatial averaging over large scales, along with which the form of the Einstein equations in general should change because of the non-linearity of the Einstein equations \\cite{Ellis:1984}. That is, (i) when invoking the RW metric and the real energy sources % of our universe in the Einstein equations for describing the cosmic evolution, the left-hand side (the geometry part) of the Einstein equations should be modified, or, from another point of view, (ii) if one insists to use the Einstein tensor corresponding to the RW metric in the left-hand side of the Einstein equations (or, equivalently, moving the above mentioned modification in the geometry part to the right-hand matter-energy part), there should appear new, effective energy sources coming from geometry (which is certainly ``dark'') and consequently the energy-momentum tensor in the right-hand side does not truly correspond to the real energy distribution of our universe. Thus, generally speaking, it is doubtful to employ the Einstein equations to describe long-time, large-scale phenomena, such as the evolution of our universe, while the spatially averaged % energy-momentum tensor and the spatially averaged metric tensor are used therein. (For more discussions about the validity and the problems of the FLRW cosmology, see \\cite{Gu:2006}.)   Instead of invoking exotic energy sources or unconfirmed physics, it is suggested \\cite{Buchert:1999mc,Schwarz:2002ba,Bene:2003fz,Rasanen:2003fy,Rasanen:2004js,Kolb:2004am,Barausse:2005nf,Kolb:2005me,Notari:2005xk,Rasanen:2005zy,Kolb:2005da,Mansouri:2005rf} that the cosmic acceleration might originate from the violation of the cosmological principle, homogeneity and isotropy, i.e., the acceleration might be induced by the inhomogeneities of the universe. % The possible change of the deceleration parameter in an inhomogeneous universe has been pointed out in Refs.\\ \\onlinecite{Partovi:1982cg} and \\onlinecite{Mashhoon:1983}. % In particular, it has been shown that the luminosity distance-redshift relation indicated by the supernova data can be reproduced in an inhomogeneous cosmological model without introducing dark energy \\cite{Mustapha:1998jb,Celerier:1999hp,Celerier:2000ew,Tomita:2000jj,Iguchi:2001sq,Alnes:2005rw,Apostolopoulos:2006eg,Garfinkle:2006sb,Biswas:2006ub,Enqvist:2006cg}. %    The current situation of our knowledge (what we know and what we do not know well) about observations, cosmic acceleration and inhomogeneities is as follows.\\vspace{0.7em}  \\noindent --- Known --- \\\\ % $\\bullet$\\hspace{0.2em} \\textit{Based on % the FLRW cosmology, the current observational results indicate the existence of the cosmic acceleration.} (That is, by invoking the FLRW cosmological model to interpret the observational data, one would conclude that the expansion of the universe is accelerating in the recent epoch.) \\vspace{0.7em}  \\noindent --- Doubts --- \\\\ % (1) \\textit{Is the FLRW cosmology a good approximation?}\\\\ (2) \\textit{Do the current observational results indicate the existence of cosmic acceleration for the real universe with complicated energy distribution?} (Even the definition of accelerating expansion is an issue for our complicated universe. This issue will be discussed in Sec.\\ \\ref{def}.)\\\\ (3) \\textit{Can the inhomogeneities of our universe explain the observational results?}\\\\ (4) \\textit{Can the inhomogeneities of our universe generate accelerating expansion?}\\vspace{0.7em}  These four doubts are far from being fully answered. This is due to the difficulties from the complexity of the real energy distribution and the non-linearity of the Einstein equations. Because of these difficulties, instead of dealing with the full Einstein equations describing the real complicated universe, usually people employ the following two approaches to sketch the possible answers to the above doubts.% \\vspace{0.2em}  \\noindent % (I) \\textit{Taking perturbative approach with a simple background} (such as the homogeneous and isotropic RW background).\\\\ % (II) \\textit{Invoking exact solutions of the Einstein equations describing an inhomogeneous universe}.\\vspace{0.5em}  \\noindent % In addition to these two approaches, there are also non-perturbative studies without invoking exact solutions, e.g., see Refs.\\ \\onlinecite{Sicka:1999cb,Buchert:1999pq,Rasanen:2008it}.   Via Approach (I) with a homogenous and isotropic background, the positive answer to Doubt (1) is supported in Refs.\\ \\onlinecite{Siegel:2005xu} and \\onlinecite{Ishibashi:2005sj}, and the negative answer to Doubts (3) and (4) supported in Refs.\\ \\onlinecite{Siegel:2005xu,Ishibashi:2005sj,Kasai:2006bt,Vanderveld:2007cq}. Nevertheless, the reliability of the perturbative approach in (I) is doubtful for investigating the late-time cosmic evolution, for which not only energy distribution, but also curvature, may not be described with perturbations \\cite{Kolb:2005da,Rasanen:2006kp}. Furthermore, the arguments of Refs.\\ \\onlinecite{Siegel:2005xu,Ishibashi:2005sj,Kasai:2006bt,Vanderveld:2007cq} have been countered in Refs.\\ \\onlinecite{Kolb:2005da,Kolb:2005ze,Rasanen:2006kp,Rasanen:2008it}.  The drawback of Approach (II) is that the exact solution invoked may be very different from the real situation of our universe. From some angle, instead of Doubts (3) and (4), Approach (II) is to answer a more general question:% \\vspace{0.2em}  \\noindent % $\\bullet$\\hspace{0.2em} \\textit{Can inhomogeneities explain observational results and generate accelerating expansion?}% \\vspace{0.2em}  \\noindent % Or, more precisely,\\vspace{0.3em}  \\noindent (3$^{\\prime}$) \\textit{Does there exist a universe (maybe very different from ours) in which the inhomogeneities can explain the observational results?}\\\\ (4$^{\\prime}$) \\textit{Does there exist a universe in which the inhomogeneities can drive the expansion to accelerate?}\\vspace{0.5em}  \\noindent % Via Approach (II), for Doubt (3$^{\\prime}$) it has been shown in Refs.\\ \\onlinecite{Mustapha:1998jb,Celerier:1999hp,Celerier:2000ew,Tomita:2000jj,Iguchi:2001sq,Alnes:2005rw,Apostolopoulos:2006eg,Garfinkle:2006sb,Biswas:2006ub,Enqvist:2006cg} that the supernova data can be explained by invoking the Lemaitre-Tolman-Bondi solution \\cite{Lemaitre:1933qe,Tolman:1934za,Bondi:1947av} that describes a spherically symmetric (but inhomogeneous) dust fluid.  In the present work, we take Approach (II) for investigating the general possibility of generating accelerating expansion via inhomogeneities, i.e.\\ Doubt (4$^{\\prime}$). % Against the common intuition in Eq.\\ (\\ref{consensus}) we find examples of accelerating expansion in the case of a spherically symmetric dust fluid described by the Lemaitre-Tolman-Bondi (LTB) solution, thereby giving support to the positive answer to Doubt (4$^{\\prime}$). % In addition to our examples, Kai et al.\\ \\cite{Kai:2006ws} also found acceleration examples based on the LTB solution, where there exists singularity around the center of the spherically symmetric system during the accelerating epoch. In contrast, in our examples the system is smooth everywhere and no singularity is involved. (For more comparison, see Sec.\\ \\ref{dom}.) % The possibility of the accelerating expansion in the LTB model was also pointed out by Paranjape and Singh in Ref.\\ \\onlinecite{Paranjape:2006cd}, in which numerical models exhibiting acceleration were constructed by an approximation where the contribution of 3-curvature dominates over the matter density.    There also existed acceleration examples for a system consisting of two or more regions \\cite{Nambu:2005zn,Rasanen:2006kp,Rasanen:2006zw,Rasanen:2008it,Wiltshire:2007fg,Paranjape:2008ai}. In these acceleration examples, connecting the separate regions smoothly and the effect of the junction between these regions are the essential issues yet to be seriously explored. % We particularly note that it needs much caution to connect two (or more) regions. In many cases the effect from the junction between separate regions is significant and should not be ignored. % A similar doubt was also raised by Paranjape and Singh in Ref.\\ \\onlinecite{Paranjape:2008ai}. % As a demonstration of how things may go wrong when the connection or the junction is not appropriately taken care, in Appendix \\ref{app:NT_example} we investigate in detail the acceleration examples studied by Nambu and Tanimoto in Ref.\\ \\onlinecite{Nambu:2005zn}, and show that actually there should be no acceleration in those cases when we properly connect the separate regions and seriously take the effect of the junction into account. %  The acceleration examples we find could be far away from the real situation of our universe. We are not proposing to employ these mathematical examples or the LTB solution to describe our universe. These acceleration examples, which provide positive answer to Doubt (4$^{\\prime}$), are to demonstrate how inhomogeneities can drive the expansion to accelerate, thereby showing how our intuition about the interplay of gravity and the cosmic evolution may go wrong [i.e.\\ against the common intuition in Eq.\\ (\\ref{consensus})]. Accordingly, the effects of inhomogeneities on the evolution of the universe should be carefully restudied.  This paper is organized as follows. In Sec.\\ \\ref{def}, a tricky issue of the definition of acceleration is discussed and two definitions to be utilized for searching acceleration examples are introduced. In Sec.\\ \\ref{ltb}, the LTB solution is described. In Sec.\\ \\ref{rad} and Sec.\\ \\ref{dom}, we present the examples of the accelerating expansion corresponding to these two definitions of acceleration, respectively. A summary and discussions follow in Sec.\\ \\ref{con}.  Throughout the present paper we will use the units where $c = 8 \\pi G = 1$. We note that in this unit system there is still one unit unspecified, and, as a result, the value of one of the dimensionful quantities can be arbitrarily set. For example, in the acceleration examples we will present, the physical size of the spherical region under consideration can be any length long (such as $1\\,$fermi, $1\\,$cm, $1\\,$pc, $1\\,$Mpc, Hubble length $H_0^{-1}$, etc.). Once the value of one of the dimensionful quantities (which are independent of $c$ and $G$) is settled, all the units for the dimensionful quantities studied in the present work are specified. ",
        "conclusions": "\\label{con}  Against the common consensus [in Eq.\\ (\\ref{consensus})] that normal matter always slows down the expansion of the universe, we have found and demonstrated examples % of the line acceleration and the domain acceleration for a spherically symmetric dust fluid described by the LTB solution without singularity. This discovery contradicts the common intuition about the interplay of gravity and the cosmic evolution. Furthermore, these examples have shown the strong correlation between acceleration and inhomogeneity. These results strongly support the suggestion that inhomogeneity can induce acceleration.  In these acceleration examples the spherically symmetric dust fluid consists of three regions: two (roughly) homogeneous regions --- the inner over-density region and the outer under-density region --- and one transition/junction region, where the inhomogeneity locates, of two homogeneous regions. Note that in these examples the energy density distribution is flat around the origin $r=0$ and therefore there is no singularity and, moreover, no cusp behavior (or weak singularity).  For further understanding the acceleration, the quantity $\\partial_t^2 \\sqrt{g_{rr}}\\,$, which characterizes the acceleration/deceleration status of an infinitesimal radial line element, is one of the good quantities to study. Naively, the regions with positive/negative $\\partial_t^2 \\sqrt{g_{rr}}$ make positive/negative contribution to acceleration. For the quantity $\\partial_t^2 \\sqrt{g_{rr}}$ to be positive we find a necessary and sufficient condition: $(a^2/r)_{,r}>0$. This condition tells us that a positive contribution to acceleration is made in the place where $a(t,r)$ increases sufficiently fast with the radial coordinate $r$. In every acceleration example we found, there exists a region with positive $\\partial_t^2 \\sqrt{g_{rr}}$ that coincides with the inhomogeneous transition/junction region quite well. This result reveals the strong correlation between acceleration and inhomogeneity, thereby giving a strong support to the suggestion that inhomogeneity can induce accelerating expansion.    It is not clear how the line acceleration and domain acceleration are related to the apparent acceleration. A similar doubt (regarding the domain acceleration) was also raised by Kai et al.\\ in Ref.\\ \\onlinecite{Kai:2006ws}. The relation between the theoretical accelerations --- line acceleration and domain acceleration --- and the apparent acceleration could be model-dependent. For our universe with complicated energy distribution the relation may be far from simple. % As a reminder, we note that the conclusion about the existence of the present cosmic acceleration indicated by observations is based on the homogeneous and isotropic FLRW cosmology. If one invokes inhomogeneous cosmology, the current observational results may not indicate the existence of the cosmic acceleration. % As shown in Appendix \\ref{app:E=0} and Ref.\\ \\onlinecite{Paranjape:2006cd}, in the cases where $E(r)=0$ there is no acceleration (for both the line and the domain acceleration). In contrast, it has been shown that it is possible to fit supernova data in the LTB model with trivial $E(r)$ \\cite{Celerier:1999hp}. This gives an example of having apparent acceleration while having no line acceleration and no domain acceleration. % One may treat a complicated universe as a large domain consisting of many different sub-domains, and from a statistical perspective it might be reasonable that the size $L_D$ of the large domain corresponds to the scale factor in the FRW metric. In many works the results from the analysis involving the quantity $L_D$ are compared with observations in this spirit \\cite{Buchert:1999pq,Rasanen:2006zw,Rasanen:2006kp,Wiltshire:2007fg,Li:2007ny,Rasanen:2008it}. Nevertheless, whether this is a good approximation is not clear yet.       One may wonder whether the acceleration in our examples is real or fake, i.e., whether it truly corresponds to the space expansion. One example leading to this concern is to consider an accelerating system consisting of two decelerating regions. Since each region is decelerating one might expect that the domain acceleration in this case does not correspond to physically observable attributes (e.g.\\ the luminosity distance-redshift relation) of an accelerating FLRW model \\cite{Ishibashi:2005sj}. Nevertheless, on the contrary, it has been shown that for a system consisting of decelerating regions it is possible to fit the supernova data or to generate apparent acceleration \\cite{Celerier:1999hp,Kolb:2005ze}.   In addition, as already pointed out in Introduction, it needs much caution to connect separate regions or to put them together. Two particular important issues (that may be closely related to each other) are: \\\\ (1) There may be singularity in the junction between separate regions. \\\\ (2) The effect from the junction between separate regions may be significant and should be taken into account seriously. \\\\ We have already had a lesson from the work by Nambu and Tanimoto in Ref.\\ \\onlinecite{Nambu:2005zn}, where they basically considered a system (of spherical symmetry) consisting of an inner and an outer decelerating FLRW region and, as a result, they found and presented examples of the domain acceleration. In this system there is singularity in the junction between these two regions, and in their treatment the contribution from the junction is ignored. As pointed out in the previous section (with details in Appendix \\ref{app:NT_example}), it is inappropriate to ignore the junction. Actually the junction makes significant contribution to the volume of the system, although it looks like a ``2D'' surface in the coordinate space. After taking care of the singularity and taking into account the volume of the junction, we find that there exists no domain acceleration in the cases studied in Ref.\\ \\onlinecite{Nambu:2005zn}.   Another concern is about the definition of acceleration, as discussed in Sec.\\ \\ref{def}. The two definitions employed in the present paper --- line acceleration and domain acceleration, on which our acceleration examples are based --- follow the scenario proposed in Sec.\\ \\ref{def} for a system of freely moving particles interacting with each other only through gravity. For the purpose of truly representing the space evolution status while avoiding the confusion from particle motion and fake frame acceleration, the length quantities invoked in this scenario to define acceleration are (1) the distance between two freely moving particles for the line acceleration and (2) the size of a spatial region with constant number of particles therein for the domain acceleration. Regarding the frame/gauge choice, we have emphasized the advantage of the synchronous gauge in which the above requirements for the length involved in the acceleration definition can be easily met. In addition, in this gauge, there exists a universal cosmic time that is the proper time of comoving observers, and for a region with its boundary fixed in the coordinate space the volume expansion rate coincides with the expansion rate of the physical proper volume.     The counter-intuitive acceleration examples we found raise two issues worthy of further investigations: % \\vspace{0.1em}  \\noindent % \\hspace*{0em} $\\bullet$ \\textit{How to understand these counter-intuitive examples?}\\\\ % \\hspace*{0em} $\\bullet$ \\textit{Can inhomogeneities explain ``cosmic acceleration''?} % \\vspace{0.3em}   Regarding the first issue, the common intuition about Eq.\\ (\\ref{consensus}) may actually stem from Newtonian gravity that is a perturbed version of general relativity in the Newtonian limit with the Minkowski space-time as the background. Accordingly, this intuition may be valid only for considering the particle motion relative to a background space-time in a perturbative framework, but may be invalid for considering the evolution of space-time that is described by general relativity. For further understanding the cosmic evolution and other topics involving the general-relativity effects as the dominant effects, it will be helpful if one can build the intuition about general relativity, i.e., with which one can make a proper guess at the behavior of the space-time geometry for an arbitrarily given energy distribution.     Regarding the second issue, the most important is whether the inhomogeneities of our universe can explain the supernova data. If yes, the next step is to see, according to observational data, how the universe evolves, in particular, whether the cosmic acceleration exists or not, in this cosmological model taking the real inhomogeneities into account. So far the examples we found may be far away from the real situation of our universe. How to benefit from these mathematical examples in order to understand the present cosmic acceleration through the inhomogeneities of our universe is an important issue currently under our investigations. No matter whether the comic acceleration can eventually be explained simply by inhomogeneity, these examples have shown that inhomogeneity may affect the cosmic evolution in a manner far beyond the usual naive intuition about the interplay of gravity and the space-time geometry. Thus, the effects of inhomogeneity on the cosmic evolution should be restudied carefully. These examples open a new perspective on the understanding of the evolution of our universe."
    },
    "0512/astro-ph0512467_arXiv.txt": {
        "abstract": "Recent observational studies have revealed star-to-star  abundance inhomogeneity among light elements (e.g., C, N, O, Na, and Al) of stars on the main sequence in the Galactic globular clusters (GCs). One of promising interpretations for this result is that the observed abundance inhomogeneity is due to the second generation of stars formed from ejecta of the first generation  of evolved  stars (e.g., AGB stars) within GCs. However it remains unclear whether and how this primordial pollution can occur within GCs. We here propose a new scenario in which primordial pollution of GCs is highly likely to occur if GCs are located in the central regions of high redshift dark matter subhalos that can host low-mass dwarf galaxies. In this scenario,  gas ejected from the first generation of stars of GCs can be effectively trapped in the deep gravitational potential of their host halos and consequently can be consumed for the formation of the second generation of stars without losing a significant amount of gas by ram pressure stripping of interstellar and intergalactic medium. During merging of these halos with the proto-Galaxy, the halos are completely destroyed owing to the strong tidal field of the Galaxy. The self-polluted GCs located initially in the central regions of the halos can survive from tidal destruction owing to their compactness and finally become the Galactic halo GCs. In this scenario, ejecta of field stars surrounding the central GCs can be also converted into stars within their host dwarfs and finally become the second generation of stars of GCs. We also discuss the origin of the difference in the degree of abundance inhomogeneity between different GCs, such as $\\omega$ Centauri and NGC 6752, in terms of the difference in physical properties between host halos from which GC originate. ",
        "introduction": "The origin of star-to-star abundance inhomogeneity observed in the Galactic globular clusters (GCs) has long been discussed based mainly on the following two working hypotheses: The primordial hypothesis and the mixing one (e.g., Cottrell \\& Da Costa 1981; Freeman \\& Norris 1981; Smith 1987; Suntzeff 1993; Kraft 1994; Gratton et al. 2004). The first hypothesis is that the observed inhomogeneity is due to the second generation of stars that were formed from  gas ejected from the first generation of evolved stars (e.g., AGB stars) of GCs (``primordial pollution'' scenario, e.g.,  Cottrell \\& Da Costa 1981). The second  is that the observed chemical inhomogeneity of GCs can result from {\\it internal processes} of stars, such as dredge-up of CN-processed material from inner hydrogen-burning regions (e.g., Smith 1987; Kraft 1994). Recent observational studies of stellar abundance of some Galactic GCs have revealed star-to-star  abundance inhomogeneity among less evolved stars on the main sequence, where deep mixing of chemical components are highly unlikely (e.g., Cannon et al. 1998 for 47 Tuc). These studies accordingly suggested that the  primordial pollution scenario is more promising than the mixing (or evolutionary) scenario (Da Costa et al. 2004; Gratton 2004 for a recent review).   One of key questions related to the primordial pollution scenario is whether and how ejecta mainly from AGB stars of the first generation of stars can be effectively trapped  within GCs and consequently converted into the second  generation of stars. Frank \\& Gisler (1976) showed that gas of GCs can be efficiently stripped by ram pressure of the Galactic halo gas for most GCs. Smith (1996) suggested that the stellar ejecta from the first generation of stars are likely to be lost entirely from GCs with small binding energies through energetic outflow of intracluster wind. Gnedin et al. (2002) demonstrated that $\\omega$ Cen  could  not enrich itself with heavy elements of AGB stars owing to efficient ram pressure stripping of the Galactic interstellar medium (ISM), if it formed and evolved in isolation. These previous studies thus appear to suggest that primordial pollution is not likely within  GCs evolving in isolation, though  numerical attempts have not yet been made to investigate the details of primordial pollution processes within GCs.    Recent numerical simulations have suggested that GCs can be formed in the central regions of dwarf galaxies embedded by low-mass dark matter halos (e.g., Bromm \\& Clarke 2002; Mashchenko \\& Sills 2005). Both observational and theoretical studies of GCs suggested that massive globular clusters such as $\\omega$ Cen and G1 were previously stellar nuclei (or nuclear star clusters) of nucleated dwarf galaxies (e.g., Freeman 1993; Bekki \\& Freeman 2003; Bekki \\& Chiba 2004). Owing to the deeper gravitational potential in the central regions of dark matter halos, primordial pollution processes can be quite different between GCs evolving in isolation and those in the nuclear regions of their host halos. Thus it is quite important and timely to discuss whether and how primordial pollution of GCs can proceed {\\it if GCs are within the central regions of low-mass dark matter halos that can host dwarf galaxy populations.}   The purpose of this Letter is to propose a new scenario in which primordial pollution can proceed very efficiently in GCs that are located in nuclear regions of their host halos at high redshifts. In this scenario, primordial pollution can proceed more efficiently in the nuclear GCs than in those evolving in isolation, because the ejecta of AGB stars are more effectively trapped in GCs (without being significantly lost through ram pressure stripping and energetic outflow of evolved stars) owing to the deeper gravitational potential of their host halos. The nuclear, self-polluted GCs can appear as the Galactic halo GCs when the host halos merge with the proto-Galaxy and are subsequently destroyed by the strong tidal field of the Galaxy (i.e., when field stars and dark matter halos surrounding the GCs are all removed by the tidal stripping).  We firstly show the three advantages of this scenario in explaining why ejecta of the first generation of stars can be effectively trapped and converted into stars. Then we discuss the origin of the difference in the degree of abundance inhomogeneity between GCs in the Galaxy in the context of the proposed chemical pollution scenario. Since many authors have already discussed advantages and disadvantages of the primordial pollution (or enrichment) processes {\\it within isolated proto-GC clouds and  GCs} (e.g., Cayrel 1986; Parmentier \\& Gilmore 2001; Thoul et al. 2002; Recchi \\& Danziger 2005), we here do not intend to discuss these points.      \\begin{figure} \\psfig{file=f1.eps,width=8.5cm} \\caption{ The radial dependences of the escape velocities ($V_{\\rm esc}$) for four different models: the SB model with $M_{\\rm h}=10^8 {\\rm M}_{\\odot}$ and $r_{0}=0.48$ kpc (thick solid), the SB one with $M_{\\rm h}=10^9 {\\rm M}_{\\odot}$ and $r_{0}=1.28$ kpc (thick dotted), the NFW model with $M_{\\rm h}=10^8 {\\rm M}_{\\odot}$ and $r_{\\rm s}=0.16$ kpc (thin solid), and the NFW one with $M_{\\rm h}=10^9 {\\rm M}_{\\odot}$ and $r_{\\rm s}=0.52$ kpc (thin dotted). For comparison, the maximum value of the wind velocity of AGB stars is shown by a dashed line. } \\label{Figure. 1} \\end{figure} ",
        "conclusions": "\\subsection{Defunct dwarfs as GC hosts}  A key question in this scenario is whether GC host dwarfs embedded in dark matter halos can be completely destroyed by the strong Galactic tidal field without the central GCs being destroyed. Recent numerical simulations have demonstrated that nuclear star clusters can survive from the strong tidal field of the Galaxy owing to their initial compactness whereas the main bodies of (nucleated) dwarfs are completely destroyed and dispersed into the Galactic halo region (e.g., Bekki \\& Freeman 2003; Tsuchiya et al. 2003). Previous simulations also showed that nucleated dwarfs can be more efficiently converted into massive star clusters (i.e., naked nuclei) under strong tidal field of groups and clusters of galaxies for the SB profiles of dark matter halos (Bekki et al. 2003). These theoretical works therefore strongly suggest that the present scenario is promising.   Recent abundance studies of stars both for the Galactic halo and for dwarfs (e.g., dwarf spheroidal, dSph)  in the Local Group have revealed that [$\\alpha$/Fe] ratios of most stars in the dwarfs are generally lower than similar metallicity Galactic halo stars (e.g., Venn et al. 2004). The higher [$\\alpha$/Fe] ratios  of the Galactic halo can be due to {\\it very early merging of low-mass dwarf galaxies}, which were destroyed to form the Galactic halo without efficient chemical enrichment of stars (Venn et al. 2004). We propose  that these dwarfs destroyed in the very early history of the Galaxy (i.e., defunct dwarfs) are host galaxies of the Galactic halo GCs with abundance inhomogeneity. This proposal is consistent with the observed higher [$\\alpha$/Fe] ratios in the Galactic GCs (e.g., Freeman 1993).    \\subsection{Diversity in primordial pollution processes}  Smith (1987) divided GCs into three classes according to (1) whether star-to-star abundance inhomogeneity can be seen for Fe-peak element and (2) whether it can be seen in CN abundance (i.e., whether GCs show bimodal CN distributions). It was found in Smith (1987) that (1) only $\\omega$ Cen and M22 show inhomogeneity in Fe-peak elements, (2) GCs with no bimodal CN distributions have lower metallicities of $-2.2 \\le {\\rm [Fe/H]} \\le -1.6$ (e.g., M92 and M15), and (3) GCs with bimodal CN distributions have higher metallicities of ${\\rm [Fe/H]} > -1.6$ (e.g., NGC 6752 and 47 Tuc). However it remains unclear why  $\\omega$ Cen shows such inhomogeneity in Fe-peak element.   The present study can provide the following  answer for the above question, by assuming that the observed inhomogeneity is solely due to the primordial pollution processes of GCs. $\\omega$ Cen was formed in the nuclear region of a massive subhalo ($\\sim 5 \\times 10^8 {\\rm M}_{\\odot}$) that was virialized at very high redshift ($z \\sim 15$) and was more massive than any other subhalos hosting the Galactic GCs at that redshift. This massive halo with a higher mass density (due to earlier virialization) could not be soon destroyed by the proto-Galactic tidal field because of its stronger self-gravity, and consequently star formation could continue for a longer time scale in its nuclear region. As a result of this, not only AGB ejecta but also some fraction of metal-enriched gas from Type II and Type Ia supernovae  could  be finally recycled and converted into new generations of stars within the $\\omega$ Cen's host halo. The initially large apocenter distance of its orbit, for which a longer time scale of dynamical friction is required for the $\\omega$ Cen's host to reach the Galactic central region,  can be also responsible for the  prolonged star formation activity (Bekki \\& Freeman 2003).   Less massive halos containing  GCs (e.g., NGC 6752) other than $\\omega$ Cen had shallower gravitational potential and were more easily destroyed by the tidal field of the proto-Galaxy in their earlier histories. Only ejecta from more massive AGB stars with smaller wind velocities (an order of $\\sim 10$ km s$^{-1}$) therefore could be converted into new stars to become the second generation of stars in GCs  without efficient chemical enrichment (i.e.,  without significant metallicity spread) due to SNe II with very large wind velocities (an order of $\\sim 1000$ km s$^{-1}$). Less massive halos with smaller pericenter and apocenter distances of their orbits would not have experienced efficient chemical pollution of GCs owing to more rapid destruction of the halos. Thus $\\omega$ Cen  can show abundance inhomogeneity both in heavier elements (e.g., Fe) and in light ones (e.g., C, N, and O) whereas other GCs can show  abundance inhomogeneity only in light ones.    \\subsection{External pollution scenario}   We have shown that the AGB ejecta can be retained in the central regions of their host dwarfs and consequently converted into stars that finally become the second generation of stars in GCs. We also have suggested  that AGB ejecta of field  stellar populations surrounding the nuclear GCs can be converted into stars to become the second generation of stars in GCs. Accordingly, it would be reasonable to say that ``external pollution'' (or modified version of the original primordial pollution scenario) rather than ``self-pollution'' is a more reasonable jargon that denotes the chemical pollution processes of GCs in the central regions of dwarfs. This external pollution scenario may well  explain not only the observed helium overabundance of $\\omega$ Cen (Bedin et al. 2004) but also the large fraction of CN-strong stars in GCs with abundance inhomogeneity (Bekki \\& Norris 2005).  One of potential  problems of this external scenario is that the physical mechanism for the selective pollution by AGB stars (not by SNe II) is not clearly understood. In order to discuss this point, we plan to investigate gas dynamics within the central $0.1-100$pc of dwarfs at very high redshifts by using chemodynamical simulations combined with the latest results of AGB yields derived by Campbell et al. (2005). It is also our future study to investigate whether this scenario  can explain the observed O-Na and Mg-Al anticorrelations of GCs. Smith \\& Norris (1982) suggested that if both the first and the second generations of stars are formed from ejecta of AGB stars with different mases,  the observed CN-bimodality in GCs can be explained. This modified primordial (and external) pollution scenario will be addressed by our future chemodynamical simulations in a more quantitatively.      Theoretical studies based on  cosmological simulations have just started extensive investigation on structural and kinematical properties of GC systems of galaxies in  a $\\Lambda$CDM Universe (e.g., Bekki 2005; Yahagi \\& Bekki 2005). These Mpc-scale simulations, combined with sub-pc scale ones on star formation within GCs, will provide more robust predictions on spatial distributions and kinematics of the Galactic GCs with different past histories of primordial pollution processes and thus be compared with corresponding observations (e.g., Carretta 2005)."
    },
    "0512/astro-ph0512184_arXiv.txt": {
        "abstract": "{ MAGIC (Major Atmospheric Gamma Imaging Cherenkov telescope) is presently the largest ground-based gamma ray telescope. MAGIC has been taking data regularly since October 2004 at the Roque de los Muchachos Observatory on the island of La Palma. In this paper the MAGIC telescope status, its performances and some preliminary results on observed gamma ray sources are presented. }  \\FullConference{International Europhysics Conference on High Energy Physics\\\\ July 21st - 27th 2005\\\\ Lisboa, Portugal}  \\begin{document} ",
        "introduction": "MAGIC (Major Atmospheric Gamma Imaging Cherenkov) telescope \\cite{Barrio_1988}, is presently the largest imaging air Cherenkov telescope in operation. MAGIC has a 17~m diameter, $f/D = 1$, parabolic reflector covering a total surface of $234~{\\mathrm m}^2$. The reflector dish is composed of 956 ($0.495 \\times 0.495~{\\mathrm m}^2$) diamond milled aluminum mirrors \\cite{Bigongiari_2004}. The reflector shape is parabolic to minimize the time spread of Cherenkov light flashes on the focal plane. Aluminum mirrors were chosen instead of glass ones to reduce the weight of the reflecting surface and allow a fast slewing of the telescope. For the same reason the telescope frame is made of carbon fiber tubes. A very short slewing time is needed to catch the $\\gamma$-ray prompt emission by GRBs. The reflected Cherenkov photons are recorded by a $3.5^{\\circ}-3.8^\\circ$ FOV hexagonal camera in the telescope focal plane, composed by 397 $0.1^\\circ$ FOV photomultiplier tubes, surrounded by 180 $0.2^\\circ$ FOV PMTs. The PMTs have hemispherical windows and only 6 dynodes to minimize the time response width. The PMT photo conversion efficiency has been enhanced up to 30\\% and extended to the UV by coating the window with wavelength shifter \\cite{Paneque_2004}. The PMT signals are transferred via optical fibers \\cite{Paneque_2003} to the electronic room where they are split and sent to trigger and digitizing systems. The trigger decision is generated by a 2-level system using the signals of the 325 innermost PMTs. Only signals above an adjustable threshold are considered. The time coincidence within 6~nS of signals from 4 adjacent PMTs is required \\cite{Meucci_2004}. The analog signals are continuously digitized by 8 bit 300~MHz Flash ADCs. If the trigger condition is fulfilled the signals stored in FADC ring buffers are written to a FIFO buffer and saved by the DAQ. ",
        "conclusions": "MAGIC has been taking data regularly since Fall 2004. About 50 possible $\\gamma$-ray sources have been observed so far and $\\gamma$-ray emission has been detected by 8 of them. MAGIC confirmed $\\gamma$-ray emission from two sources recently discovered by HESS and detected $\\gamma$-rays from 1ES~1218+304 at $z = 0.182$ for the first time. MAGIC is able to measure the emission spectrum of sources which culminate close to zenith down to 100~GeV. This limit will be lowered further by improving the data analysis technique, adopting new 2.5~GHz FADCs and building a second telescope. The construction of the second telescope, already started, will be completed by the end of 2007."
    },
    "0512/astro-ph0512116_arXiv.txt": {
        "abstract": "We have carried out a high angular resolution near-infrared imaging study of the fields of 6 quasars with 7 strong  absorption line systems at $z < 0.5$, using the Hokupa'a adaptive optics system and the QUIRC near-infrared camera on the Gemini-North telescope.   These absorption  systems include 4 classical damped Lyman-alpha absorbers (DLAs), 2 sub-DLAs, and one Lyman-limit system. Images were obtained in the H or K$'$ filters with FWHM between $0\\farcs2 - 0\\farcs5$ with the goal of detecting  the absorbing galaxies and identifying their morphologies.   Features are seen at projected separations of $0\\farcs5-16\\farcs0$ from the quasars and all of the fields show features at less than 2\\arcsec separation. We find candidate absorbers in all of the seven systems.  With  the assumption that some of these are associated with the absorbers, the absorbers are low luminosity  $ \\le 0.1 L_{H}^{*}$ or $L_{K}^{*}$; we do  not find any large bright candidate absorbers in any of our fields.  Some fields show compact features that are too faint for quantitative morphology,  but could arise in dwarf galaxies. ",
        "introduction": "Damped Lyman-alpha absorption lines in quasar spectra are believed to arise from intervening galaxies and intergalactic matter at various cosmological epochs.  The damped Lyman-alpha absorbers (hereafter DLAs) are  classically defined  as quasar absorbers with $\\log N$(H~I)$ > 20.3$ while absorbers with $19.0 < \\log  N$(H~I)$< 20.3$ are conventionally classified as sub-DLAs.  This distinction is based on the observational constraints of an early spectroscopic study \\cite[]{Wolfe86}.  Since the Ly-$\\alpha$  line shows damping wings even at $\\log N$(H~I)$ \\sim 18$ in this paper we will refer  to both the sub-DLA and DLAs as DLA systems.  At high redshifts the DLAs are believed to  contain a large fraction of the co-moving density of neutral hydrogen in galaxies and  possibly account for all of the stars visible today \\cite[e.g.][]{Wolfe95,Peroux03}.  The evolution of metallicities in these absorbers provide important  probes of the  chemical enrichment and star formation history of the Universe \\cite[]{Khare04,Kulkarni05}.  Unfortunately the connection  between DLAs and galaxies has not been clearly established.  To shed more light on this connection, it is necessary to complement the  wealth of spectroscopic data on these absorbers with information on their morphologies, luminosities, colors, and image structure from direct imaging.  It has  proven hard to obtain this information for most DLAs.  The imaging of high-$z$ DLAs has been very difficult and a large fraction of the attempts to detect the Ly-$\\alpha$ emission from high-redshift  intervening ($z_{abs} < z_{em}$) DLAs have produced either  non-detections or weak detections \\cite[e.g.][]{Smith89,Hunstead90,Lowenthal95,Djorgovski96}.  Imaging studies of low-$z$ DLAs have been more encouraging.  Although not always spectroscopically confirmed to be the absorbers, galaxies in images  of low-redshift absorber fields often show a variety of morphologies: spirals, irregulars, low  surface brightness (LSB) galaxies (e.g., Steidel et al. 1994, 1995;  LeBrun et al. 1997; Bowen et al. 2001; Cohen 2001; Turnshek et al. 2001).  Most of these previous searches had limiting flux sensitivity thresholds of $\\sim 0.2L^{*}$ and thus could not rule out LSBs, while all of the near-IR searches lacked adequate angular  resolution to rule out dwarf galaxies close to the line of sight.  It is not clear which  of the several competing scenarios for DLAs are valid: large, bright,  rotating proto-spirals (Wolfe  et al. 1986; Wolfe \\& Prochaska 1998; Prochaska \\& Wolfe 1997b, 1998),  gas-rich dwarf galaxies (York et al. 1986; Matteucci et al. 1997),  merging proto-galactic fragments with cold dark matter (e.g., Haehnelt, Steinmetz, \\& Rauch 1998; Maller et al. 2001), collapsing  halos with merging clouds   (e.g., McDonald \\& Miralda-Escud\\'e 1999), or low-surface brightness  galaxies (Jimenez, Bowen, \\& Matteucci 1999).   Here we present the first adaptive optics observations of low-redshift DLAs. We have obtained near infrared images of seven absorbers at $ 0.1 < z < 1.3$ with the University of Hawaii Hokupa'a adaptive optics system and near-infrared camera QUIRC on the Gemini-North telescope.   We discuss the observations and data reductions in Section 2.  The  analysis of the data are presented in Section 3 and the results from individual fields are discussed in Section 4.  Finally, in Section 5 we characterize our sample of low-redshift DLAs based on our measurements of the sizes, impact parameters, and image structure.   Throughout this paper we assume $\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $h =0.73$. ",
        "conclusions": "We present the first adaptive optics observations of low-redshift DLAs. The images have revealed several objects at close angular separations  to the quasar in each field.   The adaptive optics images are comparable to the HST images in resolution and several close features are seen in common with HST and with these adaptive optics images.  In addition, we report the detection of two previously unidentified objects in the fields of Q0738+313, where there are no HST images, and Q1127-145, where the HST detection is marginal.  The objects found around in these quasar absorption fields would be less than $0.1L_{*}$ if they are at the absorber redshift and most of the brighter objects appear to have disks. The census of the brighter objects in these six absorber fields is 6 disk-dominated galaxies, 3 disk+bulge galaxies,  2 bulge-dominated galaxies, 3 point sources, and 5 unconstrained objects.  In addition, five of the six fields show objects between 0\\farcs5 and 1\\farcs0 to the line of sight to the quasar.  Our census has found likely candidates for all of the DLA systems.  The KL subtraction reveals a candidate object just offset from the quasar line of sight in Q0054+144 though the HST field appears to have several faint objects distributed about the field.  The DLA in Q0235+164 appears to be the object previously identified 1\\arcsec East of the quasar \\cite[]{Yanny89}.  In Q0738+313 we find a new object to which we attribute the lower-redshift DLA. This object would have an impact parameter of $\\sim$3 kpc.  It appears to have emission extending several arcseconds to the NW.   This emission could be associated  with the jet and arm features identified by \\cite{Turnshek01} though this  emission is fainter than the new object detected here.  The DLA at $z=0.22$ in this field has been previously identified.  In Q0850+440 \\cite{Lanzetta95} find a dwarf galaxy 9\\arcsec~ from the quasar ($b\\sim25$ kpc).  We identify another object very close to the quasar line of sight as a candidate absorber.  It appears only after subtraction of the quasar but if the absorption arises from this object, then the DLA arises close to the core of a very blue galaxy ($b\\sim2$ kpc).  For Q1127-145 we find a faint diffuse  object close to the line of sight of the quasar and extending NNW several arcseconds. The absorber in Q1329+412 is identified as arising from  an object 0\\farcs7 south of the quasar.  This object is also clearly seen in both our H-band image and an HST F702W image of the field.  All candidate absorbers are faint, with luminosities less than 0.1 $L_{H}^{*}$ or $L_{K'}^{*}$.  Assuming that at least some of these objects are at the same redshift as the absorbers, we conclude that the  absorbers in our fields are associated with relatively low luminosity galaxies. Morphological analysis reveals that most of the brighter objects have a disk component.  Their sizes, inferred from the surface brightness profiles, range from small to typical scale lengths for local disk galaxies.   For reference, our Galaxy has a disk scale length of $r_d = 2.2$~kpc measured in the $K$-band \\cite[]{maihara78,jones81} and M31 has a scale length of $r_d = 3.9$~kpc in the $K$-band, and $r_d = 4$~kpc in the $I$-band \\cite[]{hiromoto83}.  Table 5 summarizes object morphologies and the  derived linear impact parameter, luminosity, and scale lengths  assuming the objects are at the redshift of the absorber.  \\cite{Rao03} have suggested  that the DLAs at $z < 1$ are dominated by dwarf or low surface brightness  galaxies. However, \\cite{Chen03}, with more photometric redshifts,  have suggested that the luminosity function of $z < 1$ DLAs could be much broader.  Our observations, at higher resolution than both of these studies, have  found all of the candidate absorbers to be faint, with significant  disk components for the majority of the objects. This suggests that  a considerable fraction of low-$z$ DLAs may be faint, low surface brightness galaxies.  Such a conclusion would appear to be consistent with the low metallicities found  in low-$z$ DLAs (e.g., \\cite{Khare04}; \\cite{Kulkarni05}; and references therein). However, it would be necessary to obtain redshift confirmations for our  candidates and to obtain similar high-resolution images of other low-$z$ DLAs to reach more definitive conclusions on the luminosity function of the absorber galaxies.  Our observations  have demonstrated the use of adaptive optics for direct high-resolution imaging  of the galaxies giving rise to quasar absorbers. Deeper observations of the same fields in the future with higher order  AO systems would help to improve the signal-to-noise ratios  in the fainter objects. Furthermore, adaptive  optics systems with laser guide stars are not constrained by the  need to have a bright guide star in the quasar field, and would thus be able to reach  higher redshift absorbers.  It is crucial to also obtain spectroscopy (or at least narrow-band imaging)  of all the fields to better constrain the redshifts of the  detected candidate absorbers. With spectroscopic PSF subtraction  procedures (such as those followed by \\cite{Moller00a}) it may be  feasible to even verify the redshifts of the objects located very close to the  line of sight of the quasar.  It is essential to expand the sample of high-resolution broad-band images,  followed with spectroscopic confirmations, for quasar absorbers at low and  high redshifts. Such a combination of high-resolution imaging and spectroscopic observations of quasar absorbers can give direct information  on their luminosities, sizes,  and star formation rates and thus the nature of these galaxies.  Performing such observations on different types of  quasar absorbers (e.g., DLAs, weak Mg II systems, C IV systems) may help to understand any trends between the absorption line  strengths and galaxy properties such as the luminosities and impact parameters.  Finally, a comparison of these properties of quasar absorbers  at low and high redshifts will allow us to study the evolution of the  absorbing galaxies with cosmological time and the connection between the absorbers and the present-day galaxies."
    },
    "0512/astro-ph0512320_arXiv.txt": {
        "abstract": "We review the observed properties of extremely hot hydrogen-deficient post-AGB stars of spectral type [WC] and PG1159. Their H-deficiency is probably caused by a (very) late helium-shell flash or a AGB final thermal pulse, laying bare interior stellar regions which are usually kept hidden below the hydrogen envelope. Thus, the photospheric element abundances of these stars allow to draw conclusions about details of nuclear burning and mixing processes in the precursor AGB stars. We summarize the state-of-the-art of stellar evolution models which simulate AGB evolution and the occurrence of a late He-shell flash. We compare predicted element abundances to those determined by quantitative spectral analyses performed with advanced non-LTE model atmospheres. A good qualitative and quantitative agreement is found. Future work can contribute to an even more complete picture of the nuclear processes in AGB stars. ",
        "introduction": "\\begin{figure*} \\begin{center} \\includegraphics[scale=.60]{fg1.eps} \\end{center} \\figcaption{\\label{fig:hrd} Complete stellar evolution track with an initial mass of $2\\msun$ from the main sequence through the Red Giant Branch phase, the Horizontal Branch phase to the Asymptotic Giant Branch phase, and finally through the post-AGB phase that includes the central stars of planetary nebulae to the final white dwarf stage. The solid line represents the evolution of a H-normal post-AGB star. The dashed line shows a born-again evolution of the same mass, triggered by a very late thermal pulse, however, shifted by approximately $\\Delta \\log \\teff = - 0.2$ and $\\Delta \\log L/\\lsun = - 0.5$ for clarity. The star shows the position of PG1159-035. } \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=.60]{fg2.eps} \\end{center} \\caption{Hot hydrogen-deficient post-AGB stars in the $g$--\\Teff--plane. We identify Wolf-Rayet central stars of early and late type \\citep[{[WCE], [WCL],} from][]{hamann:97}, PG1159 stars (from Table~\\ref{tabpg1159}) as well as two [WC]--PG1159 transition objects (Abell~30 and 78). \\Teff\\ for the [WC] stars is related to the stellar radius at $\\tau_{\\rm Ross}$=20. Evolutionary tracks are from \\citet{schoenberner:83} and \\citet{bloecker:95b} (dashed lines), \\citet{wood:86} and \\citet{2003IAUS..209..111H} (dot-dashed line) (labels: mass in M$_\\odot$).  The latter $0.604\\msun$ track is the final CSPN track following a VLTP evolution and therefore has a H-deficient composition. However, the difference between the tracks is mainly due to the different AGB progenitor evolution.}\\label{fighrd} \\end{figure*}  Post-AGB stars represent a relatively short-lived transition phase between Asymptotic Giant Branch (AGB) stars and white dwarf (WD) stars. All stars with initial masses between 1 and about 8~M$_\\odot$ become H- and He-shell burning AGB stars and end their lives as WDs with a carbon-oxygen core. The model evolution is shown in \\abb{fig:hrd} for initially $2\\msun$.  More massive Super-AGB stars up to $12\\msun$ may end as ONeMg WDs.  During the hottest phases, post-AGB stars are surrounded by a planetary nebula (PN), ionized material lost by the precursor AGB star. Canonical stellar evolution predicts that throughout all evolutionary phases these stars retain hydrogen-rich envelopes which, however, can become contaminated by processed material from the interior by dredge-up events occurring in the Red Giant Branch (RGB) and AGB stages. Such an evolution is described by the solid line in \\abb{fig:hrd}.  Already decades ago, the observation of central stars of planetary nebulae (CSPN) exhibiting emission-line spectra which are very similar to those of massive Wolf-Rayet stars with strong helium and carbon emission lines (i.e. spectral type WC) suggested the existence of hydrogen-deficient post-AGB stars \\citep[e.g.\\,][]{heap:75}. Later, the Palomar-Green Survey revealed a new spectral class of H-deficient post-AGB stars, the PG1159 stars, which are dominated by absorption lines of highly ionised He, C, and O \\citep{wesemael:85}. It is now believed that the PG1159 stars are descendants of the Wolf-Rayet CSPN and that the majority of them will evolve into WDs with helium atmospheres (i.e. non-DA white dwarfs).  The origin of the hydrogen-deficiency in [WC] and PG1159 stars is probably a late helium-shell flash, which means that a post-AGB star (or WD) re-ignites helium-shell burning and transforms the star back into an AGB star. This ``born-again AGB star'' phenomenon has been discovered in early stellar evolution modeling \\citep{fujimoto:77,schoenberner:79} and was later invoked to explain the H-deficiency observed in some hot post-AGB stars \\citep{iben:83a,herwig:99c}. A complete evolutionary track of such a born-again evolution is shown as a dashed line in \\abb{fig:hrd}. The late He-shell flash causes complete envelope mixing. Hydrogen is ingested and burnt and the surface chemistry of the star becomes dominated by the composition of the previous He/C/O-rich intershell layer (the region between the H- and He-burning shells of the AGB star, see \\kap{sec:agbevolution}).  In ``usual'', i.e. hydrogen-rich (post-) AGB stars intershell matter can be dredged up to the surface (so-called third dredge-up) and results in the pollution of the photosphere with intershell matter, that is, helium, 3$\\alpha$-burning products (C, O, Ne) as well as s-process elements which are produced by neutron capture in the He-burning environment. The study of s-process element abundances in AGB stars was and is the most important tool to reveal details of the physics of burning and mixing processes. It is essential to understand these processes, because they affect the yields from AGB stars which largely drive the chemical evolution of the Galaxy.  In contrast to these H-rich stars, [WC] and PG1159 stars do not exhibit just traces of intershell matter in their photospheres, but they are essentially made up of intershell matter. This is because the mass of the intershell is much larger than that of the hydrogen envelope and, thus, its composition dominates the mixture of both layers triggered by the late He-flash. These H-deficient post-AGB stars therefore offer the unique possibility to study intershell abundances \\emph{directly} and, hence, understand the physical processes leading to the composition.  \\begin{deluxetable}{lrcccccccccccll} \\tabletypesize{\\scriptsize} \\rotate \\tablecaption{Element abundances (mass fractions) in representative hydrogen-deficient post-AGB stars of different spectral classes. We discuss two possible evolutionary sequences: 1. [WCL] $\\longrightarrow$ [WCE] $\\longrightarrow$ [WC]--PG1159 $\\longrightarrow$ PG1159, and 2. RCB $\\longrightarrow$ extreme He-B-stars $\\longrightarrow$ He-sdO $\\longrightarrow$ O(He)\\label{tababu} } \\tablewidth{0pt} \\tablehead{ \\colhead{\\bf Spectral Class}&\\colhead{\\Teff}&\\colhead{\\logg}&&&&&&&&&&&&\\\\ \\colhead{Star}          &\\colhead{[K]}  &\\colhead{[cgs]}& \\colhead{H}& \\colhead{He}&  \\colhead{C} &\\colhead{N}& \\colhead{O}& \\colhead{F}&\\colhead{Ne}& \\colhead{Si} & \\colhead{S} & \\colhead{Fe}& note& ref } \\startdata \\multicolumn{9}{l}{\\bf{[WCL]}}\\\\ IRAS\\,21282   & 28\\,000&3.2&.10   &.43   &.46    & $<$.005       &.01   &              &   &$<$.001&&    & H present         &2\\\\ PM1-188       & 35\\,000&3.7&.01   &.39   &.47    &.01            &.07&&.03&.025&   &    & N present, Si high&2\\\\ He2-459       &77\\,000&4.4&$<$.02&.40   &.50    &               &.10   & &   & &   &    & typical He/C/O    &2\\\\ \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{[WCE]}}\\\\ NGC\\,1501     &134\\,000&6.0&      &.50   &.35&               &.15   &              &   &    &   &    & typical He/C/O    &1\\\\ Sanduleak 3   &140\\,000&6.0&      &.62   &.26    & .005          &.12   &&   &    &   &    & N present         &3\\\\ \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{[WC]--PG1159}}\\\\ Abell 78      &115\\,000&5.5&      &.33& .50   & .02           & .15&1.0\\,10$^{-5}$&   &    &&$<$.0001&Fe-deficient   &4,10,18\\\\ \\noalign{\\smallskip}\\multicolumn{9}{l}{\\bf{PG1159}}\\\\ HS1517+7403   &110\\,000&7.0&      &.85   & .13&$<$3\\,10$^{-5}$& .02  &              &   &    &   &    & C,O low           &5\\\\ HS2324+3944   &130\\,000&6.2& .17  &.35   & .42   &$<$.0003\\tablenotemark{a} &.06  & &   &    &   &    & H present         &11\\\\ PG1159-035&140\\,000&7.0&$<$.02&.33   & .48   & .001         & .17  &3.2\\,10$^{-6}$&.02&.00036&.0001&$<$.0003&typical He/C/O&6,7,9,10,19\\\\ PG1144+005    &150\\,000&6.5&      &.38& .57   & .015          & .016 &1.0\\,10$^{-5}$&.02&    &   &    & O low&8,9,10\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{RCB}}\\\\ RY Sgr        &7\\,250&0.7&6\\,10$^{-6}$&.98&.007& .003          & .0009&              &&.00040&.00045&.00020&            &12\\\\ \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{extreme He-B-stars}}\\\\ BD+10$^\\circ$2179&16\\,900&2.5&1\\,10$^{-4}$&.98&.02& .0008        & .0004&&   &.00013&.00007&.00063&            &13\\\\ \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{He-sdO}}\\\\ BD+37$^\\circ$442&60\\,000&4.0&$<$.001&.97&.025   & .003         &      &              &   &.00079& &    &&14\\\\ KS292         &  75\\,000&5.5& .32  & .65  &.023   & .013         &      &&   &    &   &    &                   &15\\\\ \\noalign{\\smallskip} \\multicolumn{9}{l}{\\bf{O(He)}}\\\\ K1-27         &105\\,000&6.5&$<$.05&.98&$<$.015& .017          &      &              &   &    &   &    & N   present &16\\\\ LoTr4         &120\\,000&5.5& .11  &.89   &$<$.010& .003          &$<$.03&&   &    &   &    & H,N present       &16\\\\ HS1522+6615   &140\\,000&5.5& .02&.97   & .01   &               &      &&   &    &   &    & H,C present       &16\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} {\\bf Sun}     && & .73 &.25   &.0029  & .00089        &.0079&5.0\\,10$^{-7}$&.0018&.00072&.00050&.0013&&17 \\enddata \\tablenotetext{a}{uncertain} \\tablerefs{ (1) \\citet{koesterke:97b}; (2) \\citet{leuenhagen:98}; (3) \\citet{koesterke:97}; (4) \\citet{werner:92}; (5) \\citet{dreizler:98a}; (6) \\citet{werner:91}; (7) \\citet{werner:96}; (8) \\citet{werner:91b}; (9) \\citet{werner:04}; (10) \\citet{werner:05b}; (11) \\citet{dreizler:98b}; (12) \\citet{asplund:00}; (13) \\citet{pandey:05}; (14) \\citet{bauer:95}; (15) \\citet{rauch:91}; (16) \\citet{rauch:98}; (17) \\citet{grevesse:01}; (18) \\citet{werner:03}; (19) \\citet{jahn:05} } \\end{deluxetable} ",
        "conclusions": ""
    },
    "0512/gr-qc0512017_arXiv.txt": {
        "abstract": "Although several models of $f(R)$ theories of gravity within the Palatini approach have been studied already, the interest was concentrated on those that have an effect on the late-time evolution of the universe, by the inclusion for example of terms inversely proportional to the scalar curvature in the gravitational action. However, additional positive powers of the curvature also provide interesting early-time phenomenology, like inflation, and the presence of such terms in the action is equally, if not more, probable. In the present paper models with both additional positive and negative powers of the scalar curvature are studied. Their effect on the evolution of the universe is investigated for all cosmological eras, and various constraints are put on the extra terms in the actions. Additionally, we examine the extent to which the new terms in positive powers affect the late-time evolution of the universe and the related observables, which also determines our ability to probe their presence in the gravitational action. ",
        "introduction": "\\label{sec:1} General Relativity has been an extremely successful theory of gravity for almost 100 years now. It has passed all tests on scales relevant for the solar system. However, all of these tests are actually relevant to the Post-Newtonian regime and not to the full version of the theory. At the same time, during the last decades, several indications have appeared, leading to the thought that maybe General Relativity is the effective part of a more general theory of gravitation, applicable only on solar system scales. Some of these indications are of theoretical origin and as an example one could name the difficulties in finding the quantum counterpart of general relativity and the fact that it has not yet been possible to include it in any grand unification scheme. The others are of observational origin, like the fact that relativity cannot explain the flat rotation curves of galaxies without introducing unseen dark matter and, of course, the fact that it cannot explain the current accelerated expansion of the universe without the introduction of both dark matter and dark energy. Also, in order to explain these cosmological observations, one has not only to introduce these exotic componets, but also to assume that they sum up to 96\\% percent of the total energy content of the universe. This, of course, does not sound appealing at all, since the nature of these components still remains a mystery, and this has lead to several attempts to modify or extend General Relativity in different ways, addressing the dark matter \\cite{mond,mond1,mond2,mond3,mond4,mond5} or the dark energy \\cite{faraoni,dvali,def,dvali2,ark} problem.  An alternative route is to modify General Relativity by abandoning the simplicity assumption that the action should be linear in the scalar curvature  $R$. So, instead of using the Einstein--Hilbert Lagrangian one can use a more general one, which depends on a generic function $f(R)$ (see \\cite{buh,bar1,bar2} for early studies). Hence, such theories are called $f(R)$ theories of gravity. Such modifications are not completely arbitrary as higher order terms in curvature invariants seem to be present in the effective Lagrangian of the gravitational field when quantum corrections or string/M-theory are considered \\cite{quant1,quant2,quant3,quant4,noji2,vassi}. Therefore, $f(R)$ theories seem also to address the theoretical concerns about General Relativity that were mentioned earlier. Phenomenological interest up until recently was restricted to the inflationary behaviour exhibited when positive powers of $R$ are present in the Lagrangian \\cite{staro}. However, it was seriously increased when it was shown that a term inversely proportional to $R$ can lead to late time expansion \\cite{capo, capo2,carroll}. Such theories, appealing as they may be, are unfortunately not free of problems. First of all they lead to fourth order differential equations which are difficult to attack. Additionally, it is doubtful whether they can pass the known solar system tests \\cite{chiba,cem} or whether they have the correct Newtonian limit \\cite{cem,dick,olmo,sot1}. It is possible that more sophisticated models may exhibit behaviour closer to the expected one \\cite{noji}, but this requires significant fine tuning of the various parameters. The most important problem of these models, however, is that they lead to unavoidable instabilities within matter in a weak gravity regime \\cite{dolgov}.  There is a further modification of gravity that one can consider which does not necessarily involve any modification of the action, but rather the use of a different variational principle. This variational principle, known as the Palatini formalism due to a historical misconception, since it was not introduced by Palatini but by Einstein himself, treats the metric and the affine connection as independent geometrical quantities. One has to vary the Lagrangian with respect to both of them to derive the field equations, in contrast with the standard metric variation, where the Lagrangian is varied with respect to the metric alone, and the connections are assumed to be the Christoffel symbols of this metric. It is also known that when the Einstein-Hilbert action is used, the Palatini variational principle leads to the Einstein equations, just like the standard metric variation \\cite{wald}. This is not true, however, for a more general action.  When used together with an $f(R)$ Lagrangian, the Palatini formalism leads to second order differential equations instead of the fourth order ones that one gets with the metric variation. At the same time, in vacuum, they straightforwardly reduce to standard General Relativity plus a cosmological constant \\cite{ferr}. This ensures us that, firstly, the theory passes the solar system tests, and secondly, that interesting aspects of GR like static black holes and gravitational waves are still present. It is also free of the instabilities discovered in \\cite{dolgov} for $f(R)$ gravity in the metric approach, which arise within matter in a weak gravity regime. Even though there was initially some debate concerning the Newtonian limit \\cite{meng2,barraco}, a recent paper \\cite{sot1} seems to have settled this matter, showing that these theories indeed have the correct behaviour. One more reason to study such a metric-affine variation is theoretical completeness. Both approaches give exactly the same result when applied to the Einstein--Hilbert action. Thus there is no real criterion so far about which one of them is better to use. Additionally, though, the Palatini variation seems to be more general since it yields GR without the need to pre-specify the relation between the metric and the connections. Finally it should be mentioned that the Palatini formalism is also closer to the picture of a Hilbert space, since is assumes that there are two sets of independent variables, and therefore seems more appealing when trying to quantize gravity.  In \\cite{vollick}, Vollick showed that using the action of \\cite{carroll}, which includes an $R^{-1}$ term, together with the Palatini variational principle can lead to a theory of gravity that predicts late time accelerated expansion of the universe. Several models with similar behaviour followed (for a short review see \\cite{meng3} and references therein). Later it was proved that including an $R^2$ term in the action could not drive inflation \\cite{meng} within this formalism. However, in \\cite{sot2} it was shown that this is not true for higher order terms in general, but merely that $R^2$  constitutes an exception due to the construction of the theory, and a model was presented in which both an $R^3$ term and an $R^{-1}$ term were present in the action. It was demonstrated that such a model can account both for early time inflation and late time accelerated expansion.  It is important, of course, to go beyond the qualitative results and use the numerous observations \\cite{Riess,sloan,2dFGRS,wmap} to get quantitative ones.  Such a study was performed in \\cite{ama}. Assuming that the gravitational action includes, besides the standard linear term, a term inversely proportional to $R$, the authors used four different sets of cosmological data to constrain it. These are the Supernovae type Ia gold set \\cite{Riess}, the CMBR shift parameter \\cite{bond}, the baryon oscillation length scale \\cite{baryonsdss} and the linear growth factor at the 2dF Galaxy Redshift Survey effective redshift \\cite{hawkins,tegmark}. However, as stated in the conclusions of \\cite{ama}, the restricted form of $f(R)$, including only a term inversely proportional to $R$, prevents the study from being exhaustive.  Even though several models of $f(R)$ gravity in the Palatini formalism have been studied, most interest was concentrated on those having terms inversely proportional to the scalar curvature. Bearing in mind that including an additional term proportional to a positive power of the curvature in the action leads to interesting phenomenology \\cite{sot2} and of course is equally, if not more, reasonable to have present in the action, here the behaviour of models with both positive and negative powers of $R$ will be investigated (we will be using the term positive powers of $R$ to imply positive powers additional to $R$ itself). By generalizing the results of \\cite{sot2} it will be shown that such models lead to curvature driven inflation and late time cosmic expansion. However, a more detailed study reveals that such a curvature driven inflation is not interesting but is actually problematic, because there is no way to put an end to it, and should therefore be avoided. Another important question is how to constrain terms in positive powers of $R$. This problem will be addressed here in detail. One of the possible constraints will arise if we want to avoid the undesirable curvature driven inflation mentioned before, but also more standard ones, related to the later evolution and the Newtonian limit will be considered. The results will also be used to check whether the presence of such terms could affect the results of \\cite{ama}. Additionally, a thorough discussion will take place, which can be considered relevant, not only to the results of this paper, but also to studies similar to \\cite{ama}.  The rest of the paper is organized as follows: In section \\ref{sec:2} the Palatini formalism is briefly reviewed. In section \\ref{sec:3} a model including both negative and positive powers of the scalar curvature in the gravitational action is examined. Its behaviour is investigated for different cosmological eras and the possible constraints for the positive power term coming from different origins are obtained. Section \\ref{sec:4} contains a thorough discussion of the physical consequences of the derived results. Also, the results of \\cite{ama} are interpreted and discussed and we examine, using the results of section \\ref{sec:3}, whether they can actually be consider general enough to include gravitational Lagrangians with positive powers in the scalar curvature as well.  Section \\ref{sec:5} contains conclusions. ",
        "conclusions": "\\label{sec:5}  Starting by setting up the formalism needed, a model with both positive and negative powers of the scalar curvature, additional to the standard Einstein--Hilbert term in the gravitational action, was studied. The exponents and the coefficients of the extra terms were left arbitrary and all eras of cosmological evolution were studied. It was shown that, even though such a model can lead to early-time inflation at large values of the scalar curvature and late-time expansion at small values, as shown in \\cite{sot2}, it is not possible to have a smooth passage from the former to the latter for any model. If the curvature dominates the expansion there seems to be no turning back to ordinary matter domination, so it is reasonable to require that this should only happen at late times. Avoiding such an inflationary regime however, does not necessarily imply that additional positive powers of the curvature should not be present, but merely that they should be seriously constrained. In that case one can also have the usual scenario where inflation is driven by an inflaton field. Other constraints were derived as well, for the coefficients of the positive power terms, by requesting that the modified Friedmann equation does not differ significantly from the standard one during big bang nucleosynthesis and that the correct Newtonian limit is obtained. Using the above constraints, it was shown that the late time evolution of the universe is not affected by the presence of additional positive power terms of the scalar curvature in the action.  As a conclusion, any result derived for a model with only negative powers of the curvature using late time observational data, as done in \\cite{ama} is general enough to be applied to a model including additional positive power terms as well. At the same time, however, since such tests are totally insensitive to the presence of positive power terms, they cannot be used to constrain them and therefore favour or disfavour their presence. To close I would also like to mention the following. The results of \\cite{ama} seem to disfavour the presence of a negative power of the curvature in the gravitational action by showing that $\\Lambda$CDM is the best fit model. However, the model that better fits the data, does not necessarily have to be the physically preferred one, unless it can also be theoretically motivated and justified. The $\\Lambda$CDM model comes with a burden, known as the coincidence problem, and this is the motivation for creating the numerous alternative models present in the literature.  In this sense, what we are interested in is not whether a model can fit the data better than the $\\Lambda$CDM one, but if we can produce a model which is theoretically motivated or at least theoretically explained and at the same time fits the current data reasonably well. Whether this is true for $f(R)$ theories of gravity in the Palatini formalism is a question that remains unanswered."
    },
    "0512/astro-ph0512050_arXiv.txt": {
        "abstract": "We present spatially resolved mid-IR spectra of NGC 1068 with a diffraction-limited resolution of 0.25\\arcsec using the Long Wavelength Spectrometer (LWS) at the Keck I telescope.  The mid-infrared image of NGC~1068 is extended along the N-S direction.  Previous imaging studies have shown the extended regions are located inside the ionization cones indicating that the mid-infrared emission arises perhaps from the inner regions of the narrow-line clouds instead of the proposed dusty torus itself.  The spatially resolved mid-IR spectra were obtained at two different slit position angles, +8.0 and -13.0 degrees across the elongated regions in the mid-IR.  From these spectra, we found only weak silicate absorption toward the northern extended regions but strong in the nucleus and the southern extended regions.  This is consistent with a model of a slightly inclined cold obscuring torus which covers much of the southern regions but is behind the northern extension.  While a detailed analysis of the spectra requires a radiative transfer model, the lack of silicate emission from the northern extended regions prompts us to consider a dual dust population model as one of the possible explanations in which a different dust population exists in the ionization cones compared to that in the dusty torus.  Dust inside the ionization cones may lack small silicate grains giving rise to only a featureless continuum in the northern extended regions while dust in the dusty torus has plenty of small silicate grains to produce the strong silicate absorption lines towards the nucleus and the southern extended regions. ",
        "introduction": "AGNs appear in a variety of types, which are often classified based on the presence of broad emission lines.  Objects are classified as type I if both broad and narrow emission lines appear in their optical spectra and as type II if only  narrow emision lines are present without broad emission lines. It is, however, a general belief that much of the observed diversity in the local universe arises from different viewing angles toward the central engine and a dusty toroidal structure around it, especially for Seyfert galaxies of types I and II.  When the dusty torus is viewed face-on, both the central engine and the broad-line regions can be seen directly causing objects to appear as Seyfert 1 galaxies.  When the dusty torus is viewed edge-on, the anisotropic obscuration created by the torus causes objects to appear as Seyfert 2 galaxies (see \\citealt{ant93} for review).  It is this crucial role played by dust in the Unified model of AGN that makes understanding dust properties very important in understanding AGN.  A significant fraction of the optical/UV/X-ray luminosity of the active nucleus is absorbed by the proposed dusty torus and reradiated at mid-infrared wavelengths.  The infrared also suffers less extinction than the optical band and is prefered for probing the proposed dusty torus or dust in general in the nucleus.  Early mid-infrared observations of Seyfert galaxies \\citep{roc91} have shown that the 9.7 $\\micron$ silicate feature appears in strong absorption in Seyfert 2's as expected from an edge-on geometry of the proposed dusty torus for type II objects.  If the dust responsible for the 9.7 $\\micron$ silicate feature belongs to the dusty torus, then the spatial distribution of the silicate absorption can provide very important clues to the physical properties (size, orientation, etc.) of the dusty torus.  It is, however, difficult to investigate the spatial distribution of dust in Seyfert galaxies because mid-IR emission from most Seyfert galaxies has not been resolved (see \\citealt{gor04}).  While thermal emission from hot dust has been considered as a dominant source of the mid-IR emission of Seyfert galaxies, the unresolved nature of the mid-IR image have left the discussion over a non-thermal origin still alive.  NGC~1068 is one of a few Seyfert galaxies whose mid-IR emission is resolved.  No other alternative mechanism like synchrotron radiation can produce emission over such an extended area, leaving heated dust grains as the likely source.  NGC 1068 is classified as a Seyfert 2 based on the presence of narrow emission lines and absence of broad emission lines.  The detection of broad emission lines in polarized light \\citep{ant85}, however, has shown that NGC 1068 harbors an obscured Seyfert 1 nucleus.  As one of the closest and brightest Seyfert 2 galaxies, NGC 1068 offers a better spatial scale for the detail investigatation of its obscured nucleus.  NGC 1068 is at 14.4 Mpc \\citep{tul98} so 1\\arcsec corresponds to 72 pc or 0.25\\arcsec (our spatial resolution) corresponds to 18 pc in physical distance.  As a part of a greater mid-infrared survery of a sample of Seyfert galaxies, NGC 1068 was selected specifically for the investigation of the physical properties of the proposed dusty torus.  NGC 1068 has been observed many times in the mid-infrared both in imaging and spectroscopic modes. The spatially resolved mid-IR images of NGC~1068 show a linear structure covering about 1\\arcsec in the north-south direction \\citep{bra93,cam93,boc00,tom01}.  Emission line imaging of O[III] showed that this disklike structure lies inside the narrow-line regions \\citep{eva91,mac94}. Since the narrow-line regions are believed to be created by the ionizing radiation from the central engine which is collimated by the dust torus, the dusty torus should be oriented in the East-West direction perpendicular to the narrow-line cones.  Thus this structure has been believed to be created by grains in the NLRs heated by the nuclear radiation.  In this model, the dusty torus is then too cold to emit significantly at 10 \\micron.  While recent mid-infrared imaging studies have provided some spatially resolved dust measurements in the nuclear region of NGC 1068, no spatially resolved spectra have been obtained in the mid-infrared until this study.  Mid-infrared spectra of NGC 1068 have also been obtained repeatly by several single aperture ground-based telescopes and Infrared Space Observatory (ISO) using various apertures from 0.4\\arcsec at VLTI to 24\\arcsec $\\times$ 24\\arcsec at ISO \\citep{lut00,lef01,jaf04,roc91,sie04}.  Most previous spectra show that the mid-infrared spectra of NGC 1068 have significant silicate absorption. Especially \\citet{stu00} and \\citet{lut00} show the AGN dominated spectra with no significant PAH emission lines.  A crude spatial map of the distribution of the silicate absorption has been provided by imaging studies using multiple narrow bands, most notably from \\citet{boc00} and \\citet{gal03}.  However, these studies have given conflicting results making spatially resolved spectra necessary to resolve the issues. For example, \\citet{boc00} has reported that the silicate feature is relatively strong in absorption on the nucleus and to the south but flat or even in emission to the north.  In contrast, \\citet{gal03} found that the silicate feature appears strong in absorption to the north but either flatter or in emission to the south.  For the first time, we present the diffraction-limited mid-infrared spectra of NGC 1068 at a spatial resolution of 0.25\\arcsec to investigate the spatial distribution and the properties of dust in the nuclear regions. ",
        "conclusions": "We used the Long Wavelength Spectrometer (LWS) at Keck I telescope with a 0.25\\arcsec slit to obtain spatially-resolved spectra of NGC 1068 at three different slit positions, slit PA = +8$^\\circ$, -13$^\\circ$, and +78.9$^\\circ$. Overall our integrated spectra agrees with the previous mid-IR spectra well. In NGC 1068, our resolved spectra showed that the silicate absorption was strongest on the nucleus (Central Engine) indicating that our line of sight transmited through the largest amount of cold obscuring dust materials to the nucleus.  Furthermore, the strength of the silicate absorption declines fast to the north but remains relatively strong to the south.  This asymmetry in the strength of the silicate absorption as well as the larger extended emission regions to the north suggests that the north pole of the torus opening is inclined towards us and that the southern extended regions are behind the obscuring dusty torus.  Based on the lack of strong silicate emission from the north, we considered the dual dust population model for AGN in which two different dust populations exist in the ionization cones and in the dusty torus.  The dust population in the ionization cones may lack small silicate grains or small grains in general and give rise to featureless continuum to the north while the dust population in the dusty torus contains plenty of small silicate grains and causes the apparent deep silicate absorption lines to the nucleus and the south in NGC 1068.  A detailed analysis using a radiative transfer model would provide a better understanding to our spatially resolved mid-IR spectra."
    },
    "0512/astro-ph0512099_arXiv.txt": {
        "abstract": " ",
        "introduction": "{\\bf A review of recent observations of the kinematics of six objects that represent the broad range of phenomena called planetary nebulae (PNe) is presented. It is demonstrated that Hubble--type outflows are predominant, consequently it is argued that ballistic ejections from the central stars could have dominated the dynamical effects of the fast winds in several, and perhaps all, of these objects. An alternative possibility, which involves an extension to the Interacting Winds model, is considered to explain the dynamics of evolved planetary nebulae.}  A consensus has been established (e.g. Kastner et al. 2003) about the basic processes for the creation of a planetary nebula (PN): an intermediate mass star (initial mass 1--8 \\msun) loses mass in its Asymptotic Giant Branch (AGB) phase at $\\leq$ 10$^{-4}$ \\msun\\ yr$^{-1}$ by emitting a `superwind' flowing at 10 - 20 \\kms\\ over $\\leq$ 10$^{5}$ yr (depending on the initial stellar mass). The star eventually becomes an 0.5 - 1.0 \\msun\\ White Dwarf (WD) which produces enough Lyman photons to ionise a substantial fraction of the circumstellar envelope recognisable as the expanding PN. The whole structure can be enveloped as well in a prior low density Red Giant (RG), similarly slow, wind. In the transition from the AGB to WD phase the outflow mass loss rate declines to 10$^{-8}$ \\msun\\ yr$^{-1}$ but increases its speed dramatically to several 1000 \\kms\\ to blow as a fast wind for an as yet unknown period. As the WD star evolves further the fast wind declines.  Obviously the real story is hugely more complicated in detail and variable between objects (Balick \\& Frank, 2002). Morphologies range from simple, spherical shells to complex poly-polar structures (e.g. NGC 2440, L{\\'o}pez et al. 1998) probably around close binary systems. The ejected, dusty, AGB, molecular superwind material is often very clumpy (e.g. the cometary globules of NGC 7293, Huggins et al 1992; Meaburn et al. 1992; O'Dell \\& Handron 1996; Meaburn et al. 1998). High--speed jets (e.g. IRAS17423-1755, Riera et al. 1995) and `bullets' (e.g. MyCn 18, Bryce et al. 1997; O'Connor et al. 2000) are found and even shell or lobe expansion velocities can range from  20 \\kms\\ to $\\geq$ 500 \\kms\\ (e.g. He2--111, Meaburn \\& Walsh 1989 and NGC 6302, Meaburn et al. 2005c). Sometimes many of these distinctly separate phenomena occur in one object reflecting the separate stages of its complex evolution. The ages of observed PNe range from the initial proto-PN stage to those of well--evolved PN $\\geq$~10$^{4}$ yr later around an ageing WD star.  In current dynamical theories of PNe  much emphasis is placed on the importance, even dominance, of the fast wind. In the elegant interacting winds (IWs) model and its variants (Kwock, Purton \\& Fitzgerald 1978; Kahn \\& West 1985; Chu et al. 1993; Mellema 1995 \\& 1997; Balick \\& Frank 2002) the shocked (10${^6}$ -- 10$^{8}$ K) fast wind can form an energy--conserving, pressure--driven `bubble' in the preceding smooth AGB wind whose density declines as distance$^{-2}$ from the star. This possibility is similar in principle to that pioneered by Dyson \\& de Vries (1972) albeit within  a stationary medium of uniform density. The characteristic shell of a simple PN is consequently formed between the shocks in the fast wind and AGB wind and, being pressure--driven by the superheated gas, is expanding faster than this ambient AGB outflow. A variant would have the momentum of the isotropic fast wind simply sweeping up the AGB outflow and accelerating an expanding shell. For the creation of a  bi-polar PN Cant{\\'o} (1978) and Barral \\& Cant{\\'o} (1981) considered something similar. Here a fast wind from a star embedded in a dense circumstellar disk forms cavities  on either side of it which are delineated by stationary shocks across which the fast wind refracts to form bi--polar, momentum--conserving, outflows parallel to the cavity walls. Again, energy conserving, elongated `bubbles', pressure driven by the shocked wind on either side of this disk, would also form expanding bi-polar lobes. Steffen \\& L{\\' o}pez (2004) examine the effects of the fast wind on a clumpy AGB wind which is more realistic than the smooth density distributions usually considered in the IWs models.  There is now an abundance of observational evidence that  fast winds exist within PNe and some evidence that they interact significantly with the circumstellar envelopes. Patriarchi \\& Perinotto (1991) discovered that 60 percent of central stars of PNe emit particle winds of 600--3500 \\kms. However, direct observational evidence of their interaction with the circumstellar medium is more limited. Collimated and truncated ablated flows, where the fast wind has mixed with, and is slowed by, photoionised gas evaporating from dense, stationary, globules (Hartquist et al. 1986; Dyson et al. 1989a; Dyson, Hartquist \\& Biro 1993) are detected in the hydrogen--deficient PNe A30 (Borkowski et al. 1995; Meaburn \\& L{\\'o}pez,1996) and A78 (Meaburn et al. 1998). Also, diffuse X--ray emission is found in the cores of five PNe ( NGC 7009, Hen 3--1475, BD+30$\\deg$ 3639, NGC 6543 \\& NGC 2392 -- Chu et al. 2001; Gruendl et al. 2001 Guerrero, Chu \\& Gruendl 2004; Chu et al. 2004; Guerrero et al. 2005). Similarly, Kastner et al. (2003) observed such diffuse X--ray emission in the core of the bi--polar PN Menzel 3 (Mz~3) and Montez et al. (2005) inside the main shell of NGC 40 which is a PN generated by a WR--type star. All of these authors interpret the X--ray emissions to be the consequences of the collisions of the fast winds with the slower moving surrounding AGB winds. They suggest that conductive cooling is occurring for the temperatures of the hot gases emitting the X-rays are far lower than if simply generated by shocks in the fast winds at their measured speeds. Nonetheless, they imply that over--pressured `bubbles' of super-heated gases are forming and driving the expansions of the ionised PNe shells as predicted by the IWs model.  The principal purpose of the present article is to examine, on the basis of observations made recently with the two Manchester echelle spectrometers (MES - Meaburn et al. 1984 and 2003), the part played by the fast winds in the creation of the well--evolved PNe, NGC 6853 (Dumbbell) and 7293 (Helix), the young PN, NGC~6543 and the outer lobes of the bi--polar (poly--polar) `PNe' NGC 6302, Mz~3 and MyCn~18 for these are all recently observed examples of the range of circumstellar phenomena broadly designated as PNe. ",
        "conclusions": "The IWs model certainly seems applicable to the very innermost shell of the young PN NGC 6543 (and very clearly to the main shell of NGC 40 -- see Sect.~1). However, considerable modification of this theory is needed to explain the current state of the evolved PNe, NGC 6853 and 7293. The fast wind could have switched off well before the 10$^{4}$ yr age of their expanding shells generating an inward acceleration of their inside surfaces. Alternatively, it remains possible that eruptive events over 10$^{4}$ yr have dominated any effects of the fast winds, active only for a few thousand years in these PNe, to create the Hubble--type expansions throughout their volumes that are currently observed.   The preponderance of Hubble--type outflows of the lobes of the bi--polar PNe, NGC 6302, Mz 3 and MyCn 18 (and see Corradi 2004 for other examples) invites the simplest interpretation; that they are consequences of ejections over short periods of time but of material with different speeds. Even when a fast wind is currently present (e.g. Mz 3) the outer high--speed lobes are shielded from it by inner shells.  The present article has been deliberately limited to considering the dynamics of a small number of objects whose motions have been well observed. It seems clear in this small sample that simple ballistic ejections, maybe involving close binary systems in some cases, could dominate the dynamical effects of the fast winds.  The author is grateful to Bob O'Dell for providing the \\hb\\ and \\heiis\\ images in Fig. 1 and to Romani Corradi for the image in Fig. 5."
    },
    "0512/astro-ph0512266_arXiv.txt": {
        "abstract": "{ We report  on the  identification of 54   embedded clusters around 217 massive   protostellar candidates   of which    34 clusters  are   new detections.  The embedded  clusters are identified as stellar  surface density  enhancements  in the 2  $\\mu$m  All Sky  Survey (2MASS) data. Because the  clusters are all associated  with  massive stars in their earliest evolutionary stage,  the clusters should  also be in an early stage of evolution.   Thus  the properties of  these   clusters should reflect  properties associated with their  formation rather than their evolution.  For each cluster, we estimate  the mass, the morphological type, the photometry and  extinction.  The clusters  in our  study, by their  association with  massive   protostars  and massive   outflows, reinstate the notion that massive stars begin to  form after the first generation   of  low  mass   stars    have completed  their  accretion phase. Further, the observed high gas densities and accretion rates at the centers  of these clusters  is consistent with the hypothesis that high mass stars form by continuing accretion onto low mass stars. ",
        "introduction": "Embedded stellar clusters, those clusters that are still surrounded by the  molecular clouds in  which they  formed, are  the youngest of the stellar clusters. As such,   the embedded clusters are of   particular interest    to understand which     properties of stellar clusters are related  to  their  origins  and which   are derived  from  subsequent evolution  (Elmegreen      et     al.~\\cite{elmegreen00}; Lada     \\& Lada~\\cite{ll03}).  For   example,    the  mass    segregation,   the concentration of higher mass stars in the centers of clusters, that is observed in many optically visible open  clusters is also seen in some of the embedded clusters.  Because the embedded clusters are too young to     have  undergone  significant   dynamical   evolution,  the mass segregation must be a property of the process of star formation in the clusters.  As another example, open clusters exhibit both hierarchical and centrally   condensed morphological  types. Observations   of both morphological types in  embedded clusters suggests that the morphology of the clusters reflects  the morphology of  the clouds from which the stars   formed  rather    than    the dynamical  evolution     of  the cluster.  Finally, the distribution of  stellar masses in the embedded clusters  ought to be   little  affected by  evolution and   therefore closest to the initial mass function (IMF).  While the embedded clusters are thought to be among the youngest clusters, if we were to identify a class of clusters in which star formation, and therefore the formation of the cluster itself, were just beginning, we would potentially be able to address some of the questions as to the causes and origins of some of the cluster properties. For example, observations of embedded clusters may suggest star formation rather than dynamical evolution as a cause of mass segregation, but there remains the question of the cause.  Is mass segregation a result of the formation of massive stars by the collisions and coalescence of lower mass stars because collisions will be more common in the high stellar density in the center of a cluster (Bonnell et al \\cite{bonnell98}; Testi et al. \\cite{testi99})? Is mass segregation a result of the formation of massive stars by continuing accretion onto existing low mass stars (Beech and Mitalas~\\cite{bm94}; Behrend \\& Maeder~\\cite{bm01}; Bernasconi and Maeder~\\cite{bm96}; Meynet and Maeder~\\cite{mm00}, Keto \\cite{keto03}), a process that requires rapid accretion to overcome radiation and thermal  pressure,  and therefore requires dense  gas as would be  found in the center  of a dense molecular clouds? Similarly, observations of actively  forming clusters  might potentially  address the relationship     of   the  star formation  to      molecular cloud structure. What  is the difference in  cloud structure leading  to the hierarchical and  centrally condensed morphological types of clusters? Do  stars always form  first in the center of  a molecular cloud or in gravitationally collapsing fragments throughout the cloud? Finally, if the lower mass  stars form  first as suggested  by several  studies of open  clusters (Herbig \\cite{herbig62}, Stahler \\cite{stahler85}), and required  by the  theories of     massive star  formation either    by coalescence or continuing accretion, then  observations of clusters in formation may potentially address the origins of the IMF. Observations of the  stellar mass  distribution  in actively  forming  clusters may allow the  opportunity  to  see  mass   distributions that  are  still evolving toward an IMF.  In this study, we test  a hypothesis that a  class of actively forming clusters, the youngest subset  of the young  embedded clusters, may be identified   by searching   for  stellar  clusters   around previously identified   massive protostellar   candidates  in isolated  molecular clouds.  These massive protostellar candidates, massive stars in their earliest  stages  of formation,  have  been  identified  as a class of luminous objects having specific IRAS colors  that are associated with other indicators of massive star formation such as dense gas and dust, water  masers, and    ultra-compact   HII (UCHII)  regions  (Palla  et al.~\\cite{palla91};                 Molinari                        et al.~\\cite{mol96},~\\cite{mol98},~\\cite{mol00};       Sridharan       et al.~\\cite{sri01};  Beuther et   al.\\cite{beu02a})  and associated with isolated  molecular  clouds.    These candidates  are   sources deeply embedded in  their molecular clouds   and therefore probably represent the first massive stars to form within the  clouds. If there were more evolved  massive stars  in  these clouds,  then we  would expect their winds, radiation pressure, and supernovae explosions  to have at least partially cleared  the region  of  gas  and dust  revealing  perhaps a classic  open cluster as  would typically be found  in  a more evolved massive star forming  region.  Thus because the  massive  protostellar candidates are  the  first massive stars  to  form within   their host molecular clouds, any associated clusters should represent clusters in their earliest evolutionary phase.  Our study is similar in its objectives to a previous survey of stellar clusters       around     Herbig     Ae/Be     stars      (Testi    et al.~\\cite{testi97},~\\cite{testi99}).    That study   also   sought  to identify very   young clusters with  active   star formation using the Herbig Ae/Be stars, which are protostars, as indicators of active star formation.  In our study we use the massive protostellar candidates as indicators  of the  earliest stages  of  stellar and therefore cluster formation.  A number  of  independent  near-infrared  (NIR)  observations  of  the regions around several  of the massive  protostellar candidates in our target  list   already show evidence   for  embedded  stellar clusters (references in Table 2).  In contrast to these previous  observations, in this study, we undertake a  systematic search for embedded clusters around all the previously identified high mass protostellar candidates in  the lists of   Molinari  et al.~(\\cite{mol96})  and  Sridharan  et al.~(\\cite{sri01}).  We identify the potential clusters  as star count density enhancements above  the mean background  level (Lada \\& Lada~\\cite{ll95}; Carpenter et al.~\\cite{car00}; Ivanov et al.~\\cite{ivanov02}) using the existing K-band      observations  of   the   2MASS    database  (Kleinmann  et al.~\\cite{2mass94}).  We  report on the  identification of 54 clusters by this technique of which  34  are new  detections. We estimate  some basic  properties   that  can be   derived  from  the   J,H,K  data in 2MASS. Finally we discuss some implications of  newer theories of high mass star formation for the formation and evolution of clusters. ",
        "conclusions": "We conducted a systematic search for clustering around 217 candidate HMPOs chosen from the combined lists of MBCP96 and Sri02. We used the 2MASS GATOR database and the technique of producing stellar surface density contours to detect clusters. We also searched for near-infrared counterparts of the 1.2\\,mm dust continuum peaks associated with all candidate HMPOs.  1) We find 54 embedded clusters associated with 217 candidate HMPOs indicating a 25\\% cluster detection rate. All targets lying in the Galactic mid-plane did not show any clusters, and we attribute this to the insensitivity of the 2MASS data to probe into the large extinctions in the Galactic mid-plane region. The detection rate for targets away from the mid-plane is 60\\%.  2) We estimate the mass of each cluster associated with massive protostellar candidates and find that the embedded cluster mass distribution function is similar to that found in a sample of all embedded clusters withing 2 kpc of the Sun.  3) Approximately equal numbers of clusters associated with massive protostellar candidates are found to have hierarchical as centrally condensed structures.  4) In about half of the detected sample, the cluster peaks and the IRAS/mm peaks coincide very well indicating a positive identification of the massive protostar with the infrared visible embedded clusters. This fraction of clusters also display the highest densities and more circular morphology in the entire sample.  5) One  hypothesis of the formation  of   massive stars by continuing accretion is that the younger stars will not  be associated with radio free-free  emission from HII regions. We  find that the near IR colors of  sources  near the dust   continuum  peaks and  that   do not  show free-free emission  are redder than  those  that are  associated  with ultra-compact HII regions.  6) The data are consistent with the hypothesis that massive stars form by continuing accretion, but the data do not discriminate against the hypothesis that massive stars form by the collisions of lower mass stars."
    },
    "0512/hep-th0512247_arXiv.txt": {
        "abstract": "{ It was recently proposed  that our universe could naturally come to be dominated by 3-branes and 7-branes if the universe is ten-dimensional.  In this paper, we explicitly  demonstrate that gravity can be localized on the intersection of three 7-branes in AdS$_{10}$ to give four-dimensional gravity. We derive the exact relations among the tensions of the branes, and show that they apply independently of the precise  distribution of energy within the  necessarily thickened  branes. We demonstrate this with  several technical  sections  showing a simple formula for  the curvature tensor  of a  diagonal metric with isometries  as well as for the curvature at a gravitational singularity. We also demonstrate a subtlety in applying Stoke's Theorem to this set-up. }  \\preprint{.} ",
        "introduction": "Even though a theory with extra dimensions must allow gravitons to propagate in all directions, four-dimensional gravity can apply even with infinitely large extra dimensions if they are sufficiently warped \\cite{rsii}. There are many ways of understanding this result but from the four-dimensional perspective, the higher-dimensional graviton will appear as a tower or spectrum of four-dimensional graviton fields with different masses, similar to the usual Kaluza-Klein case. The spectrum is gapless and continuous and contains a normalizable zero mode (or almost zero mode) that dominates the gravitational potential.  While this mechanism is simplest for a single codimension-one brane, it has been generalized to higher codimension \\cite{gs,inflargedims,gstwo}. Such setups are appealing since string theory motivates a ten-dimensional spacetime, which would make the visible universe a codimension-six brane.  It was recently pointed out \\cite{relaxing} that a generic ten-dimensional FRW cosmology would be dominated by 3-branes and 7-branes. While a 3-brane in this universe would not generally exhibit 4D gravity, the intersection of three 7-branes might. Each 7-brane is codimension-2 and can localize gravity to itself \\cite{gs}.  Intuitively, then, gravity might be localized to the intersection, as Ref. \\cite{inflargedims} showed for the codimension-1 case. Moreover, we generically expect three 7-branes to intersect over a four-dimensional spacetime surface in ten dimensions.  In this  paper , we focus on the triple 7-brane intersection and show that it can localize four-dimensional gravity. We first construct our solution explicitly and then, using the high degree of symmetry of our construction, demonstrate how to extract the tension relations about the thickened branes  in terms of the known external metric and a few parameters of the interior metric of the brane. In particular, we find the necessary tuning relations and show that for a flat four-dimensional universe we do not require an extra tensionful brane at the intersection. We calculate the graviton potential and show that gravity is localized on the intersection.  In an appendix, we present the explicit construction for the same setup with AdS$_4$ or dS$_4$ on the intersection and calculate the leading cosmological constant (c.c.)-dependent term.  It might seem surprising that we can find exact tension relations for codimension-2 branes and their intersections, given that  the codimension-2 branes should be treated as thick branes, and you would expect the tension relation to depend on the precise form of the metric on the interior. However, we will demonstrate that one can apply Stoke's theorem to relate the AdS curvature to the tension on the boundary. Our calculation in fact generalizes the surprising fact already seen in \\cite{gs,gregory, navarro} that the tension relations depended only on boundary conditions and not on the detailed form of the interior metric of a codimension-2 brane.  There is a subtlety in that we also need to take into account an interior contribution which amounts to an internal surface that depends on only a few boundary condition parameters.  To apply  Stoke's Theorem, we need to account for the curvature at a singularity and we show how to do this in the text. ",
        "conclusions": "In this paper, we have constructed 4D gravity in ten dimensions out of 7-branes, essentially as the intersection between three copies of RSII.  Due to the symmetry of the setup, we can generalize previous methods of relating the brane tension to  the curvature of spacetime outside the branes (e.g. \\cite{gs,navarro}) to extract information about the brane intersections. As usual, there is a volcano potential with an exactly solvable zero mode, as well as a continuum of massive modes.   In the course of our analysis we have derived some interesting features of Einstein's equations. We found that whenever the metric does not depend on a coordinate $x$, the corresponding component of the Ricci tensor $R^x_{\\s x}$ is a total derivative $-\\half \\nabla^2 \\log |g_{xx}|$.  We also found a formula for curvature singularities arising from the origin in polar coordinates.   We note that although we will demonstrate that gravity can be localized on the intersection of 7-branes, the filling fraction of the intersection will not in general be the most likely place for our universe to form if the branes forming it are infinite in extent.  However, it could be competitive if they loop around and form loops or some similar configuration, since such a setup would act like a 3-brane on larger scales.  This requires further study which we leave to further work.  Here we show only that the scenario of Ref. \\cite{relaxing} can consistently include four-dimensional gravity, even when no dimensions are compactified.     \\s  \\s"
    },
    "0512/nucl-th0512034_arXiv.txt": {
        "abstract": " ",
        "introduction": "The outer layers of a neutron star are formed of a solid crust composed of a Coulomb lattice of very neutron rich nuclei immersed in a nearly uniform relativistic electron gas. In the deeper layers corresponding to densities beyond the drip threshold $\\rho_{\\rm drip}\\simeq 4\\times 10^{11}\\, {\\rm g}.{\\rm cm}^{-3}$, nuclei are embedded in a sea of ``free'' neutrons which are expected to become superfluids in mature neutron stars whose temperature has dropped below the critical temperature for the onset of superfluidity (for a review of neutron star crust matter, see \\cite{Hae01} and references therein). Below the crust, at densities above $\\sim 10^{14}\\, {\\rm g}.{\\rm cm}^{-3}$ the nuclei merge into a uniform mixture of nucleons and leptons.  The crust of a neutron star, which represents only about $1\\%$ of the mass of the star and $10\\%$ of its radius, is nevertheless important to understand observational phenomena, such as for instance pulsar glitches. Pulsars are strongly magnetised rotating neutron stars whose period, ranging from milliseconds to a few seconds, very slowly increases due to the loss of energy by electromagnetic and gravitational radiation. However some pulsars have been observed to suddenly spin up. These so called ``glitches'' are characterised by a frequency jump, which varies from $\\delta \\Omega/\\Omega \\sim 10^{-9}$ (Crab) up to $\\delta \\Omega/\\Omega \\sim 10^{-6}$ (Vela), often accompanied by a permanent change in the slow down rate from $|\\delta \\dot \\Omega/\\dot \\Omega| \\sim 10^{-5}-10^{-4}$ (Crab) to  $10^{-3}-10^{-2}$ (Vela). The relaxation following a glitch lasts from days to years. It was suggested by Baym \\textit{et al.} \\cite{Baym69} soon after the discovery of the first pulsars that such very long timescales indicate the presence of superfluid components in the interior of the star. It is now widely accepted that pulsar glitches are due to a sudden tranfer of angular momentum from the faster rotating neutron superfluid to the solid crust. However the detailed mechanism of these glitches remain uncertain. In order to understand such phenomena, it is necessary to investigate the transport properties of the neutron superfluid in the inner crust \\cite{CCHII, ChamelCarter05}.    It has been recently shown that the effects of the nuclear clusters on the superfluid neutrons lead to a renormalisation of the neutron mass due to Bragg scattering \\cite{CCHI}, which gives rise to the entrainment effect according to which the momentum of the neutron superfluid is not aligned with the corresponding velocity \\cite{CCHII}. This effective mass has been calculated in the very deep layers of the crust close to the interface with the liquid core for spherical nuclei \\cite{Chamel05} and for the ``pasta'' phases where the crust is formed by a lattice of slab (``lasagna'') or cylinder (``spaguetti'') shaped nuclei \\cite{CCHI}. The effective mass has been found to be close to the bare neutron mass in the bottom of the crust while at lower densities the effective mass is strongly increased, reaching values as high as $m_\\star \\sim 15 m_n$ at the baryon density $5 \\times 10^{13}$ g.cm$^{-3}$. The purpose of the present work is to extend this calculation to much lower densities near the neutron drip point. ",
        "conclusions": "The band theory has been applied to investigate the transport properties of the neutron superfluid in the outermost layers of neutron star crust, at densities around $4\\times 10^{11}$ g.cm$^{-3}$. We have described the free neutrons with an effective Schr\\\"odinger equation, involving a mean local mass $m_n^\\oplus\\{ \\rr \\}$ and a mean potential $U_n\\{ \\rr \\}$ deduced from the results of Negele\\&Vautherin \\cite{NV73} in the extended Thomas-Fermi approximation with Skyrme nucleon-nucleon interactions. Unlike the calculations carried out by these authors, we have solved the equations beyond the Wigner-Seitz approximation by applying Bloch boundary conditions.  It has been found that the effects of the nuclear clusters are very small when the neutron gas is very dilute, meaning that the Fermi wavelength is much larger than the lattice spacing. A similar case is observed in alkali metals. However at higher densities, the effective neutron mass $m_\\star$ (which is analog to the electron optical effective mass in solid state physics) is significantly increased compared to the ``bare'' mass due to Bragg scattering (as it is also observed for the case of electrons in ordinary metals) and fluctuates with the density as a result of band effects. The results also show significant differences with the choice of the parametrisation of the Skyrme interaction. A more accurate evaluation of the effective neutron mass in the crust therefore requires a fully self-consistent treatment in which both the equilibrium structure of the crust and the transport properties are computed within the same nuclear model. However the main conclusions drawn in this paper are not expected to be altered.  The present calculations performed in the outermost layers of the inner crust combined with those carried out recently in the bottom layers \\cite{CCHI, Chamel05}, suggest that the effective neutron mass $m_\\star/m_n$ could take very large values in the middle layers of the crust where the size of the nuclear clusters is of the same order as the lattice spacing. This raises the question whether in some layers the effective neutron mass actually diverges (in the sense that the mobility vanishes ${\\cal K}=0$) due to the existence of band gaps in the energy spectrum. This question is one of the main issues in the study of any periodic materials such as for instance photonic or phononic crystals which have triggered considerable interests, both experimentally and theoretically. In our calculation, no such gaps were found for the outermost layers of the inner crust however as shown by Economou \\textit{et al.} \\cite{Economou89} in the context of photonic crystals, this does not exclude the possibility of gaps at higher densities (higher energies), in deeper layers. The existence of band gaps would have important implications for the dynamics of the star and for its cooling properties.  Very large values of the effective mass have important consequences for the understanding of pulsar glitches. These sudden increases in the rotational frequency of some pulsars are interpreted as discontinuous transfers of angular momentum between the neutron superfluid and the solid crust. The present results suggest that in the middle layers the neutron superfluid would be strongly coupled to the crust so that such layers would contribute very little to the transfer of angular momentum. In the extreme case of band gaps, the superfluid would be locked to the solid parts so that relative currents could not develop (for the same reasons that electric currents do not flow in ordinary insulators). Large enough effective masses could also trigger a Kelvin-Helmholtz instability \\cite{Andersson04}, which could provide a possible explanation for the origin of pulsar glitches. Further observational as well as theoretical investigations are required in order to sheld light on the nature of pulsar glitches and the link with the physics of ``neutronic'' crystals in neutron star crust.     \\bigskip {\\bf Acknowledgements} \\medskip  The author acknowledges financial support from the Lavoisier program of the French Ministry of Foreign Affairs. The author is also very grateful to professor Haensel for discussions.    \\appendix"
    },
    "0512/astro-ph0512500_arXiv.txt": {
        "abstract": "Our understanding of the physical and chemical structure of pre-stellar cores, the simplest star-forming sites, has significantly improved since the last IAU Symposium on Astrochemistry (South Korea, 1999). Research done over these years has revealed that major molecular species like CO and CS systematically deplete onto dust grains at the interior of pre-stellar cores, while species like N$_2$H$^+$ and NH$_3$ survive in the gas phase and can usually be detected towards the core centers. Such a selective behaviour of molecular species gives rise to a differentiated (onion-like) chemical composition, and manifests itself in molecular maps as a dichotomy between centrally peaked and ring-shaped distributions. From the point of view of star-formation studies, the identification of molecular inhomogeneities in cores helps to resolve past discrepancies between observations made using different tracers, and brings the possibility of self-consistent modelling of the core internal structure. Here I present recent work on determining the physical and chemical structure of two pre-stellar cores, L1498 and L1517B, using observations in a large number of molecules and Monte Carlo radiative transfer analysis. These two cores are typical examples of the pre-stellar core population, and their chemical composition is characterized by the presence of large freeze out holes in most molecular species. In contrast with these chemically processed objects, a new population of chemically young cores has started to emerge. The characteristics of its most extreme representative, L1521E, are briefly reviewed. ",
        "introduction": "Pre-stellar (or starless) cores are the simplest star-forming sites. They are isolated, lie nearby, and form one star (or one binary) at the time; they closely resemble the theorist's ideal of a star forming region. When observed in a molecular tracer like ammonia, a pre-stellar core appears as a centrally concentrated structure containing one or a few solar masses of material and having a typical size of about 0.1 pc (see Figure 1 and Benson \\& Myers 1989 for further global properties of cores).  \\begin{figure} \\begin{center} \\includegraphics{l1512_bm89_crop.ps} \\caption{A representative pre-stellar core, L1512. Left: optical image from the Palomar Sky Survey, where L1512 appears as a patch of obscuration against the background of stars. Right: map of the NH$_3$ emission toward the same region showing L1512 as a centrally concentrated (slightly resolved) object. Figure from Benson \\& Myers (1989). }\\label{fig:fig1} \\end{center} \\end{figure}   Pre-stellar cores are the dominant star-forming sites in nearby molecular clouds like Taurus-Auriga, where stars like our Sun are currently forming in the so-called ``isolated mode'' (e.g., Shu, Adams, \\& Lizano 1987). They are not, however, the dominant star-forming regions of our galaxy, as most stars in the Milky Way have formed in groups (e.g., Adams \\& Myers 2001), and therefore must result from the collapse of more complex gas structures. Still, star formation in isolated pre-stellar cores seems to involve most of the physical elements that we associate with the birth of a typical low-mass star, like gravitational infall, disk formation, and bipolar outflow ejection. All these elements, in fact, were first identified in stars forming in isolated cores, and they can be studied with great detail in these simple environments.  The above reasons of simplicity make pre-stellar cores ideal sites to study the initial conditions of star-formation (e.g., Ward-Thompson et al. 1999). Cores are also the most promising places to test the different competing models of star formation, in particular the (fast) turbulence driven scenario (Mac Low \\& Klessen 2004) and the (slow) magnetically-mediated star formation models (Shu et al. 1987, Mouschovias \\& Ciolek 1999). By determining how pre-stellar cores contract out of the more diffuse ambient cloud and by what mechanism cores lose their gravitational support and start collapsing to form a star, we may be able to observationally distinguish between these two models. ",
        "conclusions": ""
    },
    "0512/astro-ph0512446_arXiv.txt": {
        "abstract": "\\noindent \\rightskip=0pt We investigate the emission-line properties of galaxies with red rest-frame colors (compared to the $g-r$ color bimodality) using spectra from Data Release 4 of the Sloan Digital Sky Survey (SDSS). Emission lines are detected in more than half of the red galaxies. We focus on the relationship between two emission lines commonly used as star formation rate indicators: \\hal\\ and \\oiiw. Since \\oii\\ is the principal proxy for \\hal\\ at $z\\sim1$, the correlation between them is critical for comparison between low-z and high-z galaxy surveys.  We find a strong bimodality in \\oii/\\hal\\ ratio in the SDSS sample, which closely corresponds to the bimodality in rest-frame color. Based on standard line ratio diagnostics, most (nearly all of the) line-emitting red galaxies have line ratios typical of various types of Active Galactic Nuclei (AGN) --- most commonly ``low-ionization nuclear emission-line regions'' (LINERs), a small fraction of ``transition objects'' (TOs) and, more rarely, Seyferts. Only $\\sim6\\%$ of red galaxies have line ratios resembling star-forming galaxies. A straight line in the \\oii-\\hal\\ EW diagram separates LINER-like galaxies from other categories, provides an effective classification tool complementary to standard line ratio diagnostics. Quiescent galaxies with no detectable emission lines and those galaxies with LINER-like line ratios combine to form a single, tight red sequence in color-magnitude-concentration space.  Other than modest differences in the luminosity range they span, these two classes are only distinguished from each other by line strength. We also find that \\oii\\ equivalent widths in LINER- and AGN-like galaxies can be as large as that in star-forming galaxies. Thus, unless objects with AGN/LINER-like line ratios are excluded, \\oii\\ emission cannot be used directly as a proxy for star formation rate; this is a particular issue for red galaxies. Lack of \\oii\\ emission is generally used to indicate lack of star formation when post-starburst galaxies are selected at high redshift. Our results imply, however, that these samples have been cut on AGN properties as well as star formation, and therefore may provide seriously incomplete sets of post-starburst galaxies. Furthermore, post-starburst galaxies identifed in SDSS by requiring minimal \\hal\\ equivalent width generally exhibit weak but nonzero line emission with ratios typical of AGNs; few of them show residual star formation. This suggests that most post-starburst galaxies may harbor AGNs/LINERs. ",
        "introduction": "\\oiiw\\ emission line is a widely used, empirical star formation rate (SFR) indicator, especially at high redshift when \\hal\\ moves out of the optical window \\citep[e.g.,][]{GallagherHB89,Kennicutt92,CowieSH96,CowieHS97,Ellis97,Hammer97,Hogg98,RosaGonzalezTT02,Hippelein03}. Although its luminosity is not directly coupled to the ionizing flux, and it is very sensitive to reddening and metallicity effects, it still can be empirically or theoretically calibrated through comparison with \\hal. A variety of calibrations have been developed \\citep{Kennicutt92,Kennicutt98, Kewley04, Mouhcine05, MoustakasKT05}, enabling SFR estimations from \\oii\\ to a reasonable accuracy for star-forming galaxies.  However, many red, elliptical galaxies at $z\\sim0$ also have significant \\oii\\ and other line emission in their spectra \\citep{Caldwell84, Phillips86}.  Does \\oii\\ also indicate star formation in these galaxies? It has long been realized that star formation is not the only possible source of \\oii\\ emission in galaxies. Active Galactic Nuclei (AGN - especially LINERs), fast shockwaves, post-AGB stars, and cooling flows might also produce \\oii\\ emission. Thus, before we use \\oii\\ as a universal star formation indicator, we need to know how often and to what degree \\oii\\ emission is contaminated by sources other than star formation, especially in red, elliptical galaxies.  The assumption that \\oii\\ measures star formation has been fundamental to many studies. For instance, the lack of \\oii\\ is often used as a criterion for post-starburst galaxy identification. Post-starburst galaxies, as their name suggests, are galaxies that underwent a strong recent star formation epoch but have stopped forming stars. Their spectra can be modeled by a combination of an old stellar population (similar to that of a K giant star or an early-type galaxy) and a young population which is dominated by A stars \\citep{DresslerG83}. Therefore, these galaxies are commonly known as `K+A' or `E+A' galaxies. With time, these galaxies will display an early-type galaxy spectrum after all the A stars die out in 1 Gyr. Thus, the study of these galaxies is important for understanding galaxy evolution, given the current uncertainty in how elliptical galaxies form.  The identification of these post-starburst galaxies requires a total lack of emission lines, to be certain that star formation has truly stopped. Most works on post-starbursts \\citep[e.g.,][]{DresslerG83,Zabludoff96,Poggianti99,Balogh99,TranFI03,YangZZ04, BlakePC04,TranFI04} have used \\oii\\ as the marker, either because it is the only emission line diagnostic available at high redshifts or for the sake of facilitating comparisons between low and high redshifts and among different authors. Recently, several groups \\citep{Goto03, Quintero04, Balogh05} have employed \\hal\\ emission as the marker for star formation, rather than \\oii, for samples of galaxies from the Sloan Digital Sky Survey (SDSS). Usually, a low equivalent width (EW) threshold on \\oii, e.g., 2.5\\AA, is used as the criterion for non-detection of such emission lines, according to the sensitivity in each survey. This reflects an underlying assumption that any emission line would be coming from starforming regions. If this assumption fails to hold---for instance, if the contribution of \\oii\\ from AGN is significant---the post-starburst sample defined by such a method would be incomplete because AGN would be mistakenly discarded as star-forming galaxies. Thus, using \\oii\\ in post-starburst galaxy studies requires a better understanding of the many origins of \\oii\\ in all types of galaxies.  In many studies of AGNs, the opposite question is asked: what fraction of \\oii\\ emission in AGN spectra comes from star formation? A variety of studies of high-ionization AGN (QSOs and Seyferts) have attempted to separate the contributions from AGN and star formation to line emission. It is of special interest to the understanding of the coevolution of the nuclear black holes and the bulges. \\cite{Croom02} found by measuring the Baldwin effects \\citep{Baldwin77} that most \\oii\\ emission in QSOs might be from the host galaxies instead of the AGN. \\cite{Richards03} found that dust-reddened QSOs have stronger \\oii\\ emission than normal QSOs. One of the many possible explanations is that the host galaxies of those QSOs with higher extinctions have higher star formation rates. In contrast, studies by \\cite{Ho05} show that in quasar spectra there is very little \\oii\\ emission beyond that expected from the AGN itself, indicating a supressed star formation efficiency in quasar host galaxies.  Unlike the controversy in Type I QSOs, the picture is a little clearer for Seyfert 2s and Type II QSOs. \\cite{Gu06} found that \\oiii/\\hb\\ and \\oiii/\\oii\\ ratios are lower in those Seyfert 2 galaxies with significant star formation, which indicates that star formation could contribute substantially or even dominate the \\hb\\ and \\oii\\ emission. Consistently, \\cite{KimHI06} found that Type II QSOs exhibit significantly enhanced \\oii/\\oiii\\ ratios relative to Type I QSOs. Combined with their high \\oiii\\ luminosities, the high \\oii/\\oiii\\ ratio is best explained by a high level of star formation. However, the problem is far from settled. Although the \\oii\\ emission in QSO spectra could be a combination of two or many origins, we still have no idea of their proportions. In addition, QSO hosts are rare among galaxies. Narrow-line, low-luminosity AGN and LINERs (which may or may not be AGN) are much more abundant. The dominant origin of \\oii\\ emission among these galaxies is still unknown.  A variety of approaches can help to distinguish emission lines having different origins. One commonly used method, which is relatively reliable, is emission-line diagnostics \\citep[etc.]{BPT, VeilleuxO87, RolaTT97, KauffmannHT03}, as also used by many authors mentioned above in QSO studies. This method is effective in distinguishing the two major origins of emission: star formation versus AGN. We will focus on using this method to investigate the origin of \\oii\\ emission in red galaxies in this paper.  The large SDSS redshift survey provides a perfect sample for such a study. Its wide wavelength coverage covers many emission lines from \\oii\\ to \\hal, from which many line ratio diagnostics can be created. At the same time, moderate resolution still allows reasonable line profile fitting to be conducted, giving good control of errors in the line strength measurement. Finally, its unprecedented huge sample size produces reliable statistics. Previous studies of line ratio diagnostics \\citep[e.g.,][]{RolaTT97} used much smaller datasets with less uniform sampling.  In this paper, we present the discovery of a bimodality in \\oii/\\hal\\ ratio among galaxies. One mode is largely associated with star-forming galaxies. Galaxies in the other mode generally have line ratios similar to LINERs, with \\oii\\ emission produced by a mechanism not associated with ordinary star formation. Narrow-line Seyferts and Transition Objects mostly fall in between the two dominant populations, consistent with the picture that both star formation and AGN (or some other source) might contribute substantially to the emission in these objects. The \\oii/\\hal\\ bimodality we find also has important implications for post-starburst studies.  The paper is structured as follows. In \\S\\ref{sec:data}, we describe the data used. (The details of our measurement methods are descibed in the Appendices.) In \\S\\ref{sec:comp}, we compare \\oii\\ emission-line strength with \\hal\\ for all types of galaxies. We demonstrate that red and blue galaxies show different \\oii-\\hal\\ correlations, which appear to reflect different emission origins. Classification based on this \\oii/\\hal\\ bimodality are made and investigated later in \\S\\ref{sec:agn} with the standard line ratio diagnostics. Possible origins of the emission in red galaxies are discussed. With the conclusions drawn from that, we show in \\S\\ref{sec:implications} how the selection of post-starbursts using \\oii\\ leads to an incomplete sample and explore the nature of emission in post-starbursts. We summarize in \\S\\ref{sec:summary}. The cosmology used is a flat $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology with density parameter $\\Omega_m =0.3$. ",
        "conclusions": "\\label{sec:summary}  We have measured the fluxes and EWs of \\oii, \\hal\\, and several other emission lines for $\\sim 300,000$ galaxies in the SDSS DR4 main galaxy sample after careful subtraction of the stellar continua. From the comparison between \\oii\\ and \\hal, combined with a variety of line ratio diagnostics, we have investigated the origin of line emission in red galaxies and the implications for post-starburst galaxy studies. Our main conclusions are as follows:  \\begin{enumerate} \\item  Galaxies display a bimodality in \\oii/\\hal\\ ratio, which corresponds closely to the bimodality in their rest-frame colors. In blue galaxies, the \\oii/\\hal\\ ratio generally matches the expectation for star-forming HII regions. However, in most red galaxies, the \\oii/\\hal\\ ratio is unusually high and difficult to reconcile with predictions for star-forming HII regions.  \\item About 52\\% of all red galaxies have detectable line emission; 38\\% have detectable \\oiiw\\ emission. More than 29\\% of all red galaxies have line ratios characteristic of LINERs, while less than 17\\% are TOs or Seyferts, and $\\sim6\\%$ have lines dominated by star formation (probably dusty starbursts and edge-on spirals). Further information about the spatial distribution of emission in these galaxies and/or X-ray, UV and radio observation is necessary to firmly identify the origin of the line emission in each case.  \\item LINER-like galaxies make up the high-\\oii/\\hal\\ mode in the \\oii/\\hal\\ bimodality. Other than modest differences in the luminosity range spanned, they are essentially indistinguishable from quiescent galaxies in the color-magnitude-concentration space. The combination of LINER-like and quiescent galaxies defines a uniform red sequence in color-magnitude-concentration space, which can be effectively selected by a simple division in the \\oii-\\hal\\ EW diagram. The remaining red galaxies, those with low \\oii/\\hal\\ ratio, are mostly TOs, dusty-starforming galaxies, and a small fraction of Seyferts. This group also have lower galactic concentration and slightly bluer color than the LINER/Quiescent red sequence.  \\item Post-starburst galaxies, identified in the SDSS dataset using a lack of \\hal\\ emission to indicate that star formation has ceased, often exhibit significant \\oii\\ EW. Their positions in the \\oii-\\hal\\ plot and BPT diagrams are the same as for Seyferts, LINERs and TOs, which suggest they may harbor AGNs. Less than 5\\% of this post-starburst sample show evidence for residual star formation.  \\item \\oii\\ emission in red galaxies and in post-starburst galaxies (i.e., in cases where it is not induced by star-formation) can be as strong as in star-forming galaxies. Locally at least, \\oii\\ can only be used as a star-formation indicator for blue galaxies.  \\item More than half of all the post-starbursts have detectable \\oii\\ emission. Post-starburst samples defined using \\oii\\ as a star-formation indicator will be very incomplete, especially if AGN/LINER rates and intensity were higher in the past. We recommend using \\hb\\ as an alternative for high-redshift studies.  \\end{enumerate}"
    },
    "0512/astro-ph0512393_arXiv.txt": {
        "abstract": "{The ``central engine'' of AGN is thought to be powered by accretion on a central nucleus believed to be a super-massive black hole. The localization and exact mechanism of the energy release in AGN are still not well understood.} {We present observational evidence for the link between variability of the radio emission of the compact jet, optical and X-ray continua emission and ejections of new jet components in the radio galaxy 3C\\,390.3.} {The time delays between the light curves of the individual jet components and the light curve of the optical continuum are estimated by using minimization methods and the discret correlation function.} {We find that the variations of the optical continuum are correlated with radio emission from a stationary feature in the jet. This correlation indicates that the source of variable non-thermal continuum radiation is located in the innermost part of the relativistic jet. } {We suggest that the continuum emission from the jet and counterjet ionizes material in a subrelativistic outflow surrounding the jet, which results in a formation of two conical regions with broad emission lines (in addition to the conventional broad line region around the central nucleus) at a distance $\\ga$ 0.4 parsecs from the central engine. Implications for modeling of the broad-line regions are discussed.} ",
        "introduction": "The variable continuum flux in AGN, signaling the activity of the central engine, is detected throughout the entire electromagnetic spectrum, on time-scales from days to years \\citep{peterson02,zheng,wamsteker,shapo}. The continuum flux is believed to be responsible for ionizing the cloud material in the broad-line region (BLR). Localization of the source of the variable continuum emission in AGN is therefore instrumental for understanding the mechanism for release and transport of energy in active galaxies. In radio-quiet AGN, representing about 90\\,\\% of the AGN population, the presence of rapid X-ray flux variations and iron emission line (Fe K$\\alpha$) indicates that most of the soft X-ray emission originates from the accretion disk \\citep{mushotzky93}. In radio-loud AGN, the activity of the central engine is accompanied by highly-relativistic collimated outflows (jets) of plasma material formed and accelerated in the vicinity of the black hole \\citep{ferrari98}. Inhomogeneities in the jet plasma appear as a series of compact radio knots (jet components) observed on scales ranging from several light weeks to about a kiloparsec \\citep{alef,kellermann04}. Continuum emission from the relativistic jet dominates at all energies \\citep[see][] {ulrich,worrall}, swamping the X-ray emission associated with the accretion flow.  Hence, the continuum variability in radio-loud AGN may be related to both the jet and the instabilities of accretion flows \\citep{mushotzky93,ulrich} near the central engine.  The unification scheme \\citep{urry95} of radio-loud AGN suggests that the central powerful optical continuum and broad emission lines are viewed directly in radio quasars and BL Lacs, whereas in radio galaxies these can be hidden by an obscuring, dusty torus, and therefore, the bulk of continuum and broad-line emission in radio galaxies may be attributed to the relativistic jet rather than the central engine. The presence of a positive correlation between beamed synchrotron emission from the base of the jet and optical nuclear emission in radio galaxies suggests that the optical emission is non-thermal and may originate from a relativistic jet \\citep{chiaberge99,chiaberge02,hardcastle00}.  The detection of a correlation between the radio and optical emission variability from the nuclear region would be the most direct evidence of optical continuum emission coming from the jet. No observational evidence has yet been reported for a link between optical/UV continuum variability and the radio jet in radio-loud AGN. In Section 2, we analyse the structure and kinematics of the pc-scale jet in \\object{3C 390.3} and look for correlations between variable emission of jet components and nuclear optical emission on scales less than one parsec. The structure and emission mechanism of the sub-parsec-scale region around the central nucleus is discussed in Section 3. In Section 4, we discuss the results and draw conclusions.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{f1.eps}} \\caption{Radio structure of 3C 390.3 observed in 1998.21 with very long baseline interferometry at 15\\,GHz (2\\,cm). Innermost fraction of the jet is shown in the inset.  The resolving point-spread function (beam) plotted in the upper left corner is $0.87 \\times 0.55$ mas oriented at an angle of 8.0 degrees (clockwise rotation).  The peak flux density in the image is 190\\,mJy/beam ($3.1\\times 10^9$\\,K) and the rms noise is 0.2\\,mJy/beam ($1\\,\\mathrm{Jy}=10^{-26}\\,\\mathrm{W}\\,\\mathrm{m}^{-2}\\,\\mathrm{Hz}^{-1}$). The contours are drawn at $1,\\,\\sqrt{2},\\,2\\,...$ of the lowest contour shown at 0.6\\,mJy/beam. The structure observed is quantified by a set of two-dimensional, circular Gaussian features (shaded circles) obtained from fitting the visibility amplitudes and phases \\citep{pearson}. Similar fits have been obtained for ten observations from the 15\\,GHz VLBA survey database \\citep{kellermann04}, for the purpose of cross-identifying and tracing different features in the jet. The labels mark three stationary features (D, S1, and S2) and a subset of moving components (C2--C6) identified in the jet. Note that two more components, C7 and C8, have been first identified in the jet in the VLBA images at later epochs (see Fig.~\\ref{rfit}).} \\label{rmap} \\end{figure} ",
        "conclusions": "The large distance of the BLR\\,2 from the central engine challenges the existing models in which the broad-line emission is localized exclusively around the disk or near the central engine \\citep{peterson02}. It also questions the assumption of virialized motion in the BLR \\citep{kaspi00}, which forms the foundation of the method for estimating black hole masses from reverberation mapping \\citep{peterson02}. Time delays and profile widths measured during periods when the jet emission is dominant may not necessarily reflect the Keplerian motion in the disk, but rather trace the rotation and outward motion in an outflow. This can result in large errors in estimates of black hole masses made from monitoring of the broad-line emission.  In the case of 3C\\,390.3, the black hole mass ($2.1\\times10^9$ solar masses, $M_{\\odot}$) estimated effectively from the measurements near the maximum in the continuum light curve \\citep{shapo} is significantly larger than the values ($3.5$--$4 \\times 10^8\\,M_{\\odot}$) reported in other works \\citep{wandel99,kaspi00}. This difference is reconciled by considering the line width and the time delay between the optical continuum and line fluxes near the minimum of the continuum light curve, which yields $M_\\mathrm{bh} = 3.8\\times 10^8\\,M_{\\odot}$. The possible existence of an outflow-like region in a number of radio-loud AGN should be taken into account when estimates of the nuclear mass are made from the variability of broad emission lines.  While the continuum emission at the base of the counterjet is likely to be too weak to be detected because of the relativistic dimming and large opacity in the disk, the optical/UV/X-ray line emissions from BLR\\,2 in the counterjet are free from relativistic effects and have better chances to penetrate the absorbing medium towards the observer. Detection of a time lag between correlated variability of emission lines from the far-side and near-side BLRs would be the most direct observational evidence for double BLRs ionized by the continuum radiation from the bases of the jet and counterjet.  If the jets are intrinsically symmetric then a time delay of $\\sim 2$ years is expected for such variations in 3C\\,390.3 (assuming the distance of $\\sim 1$ pc between oppositely positioned BLRs and the jet viewing angle of 50$^{\\circ}$).  The presence of the BLR2 in radio-loud AGN is capable of explaining some of the spectral characteristics of emission lines. Depending on the orientation of the jet, the approaching and rotating outflow material in the BLR2 will imprint prominent signatures on the emission lines. At small viewing angles of the jet the BLR2 may produce blueshifted and single-peaked broad emission lines, while non-shifted and double-peak emission lines \\citep{eracleous03} will be observed at large angles of the BLR2 to the line of sight. Similar characteristics will have the narrow emission lines \\citep{boroson05} ionized in the approaching subrelativistic outflow by the beamed continuum emission of the jet.  The principal results of this work and its implications are: \\begin{enumerate} \\item Analysis of combined radio VLBI, optical/UV and X-ray data reveals significant correlations between variable optical continuum flux (5100\\AA) and radio flux density (15\\,GHz) of D and S1 stationary components of the jet (at the 99\\% and 95\\% confidence levels respectively). The optical emission follows the radio flares with the mean $\\tau_{\\rm D-opt}\\sim1.4$ yr and $\\tau_{\\rm S1-opt}\\sim0.4$ yr. This finding indicates for the physical link between the jet and optical continuum: the variable optical continuum emission is located in the innermost part of the jet near to component S1 and it is of non-thermal origin. This link is also supported by the correlation between the local maxima in the optical continuum light curve and the epochs at which the moving components of the jet pass the stationary radio feature S1.  \\item These results have important implications for the structure of the sub-parsec-scale nuclear region. We suggest that (i) the characteristics of the long-term variability of optical continuum emission are related to the properties of the sub-parsec scale jet (ejection rate of the radio components, its structure and kinematics), (ii) the bulk of optical variable emission originates in the jet at a distance more than or equal to 0.4\\,pc from the central engine, and (iii) the continuum radiation of the jet forms two BLRs (associated with the jet and counterjet) in the subrelativistic outflow around the jet. \\end{enumerate}"
    },
    "0512/astro-ph0512387_arXiv.txt": {
        "abstract": "Multi-wavelength observations of Galactic black hole transients during outburst decay are instrumental for our understanding of the accretion geometry and the formation of outflows around black hole systems. \\SO, a black hole transient observed intensely in X-rays and  also covered in the radio band during its 2003 decay, provides clues about the changes in accretion geometry during state transitions and also the general properties of X-ray emission during the intermediate and the low-hard states. In this work, we report on the evolution of spectral and temporal properties in X-rays and the flux in the radio band with the goal of understanding the nature of state transitions observed in this source. We concentrate on the transition from the thermal dominant state to the intermediate state that occurs on a timescale of one day. We show that the state transition is associated with a sudden increase in power-law flux. We determine that the ratio of the power-law flux to the overall flux in the 3--25 keV band must exceed 0.6 to observe strong timing noise. Even after the state transition, once this ratio was below 0.6, the system transited back to the thermal dominant state for a day. We show that the emission from the compact radio core does not turn on during the transition from the thermal dominant state to the intermediate state but does turn on when the source reaches the low-hard state, as seen in 4U~1543$-$47 and GX~339$-$4. We find that the photon index correlates strongly with the QPO frequency and anti-correlates with the rms amplitude of variability. We also show that the variability is more likely to be associated with the power-law emission than the disk emission. ",
        "introduction": "\\label{sec:intro}  The Galactic black hole transients  show several correlated spectral and temporal variability properties during outbursts, denoted as spectral states. During the initial rise and at the end of the decay before quiescence, these transients are usually in the ``low-hard'' state (LHS). In this state, a hard power-law component dominates the X-ray spectrum, and strong variability ($>$ 20\\% rms amplitude) and quasi-periodic oscillations (QPOs) are often observed. In between the rise and the decay, the source may evolve through a combination of ``thermal dominant'' and ``steep power-law'' states. In the thermal dominant state (TDS), the soft disk component dominates the spectrum, and the timing noise is very low or absent. In the steep power law state, the power-law flux in the 2--20 keV band accounts for more than 50\\% of the flux and has a photon index (\\IN) greater than 2.4. Moderate variability and QPOs are observed in this state. There also exist intermediate states (IS), where source characteristics do not fit into the steep power law, TDS, or LHS, but show various combinations of these states \\citep[see][for detailed discussion of spectral states]{McClintock03}. Throughout this work, we will use the name IS for the particular intermediate state between the TDS and the LHS during the outburst decay. Even though these states were historically characterized using X-ray observations, changes in other bands occur as well. In the TDS, the radio emission from the compact core is quenched \\citep{Fender99,Corbel00}. Optically thin outflows are sometimes detected during state transitions \\citep{Fender01_c,Corbel01}, and powerful, compact jets are always observed in the LHS \\citep{Fender01b}. The optical and infrared emission also show state dependent properties \\citep{Kalemci05,Homan05_a,Corbel02}.  The multi-wavelength observations made during the decaying portion of the outbursts provide valuable information about black hole transients because of the very high probability of observing transitions from the TDS to the IS, and eventually to the LHS \\citep{Kalemci_tez}. The changes during the transitions can reveal the geometry and the physical environment of these systems before and after the transitions \\citep{Esin97,Zdziarski02_2}. The LHS contains additional information due to strong variability, and strong radio emission, both correlating with spectral parameters. Our group has been observing these transients during outburst decay in X-rays with the \\emph{Rossi X-ray Timing Explorer} (\\rxte) and in radio to understand the evolution before, during and after the state transitions \\citep{Kalemci01,Tomsick01b,Kalemci02,Tomsick03_2,Kalemci05}. Our emphasis is on state transitions, and especially on understanding the changes in X-rays while the radio jet is turning on. A uniform analysis of all black hole transients observed with \\wsim daily coverage with \\rxte\\ during outburst decay between 1996 and 2001 provided important information on the evolution of spectral and temporal parameters during the decay \\citep{Kalemci03}. The sharpest change indicating a state transition is observed to be a jump in the rms amplitude of variability from less than a few percent to more than tens of percent in less than a day. This change in the rms amplitude is almost always accompanied by a sharp increase in the power-law flux. There is also evidence that the strong rms noise is only observed when the power-law flux from the source is above a certain percentage of the total flux. This sharp change in rms amplitude of variability is noted as the time of state transition from the TDS to a harder state in \\cite{Kalemci03}, and the same definition will be applied here. During the outburst decay, the photon index, the disk temperature, and the disk flux usually decrease slowly. Often, late in the outburst, the disk flux becomes undetectable. After the transition, characteristic frequencies of the power density spectrum also decrease with time.   \\SO\\ was discovered with \\emph{Ariel 5} \\citep{Kaluzienski77_iauc1} and \\emph{HEAO-1} \\citep{Doxsey77_iauc} satellites in August 1977. After a couple of detections in 1984 with \\emph{EXOSAT} \\citep{Reynolds99}, and in 1996 with TTM/COMIS on \\emph{Mir-Kvant} \\citep{Emelyanov00}, the source was detected in outburst again in March 2003 with \\integral\\ \\citep{Revnivtsev03_atel}, and \\rxte\\ \\citep{Markwardt03_atel}. The radio \\citep{Rupen03_atel}, infrared \\citep{Baba03_iauc} and optical \\citep{Steeghs03_atel} counterparts were quickly identified during the 2003 outburst. The radio observations revealed relativistic jets \\citep[$v/c \\eqsim 0.8$,][]{Rupen04, Corbel05}. Large scale jets were also detected in X-rays with the \\chandra\\ at the end of the outburst \\citep{Corbel05}. The X-ray observations with \\rxte\\ and \\integral\\ indicate that the source went through several spectral states before fading at the end of 2003 \\citep{Markwardt03_atel, Homan03_atel, Kretschmar03_atel, Grebenev03_atel, Tomsick03_atel, Parmar03, Joinet05}. Even though there is no mass measurement of the compact object, the X-ray spectral and temporal properties, and a high frequency QPO pair with frequencies similar to those of other black hole sources \\citep{Homan05} establish this source as a very likely black hole.   In this work, we will characterize the X-ray and radio properties of \\SO\\ during the outburst decay in 2003, compare these properties to the general properties of black hole transients, and discuss the unique properties of this source in detail. We will especially concentrate on the triggering mechanism for the state transitions during the outburst decay. ",
        "conclusions": "We analyzed the \\rxte\\ X-ray observations of \\SO\\ during its outburst decay in 2003. In addition to the X-ray observations, we also obtained radio fluxes and discussed their evolution with respect to the state transitions. The evolution of the spectral and temporal fit parameters show similar properties to those of other sources. The transition from the TDS to the IS is marked by a strong increase in the rms amplitude of variability and power-law flux. This transition may correspond to the emergence of a secondary soft component, such as synchrotron radiation, or a secondary power-law component. Future multi-wavelength observations of these transients may help making the distinction. At the time of this transition the radio core was not detected. Three days after the transition, the source went back to the TDS for one day, indicating that the IS is not stable. After the source returned to the IS, it gradually reached the low hard state. The core is detected in the radio band when the X-ray spectrum is totally dominated by the power-law emission.  There is a strong correlation between the photon index and the QPO frequency. At high QPO frequencies and photon indices, the correlation shows a turnover (or saturates). The threshold PLR for the state transition, and the correlation between the PLR and the ratio of rms amplitude of variability in  6--15 keV and 3--6 keV bands show that the DBB emission dilutes the amplitude of variability. When the DBB emission is absent, the rms amplitude of variability in two bands are similar. This may indicate that the origin of the variability and the QPO is the corona itself, rather than the disk."
    },
    "0512/astro-ph0512178_arXiv.txt": {
        "abstract": "We investigate the time evolution of luminous accretion disks around black holes, conducting the two-dimensional radiation-hydrodynamic simulations. We adopt the $\\alpha$ prescription for the viscosity. The radial-azimuthal component of viscous stress tensor is assumed to be proportional to the total pressure in the optically thick region, while the gas pressure in the optically thin regime. The viscosity parameter, $\\alpha$, is taken to be 0.1. We find the limit-cycle variation in luminosity between high and low states. When we set the mass input rate from the outer disk boundary to be $100 L_{\\rm E}/c^2$, the luminosity suddenly rises from $0.3L_{\\rm E}$ to $2L_{\\rm E}$, where $L_{\\rm E}$ is the Eddington luminosity. It decays after retaining high value for about $40$ s. Our numerical results can explain the variation amplitude and duration of the recurrent outbursts observed in microquasar, GRS 1915+105. We show that the multi-dimensional effects play an important role in the high-luminosity state. In this state, the outflow is driven by the strong radiation force, and some part of radiation energy dissipated inside the disk is swallowed by the black hole due to the photon-trapping effects. This trapped luminosity is comparable to the disk luminosity. We also calculate two more cases: one with a much larger accretion rate than the critical value for the instability and the other with the viscous stress tensor being proportional to the gas pressure only even when the radiation pressure is dominant. We find no quasi-periodic light variations in these cases. This confirms that the limit-cycle behavior found in the simulations is caused by the disk instability. ",
        "introduction": "Microquasars in our Galaxy display large flux variability in X-ray band. Such luminosity variations are thought to reflect violent phenomena in accretion disks around black holes. The dwarf-nova type disk-instability is responsible for the luminosity changes on the time-scale of months (e.g., Mineshige \\& Wheeler 1989; for a review Kato et al. 1998, \\S 5). In contrast, the mechanism causing the short-term variability ($10-100$ s) is not well understood yet. One of the plausible mechanisms for it is thermal and secular instability which arises when the radiation pressure becomes dominant over gas pressure (e.g., Lightman \\& Eardley 1974; Shibazaki \\& H\\=oshi 1975; Pringle 1976; Shakura \\& Sunyaev 1976). S-shaped sequence on the $\\dot{M}_{\\rm acc}-\\Sigma$ (mass accretion rate vs. surface density) plane was completed by the finding of a slim-disk branch (Abramowicz et al. 1998), in which the viscous heating is balanced by the advective cooling. It was suggested that the disk may undergo limit-cycle oscillations, like the dwarf nova outbursts (for a review Kato et al. 1998, \\S 10). This instability is expected to occur when the mass accretion rate is comparable to or moderately exceeds the critical value, $L_{\\rm E}/c^2$, where $L_{\\rm E}$ is the Eddington luminosity and $c$ is the velocity of light.    The limit-cycle oscillations have been investigated by one-dimensional (1D), vertically-integrated approach. The quasi-periodic outbursts were first demonstrated by Honma et al. (1991) using the time-dependent 1D simulations of the accretion disks and have then been investigated in detail (Szuszkiewicz \\& Miller 1997, 1998, 2001; Watarai \\& Mineshige 2003). However, the 1D model cannot treat the multi-dimensional motion, i.e., convection, circulation, and outflow, although they would influence the flow dynamics significantly via the transport of mass, momentum/angular momentum, and energy. The modification of the accretion disk model, where the mass ejection from the disk surface is taken into account, was suggested by Nayakshin et al. (2000) (see also Janiuk et al. 2000; 2002 and Janiuk \\& Czerny 2005). But they still used simplified 1D method without treating the multi-dimensional flow motion. We need to perform at least two-dimensional (2D) analysis in order to understand the accretion flow correctly.  Recently, Teresi et al. (2004a, 2004b) first reproduced the quasi-periodic luminosity variations by a 2D smoothed particle hydrodynamic (SPH) simulations. However, their simulations focused on the time evolution of the disk itself and do not treat the optically thin regime correctly, i.e., the disk surface as well as the atmosphere surrounding the disk. They assumed the equilibrium between gas and radiation without treating the energy of gas and radiation separately. In their simulations, the transport of the radiation energy was not solved in the optically thin regime, though the radiation flux was evaluated by using the diffusion approximation deep inside the disk. The thermal energy was extracted from the SPH particles located near the disk surface. It should be calculated by solving the interaction between gas and radiation as well as radiative transfer. Their assumptions and method are valid only in the optically thick regime, i.e., deep inside the disk. As a result, the outflow was not found in their simulations even though the luminosity exceeded the Eddington luminosity in the high-luminosity state. The outflow would have an important role on the evolution of the disk through the extraction of the mass, momentum/angular momentum, and energy from the disk. In addition, the simulations for stable disks were not performed in their work. A comparison of their results with those of simulations for stable disks is essential in order to understand the physical mechanism of the disk oscillations. Hence, it is important to confirm the limit-cycle behavior by the grid-based simulations, in which the multi-dimensional radiation hydrodynamic (RHD) equations are properly solved. It is also important to investigate the time evolution of not only an unstable disk but also a stable disk.   The photon trapping is also basically a multi-dimensional effect, by which some or large part of photons generated inside the disk is swallowed by the black hole in supercritical accretion regime, leading to a reduction of the energy conversion efficiency (Ohsuga et al. 2002, 2003, 2005). It has been reported by the analysis of the X-ray data of GRS 1915+105 that the disk luminosity is comparable to or exceeds the Eddington luminosity in the high-luminosity state (Yamaoka et al. 2001), implying the supercritical accretion. Thus, the photon trapping is expected to appear in high-luminosity state of the quasi-periodic oscillations.   Here, by solving full set of 2D RHD equations including viscosity term, we report the 2D RHD model for the quasi-periodic oscillations of the accretion disks around black holes. Our numerical simulations carefully treat the 2D effects, including the outflow motion and the photon trapping. We also show in our simulations that the disk is unstable on condition that the mass accretion rate is moderately larger than the critical value and the viscous stress tensor is proportional to the total pressure. It is consistent with the disk theory and proves that the bursting behavior in our simulations arises from the disk instability in the radiation-pressure dominant region. In \\S 2, our model and numerical method are described. We present the numerical results in \\S 3. \\S 4 and \\S5 are devoted to discussion and conclusions. ",
        "conclusions": "By performing the 2D RHD simulations, we investigate the time evolution of the accretion disks around the black holes and find the limit-cycle oscillations. In this study, we carefully treat the 2D effects including the outflow and the photon trapping. We summarize our results as follows:  (1) The luminosity sharply rises from $\\sim 0.3L_{\\rm E}$ to $\\sim 2L_{\\rm E}$ in our simulations. Duration of the high-luminosity state is about $30-50$ s. The resulting variation amplitude and duration nicely fit to the observations of microquasar, GRS 1915+105.  (2) It is also found that the 2D effects are significant in the high-luminosity state. In this state, the trapped luminosity is comparable to the luminosity due to the effective photon trapping. The outflow, circular motion, and patchy structure appear in this state. The outflow is driven by the strong radiation force.   (3) The physical mechanism, which causes the limit-cycle oscillations, is the thermal instability in the radiation-pressure dominant region. In our simulations, the disk is stabilized when the mass accretion rate highly exceeds the critical value. The disk does not exhibit the oscillations if the viscous stress tensor is proportional to the gas pressure only."
    },
    "0512/astro-ph0512208_arXiv.txt": {
        "abstract": "We study the effects of galaxy formation on the Sunyaev-Zel'dovich effect (SZE) observable-mass relations using high-resolution cosmological simulations. The simulations of eleven individual clusters spanning a decade in mass are performed with the shock-capturing Eulerian adaptive mesh refinement N-body+gasdynamics ART code.  To assess the impact of galaxy formation, we compare two sets of simulations performed in an adiabatic regime (without galaxy formation) and those with several physical processes critical to various aspects of galaxy formation: radiative cooling, star formation, stellar feedback and metal enrichment.  We show that a SZE signal integrated to a sufficiently large fraction of cluster volume correlates strongly with its enclosed mass, independent of details of gas physics and dynamical state of the cluster.  The slope and redshift evolution of the SZE flux-mass relation are also insensitive to processes of galaxy formation and are well characterized by a simple self-similar cluster model.  Its normalization, on the other hand, is significantly affected by gas cooling and associated star formation.  Our simulations show that inclusion of these processes suppresses the normalization by $\\approx 30-40\\%$.  The effect is due to a decrease in gas mass fraction, which is offset slightly by an increase in gas mass-weighted temperature. Gas cooling and star formation also cause an increase in total mass and modify the normalization by a few percent.  Finally, we compare the results of our simulations to recent observations of the SZE scaling relations obtained using 36 OVRO/BIMA SZE$+${\\it Chandra} X-ray observations. The comparison highlights the importance of galaxy formation in theoretical modelling of clusters and shows that current generation of simulations produce clusters with gross properties quite similar to their observed counterparts. ",
        "introduction": "\\label{sec:intro}  The Sunyaev-Zel'dovich effect (SZE) is a potentially powerful observational tool for cosmology.  It is a small distortion in the cosmic microwave background (CMB) spectrum caused by scattering of CMB photons off a distribution of high energy electrons in dense structures such as clusters of galaxies \\citep{sunyaev_etal70,sunyaev_etal72}.  This effect has the unique property that its signal is independent of redshift, making it particularly well suited for deep cluster surveys \\citep[e.g.,][]{holder_etal00,weller_etal02}. The next generation of SZE instruments, such as the South Pole Telescope (SPT) and the Atacama Cosmology Telescope (ACT), should be capable of mapping a fairly large portion of the sky and finding a large number ($\\gtrsim 10^4$) of clusters out to high-redshift.  Such large and homogeneous sample of galaxy clusters will enable direct and precise measurements of their number density as a function of redshift, and the expected survey yields will be sufficient to provide one of the most powerful constraints on the nature of dark energy \\citep{wang_etal98,haiman_etal01}.  To realize the full statistical power of the upcoming SZE surveys, however, systematic uncertainties would have to be controlled at a level comparable to statistical uncertainties.  One of the main sources of the systematic uncertainties lie in the relation between the SZE observable and cluster mass as a function of redshift.  Making this connection is important because the cluster mass is not directly observable.  For the future SZE surveys, the requirement is to control systematic uncertainties in the SZE observable-mass relations to better than $\\sim$5\\% at all redshift \\citep[e.g.,][]{carlstrom_etal02}.  This poses a serious challenge to both observers and theorists. To date, observational studies of the SZE scaling relations have been performed using two largest datasets of SZE measurements obtained by the OVRO/BIMA cm-wave imaging experiment \\citep{cooray_etal99,mccarthy_etal03,laroque05} and multiband SuZIE experiment \\citep{benson_etal04} along with X-ray observations. Analyzing a sample of 36 clusters observed with both the OVRO/BIMA SZE imaging and {\\it Chandra} X-ray observations, \\citet{laroque05} showed that there are tight correlations between the observed SZE flux and X-ray temperature and cluster masses.  The observed regularity of the SZE effect of clusters is encouraging news, but further progress is clearly needed for the future SZE surveys.  The observational situation is expected to improve rapidly with the advent of a number of dedicated SZE survey instruments, which will dramatically increase the sample size and the number of low-mass clusters.  On a theoretical side, a number of groups have studied the SZE scaling relations using semi-analytic models \\citep{verde_etal02,mccarthy_etal03} and cosmological simulations \\citep{metzler_etal98,white_etal02,daSilva_etal04,diaferio_etal05,motl_etal05}. Motivated by results of X-ray observations in a past decade \\citep[see][for a review and references therein]{voit_etal05}, recent studies have focused on studying the effects of non-gravitational physical processes, including gas cooling, star formation and energy feedback, on the SZE scaling relations.  One of the main results is that a SZE signal integrated to a sufficiently large fraction of cluster volume is an extremely good proxy for its enclosed mass, independent of details of gas physics and dynamical state of a cluster \\citep[see e.g.,][]{motl_etal05}.  The slope and the redshift evolution of the SZE scaling relation also appear to be insensitive to details of cluster gas physics \\citep[e.g.,][]{daSilva_etal04}.  While these results are encouraging for cosmological applications, these previous studies have focused on simulating a large number of clusters, and the resolution was inevitably limited to capture relevant cluster physics. As such, the impact of galaxy formation on normalization of the SZE scaling relations has not yet converged among different simulations.  It is therefore important to push theoretical modelling of SZE scaling relations and check previous results using higher-resolution cluster simulations.  In this paper, we present such study using high-resolution cosmological simulations of cluster formation.  Although the statistic is limited, our cluster sample spans over a decade in mass and provide a good leverage on scaling relations.  The mass resolution of our simulations is more than an order of magnitude higher than that in previous studies.  This work is therefore complimentary to the previous studies in a literature.  Using these simulations, we study the impact of gas cooling and star formation on the SZE scaling relations, including their normalization, slope and redshift evolution.  To test the results of our simulations, we compare our results to recent observations of the SZE scaling relations based on a sample of 36 clusters obtained using the OVRO+BIMA SZE and {\\it Chandra} X--ray telescopes \\citep{laroque05}.  The paper is organized as follows. In \\S~\\ref{sec:model} we define observational quantities and present relevant scaling laws predicted by a self-similar cluster model.  We describe simulations presented in this paper in \\S~\\ref{sec:sim} and present results and comparisons to previous studies and recent observational results in \\S~\\ref{sec:results}.  Finally, in \\S~\\ref{sec:discussion} we discuss our conclusions and their implications for SZE cluster surveys. ",
        "conclusions": "\\label{sec:discussion}  We have presented the analysis of the Sunyaev-Zel'dovich effect (SZE) scaling relations using high-resolution simulations of galaxy clusters formed in a concordance $\\Lambda$CDM cosmology.  The simulations of eleven individual clusters spanning a decade in mass ($M_{500c}=3.5\\times 10^{13}$ to $9\\times 10^{14}h^{-1}M_{\\odot}$) are performed with the shock-capturing Eulerian adaptive mesh refinement N-body+gasdynamics ART code.  We study the effects of gas cooling and star formation on the SZE scaling relations and their redshift evolution between $z$=0 and $z$=2 by comparing two sets of simulations performed with and without these processes included.  The main results are summarized as follows.  \\newcounter{bean} \\begin{list} {\\arabic{bean}.}{ \\usecounter{bean} \\setlength{\\parsep}{+0.03in} \\setlength{\\leftmargin}{+0.15in} \\setlength{\\rightmargin}{+0.15in}}  \\item{} The SZE signal integrated to a sufficiently large fraction of cluster volume correlates very strongly with the enclosed total cluster mass, independent of details of gas physics and dynamical state of clusters.  The rms scatter of the SZE-total mass relation is about 10-15\\%.  \\item{} The slope of the relation in the adiabatic run is in a very good agreement with the predicted slope of the self-similar model in the entire redshift range.  We find that the impact of galaxy formation on the slope is small ($\\lesssim 0.1$).  \\item{} The redshift evolution of the SZE-total mass relation is consistent with the self-similar model between $z$=0 and $z$=2: (a) the slope is constant with redshift and (b) the normalization evolves with redshift according to the self-similar evolution model.  \\item{} Gas cooling and star formation significantly modify the normalization of the SZE flux-total mass relation.  Inclusion of these physical processes causes a decrease in the normalization by 41, 34 and 27\\% at $r_{2500c}$, $r_{500c}$ and $r_{180m}$, respectively.  The decrease is due to a large decrease in gas fraction, which is offset somewhat by an increase in mass-weighted temperature.  Gas cooling and star formation also cause an increase in total cluster mass and hence modify the normalization by a few percent.  \\item{} The integrated SZE signal also correlates strongly with gas mass and mass-weighted temperature of clusters.  The results (1)-(3) apply equally well for these relations, except that the redshift evolution of the normalization exhibits some deviations from the self-similar model, which increases toward higher redshifts.  Gas cooling and star formation also significantly modify the normalization of these relations.  \\item{} The SZE flux-gas mass relation in the simulations with gas cooling and star formation is in a better good agreement with the observed relation for a sample of 36 OVRO/BIMA SZE+{\\it Chandra} X-ray observations \\citep{laroque05} than the simulations neglecting galaxy formation.  \\end{list}  These results have a number of important implications for cosmological studies with upcoming SZE cluster surveys.  First and foremost, the SZE fluxes of clusters exhibit a remarkable regularity at all redshifts, and the SZE signal integrated to a sufficiently large fraction of cluster volume is insensitive to merging events \\citep[see e.g.,][]{motl_etal05} or properties of a cluster core (see \\S~\\ref{sec:profiles}).  This indicates that the integrated SZE flux is an extremely good proxy for cluster mass.  Second, the slope and the redshift evolution of the SZE scaling relations are insensitive to details of cluster gas physics, and they are well characterized by a simple self-similar cluster model between $z$=0 and 2.  The simplicity of their redshift evolution implies that the self-calibration \\citep{hu03,majumdar_etal04} will be effective.  Finally, these results appear to be very robust, as the same conclusions have been reached using simulations with very different numerical techniques, resolution and implementation of various physical processes incorporated in simulations (see \\S~\\ref{sec:comp}).  Despite the simplicity of redshift evolution, the normalization of the SZE flux-mass relation is much less understood, because the effect is closely related to the cold baryon fraction, which has not yet converged among different simulations.  Using high-resolution cluster simulations, we show that processes of galaxy formation have a significant impact on the normalization of the SZE scaling relations. Gas cooling and star formation suppress the normalization by $\\approx 30-40\\%$, primarily due to a large reduction in cluster gas mass fraction.  Interestingly, the SZE scaling relations in these simulations are in a reasonably good agreement with recent observations.  Moreover, the gas mass fractions in these simulations compare well with measurements from deep {\\it Chandra} X--ray observations of nearby clusters (Kravtsov et al. in preparation). Despite these successes in matching SZE and X--ray observations, the current cluster simulations may still suffer from the \"overcooling\" problem since the fraction of baryons in the cold gas and stars within a virial radius at $z$=0 is in a range ~0.25-0.35, at least a factor of two higher than observational measurements \\citep[see][and discussions and references therein]{kravtsov_etal05}.  Given the importance of these issues, further efforts from observers and theorists are needed to better understand the SZE scaling relations for cosmological studies.  Observationally, it is important to increase the sample size and the number of low-mass clusters. Comparisons of different instruments will also help resolve systematic uncertainties among different instruments \\citep[see e.g.,][]{benson_etal04}.  It is also critical to understand systematic uncertainties in measurements of cluster mass through detailed and extensive comparisons of X--ray, SZE and optical observations.  With the advent of a number of dedicated SZE survey instruments, the observational situation is expected to improve rapidly.  Theoretically, it is important to push detailed theoretical modelling of cluster formation to investigate roles of various physical processes, such as thermal conduction and AGN feedback, in shaping the properties of the ICM.  For cosmological application, it is also critical to understand projection effects \\citep{white_etal02} and nature of scatter in observable-mass relations \\citep{lima_etal05}. Numerical simulations of cluster formation will likely provide important insights into these issues and help assess the effectiveness of the self-calibration technique \\citep{hu03,majumdar_etal04} for future cluster surveys."
    },
    "0512/astro-ph0512514_arXiv.txt": {
        "abstract": " ",
        "introduction": "Galaxies in clusters are sensitive to environmental effects like the cluster tidal field, gravitational encounters with other galaxies, galaxy mergers,  ram pressure stripping and accretion of gas (see e.g. Moore et al. 1998, Vollmer et al. 2001). Such external events dramatically affect their structure, triggering internal  perturbations  like bars or oval distorsions (e.g. Bournaud \\& Combes 2002), spirals,  warps (Huang \\& Carlberg 1997)  or lopsidedness (Bournaud et al. 2005). Their kinematics is also disturbed, as revealed by long-slit spectroscopy 1-D rotation curves (Rubin et al. 1999).  High-resolution \\ha\\ velocity fields were obtained for 30 Virgo cluster galaxies (Chemin et al. 2005) in order to study the  degree of perturbation of their 2-D kinematics and the influence of the environment on the kinematics. The harmonic analysis is a powerful tool to detect kinematical anomalies, as already shown on \\hi\\ velocity fields (Schoenmakers, Franx \\& de Zeeuw 1997). This technique is applied to the \\ha\\ velocity field of NGC 4254 (Figure~\\ref{fig1}). ",
        "conclusions": ""
    },
    "0512/nucl-th0512020_arXiv.txt": {
        "abstract": "In cores of supernovae and crusts of neutron stars, nuclei can adopt interesting shapes, such as rods or slabs, etc., which are referred to as nuclear ``pasta.'' Recently, we have been studying the pasta phases focusing on their dynamical aspects with quantum molecular dynamic (QMD) approach. We review our findings on the following topics: dynamical formation of the pasta phases by cooling down the hot uniform nuclear matter; a phase diagram on the density versus temperature plane; structural transitions between the pasta phases induced by compression and their mechanism. Properties of the nuclear interaction used in our works are also discussed. ",
        "introduction": "In ordinary matter, atomic nuclei are roughly spherical. This may be understood in the liquid drop picture of the nucleus as being a result of the forces due to the surface tension of nuclear matter, which favors a spherical nucleus, being greater than those due to the electrical repulsion between protons, which tends to make the nucleus deform. When the density of matter approaches that of atomic nuclei, i.e., the normal nuclear density $\\rho_0$, nuclei are closely packed and the effect of the electrostatic energy becomes comparable to that of the surface energy. Consequently, at subnuclear densities around $\\rho\\simeq\\rho_0/2$, the energetically favorable configuration is expected to have remarkable structures: the nuclear matter region (i.e., the liquid phase) is divided into periodically arranged parts of rodlike or slablike shape, embedded in the gas phase and in a roughly uniform electron gas. Besides, there can be phases in which nuclei are turned inside out, with cylindrical or spherical bubbles of the gas phase in the liquid phase. These phases with nonspherical nuclei are often referred to as nuclear ``pasta'' phases because nuclear slabs and rods look like ``lasagna'' and ``spaghetti.'' Likewise, spherical nuclei and bubbles are called ``meatballs'' and ``cheese,'' respectively.   In equilibrium dense matter in supernova cores and neutron stars, existence of the pasta phases has been predicted by Ravenhall {\\it et al.} \\cite{rpw} and Hashimoto {\\it et al.} \\cite{hashimoto}. Since these seminal works, properties of the pasta phases in equilibrium states have been investigated with various nuclear models. They include studies on phase diagrams at zero temperature \\cite{lorenz,oyamatsu,sumiyoshi,gentaro,williams} and at finite temperatures \\cite{lassaut}. These earlier works have confirmed that, for various nuclear models, the nuclear shape changes as: $\\mbox{sphere} \\rightarrow \\mbox{cylinder} \\rightarrow \\mbox{slab} \\rightarrow \\mbox{cylindrical hole} \\rightarrow \\mbox{spherical hole} \\rightarrow \\mbox{uniform}$, with increasing density.   In these earlier works, however, a liquid drop model or the Thomas-Fermi approximation is used with an assumption on the nuclear shape (except for Ref.\\ \\cite{williams}). Thus the phase diagram at subnuclear densities and the existence of the pasta phases should be examined without assuming the nuclear shape. It is also noted that at temperatures of several MeV, which are relevant to the collapsing cores, effects of thermal fluctuations on the nucleon distribution are significant. However, these thermal fluctuations cannot be described properly by mean-field theories such as the Thomas-Fermi approximation used in the previous work \\cite{lassaut}.   In contrast to the equilibrium properties, dynamical or non-equilibrium aspects of the pasta phases had not been studied until recently except for some limited cases \\cite{formation,review}. Thus it had been unclear even whether or not the pasta phases can be formed and the transitions between them can be realized during the collapse of stars and the cooling of neutron stars, which have finite time scales.   To solve the above problems, molecular dynamic approaches for nucleon many-body systems are suitable. They treat the motion of the nucleonic degrees of freedom and can describe thermal fluctuations and many-body correlations beyond the mean-field level.   Using the framework of QMD \\cite{aichelin}, which is one of the molecular dynamic methods, we have solved the following two major questions \\cite{qmd_transition,qmd_cold,qmd_hot}. \\begin{itemize} \\item {\\it Question} 1: Whether or not the pasta phases are formed by cooling down hot uniform nuclear matter in a finite time scale much smaller than that of the neutron star cooling? \\item {\\it Question} 2: Whether or not transitions between the pasta phases can occur by the compression during the collapse of a star? \\end{itemize}   The pasta phases have recently begun to attract the attention of many researchers (see, e.g., Refs.\\ \\cite{burrows,martinez} and references therein). The mechanism of the collapse-driven supernova explosion has been a central mystery in astrophysics for almost half a century (e.g., Ref.\\ \\cite{bethe}). Previous studies suggest that the revival of the shock wave by neutrino heating is a crucial process. As has been pointed out in Refs.\\ \\cite{gentaro,qmd_cold} and elaborated in Refs.\\ \\cite{horowitz1,horowitz2,future}, the existence of the pasta phases instead of uniform nuclear matter increases the neutrino opacity of matter in the inner core significantly due to the neutrino coherent scattering by nuclei \\cite{freedman,sato}; this affects the total energy transferred to the shocked matter. Thus the pasta phases could play an important role in the future study of supernova explosions. Our recent work \\cite{qmd_transition} strongly suggests the possibility of dynamical formation of the pasta phases from a crystalline lattice of spherical nuclei; effects of the pasta phases on the supernova explosions should be seriously discussed in the near future. ",
        "conclusions": "We approached the two questions posed in Section I using the framework of QMD. According to the results of our simulations, our answer is strongly affirmative for both questions. The nuclear interaction used in our simulations shows generally reasonable properties at subnuclear densities not only for symmetric nuclear matter but also for neutron matter. This result also supports our conclusion.          \\begin{theacknowledgments} G. W. appreciates C. J. Pethick for his valuable comments and hospitality at NORDITA. The research reported in this article grew out of collaborations with Kei Iida, Toshiki Maruyama, Katsuhiko Sato, Kenji Yasuoka and Toshikazu Ebisuzaki. Further research currently in progress is performed using RIKEN Super Combined Cluster System with MDGRAPE-2. This work was supported in part by the Nishina Memorial Foundation, by the JSPS Postdoctoral Fellowship for Research Abroad, by the Japan Society for the Promotion of Science, by the Ministry of Education, Culture, Sports, Science and Technology through Research Grant No. 14-7939, and by RIKEN through Research Grant No. J130026. \\end{theacknowledgments}"
    },
    "0512/astro-ph0512272_arXiv.txt": {
        "abstract": "We perform numerical simulations of the Kelvin-Helmholtz instability in the mid-plane of a protoplanetary disk. A two-dimensional corotating slice in the azimuthal--vertical plane of the disk is considered where we include the Coriolis force and the radial advection of the Keplerian rotation flow.  Dust grains, treated as individual particles, move under the influence of friction with the gas, while the gas is treated as a compressible fluid. The friction force from the dust grains on the gas leads to a vertical shear in the gas rotation velocity. As the particles settle around the mid-plane due to gravity, the shear increases, and eventually the flow becomes unstable to the Kelvin-Helmholtz instability. The Kelvin-Helmholtz turbulence saturates when the vertical settling of the dust is balanced by the turbulent diffusion away from the mid-plane. The azimuthally averaged state of the self-sustained Kelvin-Helmholtz turbulence is found to have a constant Richardson number in the region around the mid-plane where the dust-to-gas ratio is significant. Nevertheless the dust density has a strong non-axisymmetric component. We identify a powerful clumping mechanism, caused by the dependence of the rotation velocity of the dust grains on the dust-to-gas ratio, as the source of the non-axisymmetry. Our simulations confirm recent findings that the critical Richardson number for Kelvin-Helmholtz instability is around unity or larger, rather than the classical value of 1/4. ",
        "introduction": "One of the great unsolved problems of planet formation is how to form planetesimals, the kilometer-sized precursors of real planets \\citep{Safronov1969}. At this size solid bodies in a protoplanetary disk can attract each other through gravitational two-body encounters, whereas gravity is insignificant between smaller bodies. Starting from micrometer-sized dust grains, the initial growth is caused by the random Brownian motion of the grains \\citep[e.g.][see \\cite{Henning+etal2006} for a review]{BlumWurm2000,DullemondDominik2005}. The vertical component of the gravity from the central object causes the gas in the disk to be stratified with a higher pressure around the mid-plane. Even though the dust grains do not feel this pressure gradient, the strong frictional coupling with the gas prevents small grains from having any significant vertical motion relative to the gas. However, once the grains have coagulated to form pebbles with sizes of a few centimeters, the solids are no longer completely coupled to the gas motion.  They are thus free to fall, or sediment, towards the mid-plane of the disk. The increase in dust density opens a promising way of forming planetesimals by increasing the local dust density around the mid-plane of the disk to values high enough for gravitational fragmentation of the dust layer \\citep{Safronov1969,GoldreichWard1973}.  There are however two major unresolved problems with the gravitational fragmentation scenario. Any global turbulence in the disk causes the dust grains to diffuse away from the mid-plane, and thus the dust density is kept at values that are too low for fragmentation. A turbulent $\\alpha$-value of $10^{-4}$ is generally enough to prevent efficient sedimentation towards the mid-plane \\citep{WeidenschillingCuzzi1993}, whereas the $\\alpha$-value due to magnetorotational turbulence \\citep{BalbusHawley1991,Brandenburg+etal1995,Hawley+etal1995,Armitage1998} is from a few times $10^{-3}$ (found in local box simulations with no imposed magnetic field) to 0.1 and higher (in global disk simulations). The presence of a magnetically dead zone around the disk mid-plane \\citep{Gammie1996,Fromang+etal2002,Semenov+etal2004} may not mean that there is no turbulence in the mid-plane, as other instabilities may set in and produce significant turbulent motion \\citep{Li+etal2001, KlahrBodenheimer2003}. The magnetically active surface layers of the disk can even induce enough turbulent motion in the mid-plane to possibly prevent efficient sedimentation of dust \\citep{FlemingStone2003}. The presence of a dead zone may actually {\\it in itself} be a source of turbulence. The sudden fall of the accretion rate can lead to a pile up of mass in the dead zone, possibly igniting the magnetorotational instability in bursts \\citep{Wuensch+etal2005} or a Rossby wave instability \\citep{VarniereTagger2005}.  The second major problem with the gravitational fragmentation scenario is that even in the absence of global disk turbulence, the dust sedimentation may in itself destabilize the disk. Protoplanetary disks have a radial pressure gradient, because the temperature and the density fall with increasing radial distance from the central object, so the gas rotates at a speed that is slightly below the Keplerian value. The dust grains feel only the gravity and want to rotate purely Keplerian. Close to the equatorial plane of the disk, where the sedimentation of dust has increased the dust-to-gas ratio to unity or higher, the gas is forced by the dust to orbit at a higher speed than far away from the mid-plane where the rotation is still sub-Keplerian. Thus there is a vertical dependence of the gas rotation velocity. Such shear flow can be unstable to the Kelvin-Helmholtz instability (KHI), depending on the stabilizing effect of vertical gravity and density stratification. A necessary criterion for the KHI is that the energy required to lift a fluid parcel of gas and dust vertically upwards by an infinitesimal distance is available in the relative vertical motion between infinitesimally close parcels \\citep{Chandrasekhar1961}. The turbulent motions resulting from the KHI are strong enough to puff up the dust layer and prevent the formation of an infinitesimally thin dust sheet around the mid-plane of the disk \\citep{Weidenschilling1980,WeidenschillingCuzzi1993}.  Modifications to the gravitational fragmentation scenario have been suggested to overcome the problem of Kelvin-Helmholtz turbulence. \\citeauthor{Sekiya1998} (1998, hereafter referred to as S98) found that if the mid-plane of the disk is in a state of constant Richardson number, as expected for small grains whose settling time is long compared to the growth rate of the KHI, then an increase in the global dust-to-gas ratio can lead to the formation of a high density dust cusp very close to the mid-plane of the disk, reaching potentially a dust-to-gas ratio of 100 already at a global dust-to-gas ratio that is 10 times the canonical interstellar value of $0.01$. The appearance of a superdense dust cusp in the very mid-plane has been interpreted by \\cite{YoudinShu2002} as an inability of the gas (or of the KHI) to move more mass than its own away from the mid-plane.  As a source of an increased value of the global dust-to-gas ratio, \\cite{YoudinShu2002} suggest that the dust grains falling radially inwards through the disk pile up in the inner disk. A slowly growing radial self-gravity mode in the dust density has also been suggested as the source of an increased dust-to-gas ratio at certain radial locations \\citep{Youdin2005a,Youdin2005b}. Trapping dust boulders in a turbulent flow is a mechanism for avoiding the problem of self-induced Kelvin-Helmholtz turbulence altogether \\citep{BargeSommeria1995,KlahrHenning1997,HodgsonBrandenburg1998,Chavanis2000,Johansen+etal2004}. If the dust can undergo a gravitational fragmentation locally, because the boulders are trapped in features of the turbulent gas flow such as vortices or high-pressure regions, then there is no need for an extremely dense dust layer around the mid-plane.  \\citet*{Johansen+etal2006} found that meter-sized dust boulders are temporarily trapped in regions of slight gas overdensity in magnetorotational turbulence, increasing the dust-to-gas ratio locally by up to two orders of magnitude. They estimate that the dust in such regions should have time to undergo gravitational fragmentation before the high-pressure regions dissolve again.  \\cite{FromangNelson2005}, on the other hand, find that vortices can even form in magnetorotationally turbulent disks, keeping dust boulders trapped for hundreds of disk rotation periods. The KHI cannot operate inside a vortex because there is no radial pressure gradient, and thus no vertical shear, in the center of the vortex \\citep{KlahrBodenheimer2006}.  From a numerical side it has been shown many times that a pure shear flow, i.e.\\ one that is not explicitly supported by any forces, is unstable, both with magnetic fields \\citep{Keppens+etal1999, KeppensToth1999} and without \\citep{BalbusHawleyStone1996}. But the key point here is that the vertical shear formed in a protoplanetary disk is due to the sedimentation of dust, and that the shear is able to regenerate as the dust falls down again, thus keeping the flow unstable to KHI. The description of the full non-linear outcome of such a system requires numerical simulations that include dust that can move relative to the gas.  Linear stability analysis of dust-induced shear flows in protoplanetary disks have been performed for simplified physical conditions \\citep{SekiyaIshitsu2000}, but also with Coriolis forces and Keplerian shear included \\citep{IshitsuSekiya2002,IshitsuSekiya2003}. Recently \\citeauthor{GomezOstriker2005} (2005, hereafter referred to as GO05) took an approach to include the dust into their numerical simulations of the Kelvin-Helmholtz instability by having the dust grains so extremely well-coupled to the gas that they always move with the instantaneous velocity of the gas. This is indeed a valid description of the dynamics of tiny dust grains. However, the strong coupling to the gas does not allow the dust grains to fall back towards the mid-plane. Thus the saturated state of the Kelvin-Helmholtz turbulence can not be reached this way.  In this paper we present computer simulations where we have let the dust grains, represented by particles each with an individual velocity vector and position, move relative to the gas. This allows us to obtain a state of self-sustained Kelvin-Helmholtz turbulence from which we can measure quantities such as the diffusion coefficient and the maximum dust density. A better knowledge of these important characteristics of Kelvin-Helmholtz turbulence is vital for our understanding of planet formation. ",
        "conclusions": "\\label{ch:conclusions}  The onset of the Kelvin-Helmholtz instability in protoplanetary disks has been known for decades to be the main obstacle for the formation of planetesimals via a gravitational collapse of the particle subdisk. Thus the study of the Kelvin-Helmholtz instability is one of the most intriguing problems of planetesimal formation. It is also a challenging problem to solve, both analytically and numerically, because of the coevolution of the two components gas and dust. Whereas turbulence normally arises from the gas flow alone, in Kelvin-Helmholtz turbulence the dust grains take the active part as the source of turbulence by piling up around the mid-plane and thus turning the energetically favored vertical rotation profile into an unstable shear. Planetesimal formation would be deceptively simple could the solids only sediment unhindered, but nature's dislike of thin shear flows precludes this by making the mid-plane turbulent.  In the current work we have shown numerically that when the dust particles are free to move relative to the gas, the Kelvin-Helmholtz turbulence acquires an equilibrium state where the vertical settling of the solids is balanced by the turbulent diffusion away from the mid-plane. For cm-sized pebbles and dm-sized rocks, we find that the dust component forms a layer that has a constant Richardson number. We thus confirm the analytical predictions by \\cite{Sekiya1998} for the first time in numerical simulations.  In the saturated turbulence we find the formation of highly overdense regions of solids, not in the mid-plane, but embedded in the turbulent flow. The clumping is very related to the streaming instability found by \\cite{YoudinGoodman2005}. Dust clumps with a density that is equal to or higher than the gas density orbit at the Keplerian velocity, so the clumps overtake sub-Keplerian regions of lower dust density.  Thus the dense clumps continue to grow in size and in mass. The final size of a dust clump is given by a balance between this feeding and the loss of material in a rarefaction tail that is formed behind the clump along the sub-Keplerian stream. The gravitational fragmentation of the single clumps into planetesimals is more likely than the whole dust layer fragmenting, because the local dust density in the clumps can be more than an order of magnitude higher than the azimuthally averaged mid-plane density. This process is very much related to the gravoturbulent formation of planetesimals in turbulent magnetohydrodynamical flows \\citep*{Johansen+etal2006}.  A full understanding of the role of Kelvin-Helmholtz turbulence in protoplanetary disks must eventually rely on simulations that include the effect of the Keplerian shear, so future simulations have to be extended into three dimensions. One can to first order expect that growth rates of the KHI larger than the shear rate $\\varOmega_0$ are required for a mode to grow in amplitude faster than it is being sheared out \\citep {IshitsuSekiya2003}, but so far it is an open question in how far the radial shear changes the appearance of the self-sustained state of Kelvin-Helmholz turbulence. Including furthermore the self-gravity between the dust particles, it will become feasible to study the formation of planetesimals in one self-consistent computer simulation and possibly to answer one of the outstanding questions in the planet formation process."
    },
    "0512/gr-qc0512003_arXiv.txt": {
        "abstract": "In this paper we analyze in detail some aspects of the proposed use of Ajisai and Jason-1, together with the LAGEOS satellites, to measure the general relativistic Lense-Thirring effect in the gravitational field of the Earth.  A linear combination of the nodes of such satellites is the proposed observable. The systematic error due to the mismodelling in the uncancelled even zonal harmonics would be $\\sim 1\\%$ according to the latest present-day CHAMP/GRACE-based Earth gravity models. In regard to the non-gravitational perturbations especially affecting Jason-1, only relatively high-frequency harmonic perturbations should occur: neither semisecular nor secular bias of non-gravitational origin should affect the proposed combination: their maximum impact is evaluated to $\\sim 4\\%$ over 2 years. Our estimation of the root-sum-square total error is about 4-5$\\%$ over at least 3 years of data analysis required to average out the uncancelled tidal perturbations. ",
        "introduction": "The most recent and relatively accurate test of the general relativistic gravitomagnetic Lense-Thirring effect on the orbit of a test particle (Lense and Thirring 1918; Barker and O'Connell 1974; Cugusi and Proverbio 1978; Soffel 1989; Ashby and Allison 1993; Iorio 2001) in the gravitational field of the Earth\\footnote{A more precise  (6$\\%$ on average) test of the Lense-Thirring effect was recently reported by Iorio (2006a) in the gravitational field of Mars. } was performed by Ciufolini and Pavlis (2004), who analyzed the laser data of the LAGEOS and LAGEOS II satellites according to a suitable combination of the residuals of their nodes proposed in (Ries et al. 2003a; 2003b; Iorio and Morea 2004) \\eqi\\delta\\dot\\Omega^{\\rm LAGEOS}+ c_1\\delta\\dot\\Omega^{\\rm LAGEOS\\ II }\\sim 48.1.\\lb{iorform}\\eqf  Let us briefly recall the linear combination approach from which \\rfr{iorform} originates. The combinations are obtained by explicitly writing down the expressions of the residuals of $N$ orbital elements (the nodes of different satellites in our case) in terms of the classical secular precessions induced by the mismodelled part of $N-1$ even zonal harmonic coefficient $\\delta J_{\\ell},\\ \\ell=2,4,...$ of the multipolar expansion of the terrestrial gravitational potential (see also Section \\ref{gravzon} and \\rfr{prc} for the meaning of the coefficients $\\dot\\Omega_{.\\ell}$) and the Lense-Thirring effect $\\dot\\Omega_{\\rm LT}$ considered as an entirely unmodelled feature of motion \\eqi\\delta\\dot\\Omega^{(i)}=\\sum_{\\ell=2}^{2(N-1)}\\dot\\Omega^{(i)}_{.\\ell}\\delta J_{\\ell}+\\dot\\Omega^{(i)}_{\\rm LT}\\mu_{\\rm LT},\\ i=1,2...N,\\lb{syst}\\eqf and solving the resulting algebraic non-homogeneous linear system of $N$ equations in $N$ unknowns of \\rfr{syst} with respect to the scaling parameter $\\mu_{\\rm LT}$ which is 1 in the Einsteinian theory and 0 in Newtonian mechanics. The obtained coefficients weighing the satellites' orbital elements depend on their semimajor axes $a$, eccentricities $e$ and inclinations $i$: they allow to cancel out the impact of the $N-1$ even zonal harmonics considered. In \\rfr{iorform} the  value of the secular trend predicted by the General Relativity Theory is 48.1 milliarcseconds per year (mas yr$^{-1}$), and $c_1=0.546$. The coefficient $c_1$ makes the combination of \\rfr{iorform} insensitive to the biasing action of only the first even zonal $J_2$  and its temporal variations.  The other even zonal harmonics $J_{\\ell\\geq 4}$, along with their secular variations $\\dot J_{\\ell\\geq 4}$, do affect \\rfr{iorform} inducing a systematic error in the measurement of the Lense-Thirring effect, whose correct and reliable evaluation is of crucial importance for the reliability of such an important test of fundamental physics. Ciufolini and Pavlis (2004), who used the GRACE-only Earth gravity model EIGEN-GRACE02S (Reigber et al 2005a), claimed a total error of $5\\%$ at 1-sigma and $10\\%$ at 3-sigma. Such estimates were criticized by Iorio (2005; 2006b) for various reasons. His evaluations, based on the analysis of different gravity model solutions and on the impact of the secular variations of the uncancelled even zonals, point toward a more conservative $\\sim 20\\%$ total error at 1-sigma.  The major drawbacks of the combination of \\rfr{iorform} are as follows \\begin{itemize} \\item It is mainly affected  by the low-degree even zonal harmonics $J_4, J_6$. The combination of  \\rfr{iorform} is practically insensitive to the even zonal harmonics of degree higher than $\\ell=12-14$ in the sense that the error induced by the uncancelled zonals does not change if the terms of degree higher than $\\ell=12-14$ are neglected in the calculation, as fully explained in Section \\ref{gravzon}. Unfortunately, the major improvements from the present-day and forthcoming GRACE models are mainly expected just for the medium-high degree even zonal harmonics which do not affect \\rfr{iorform}. Instead, the low-degree even zonals should not experience notable improvements, as showed by the most recent long-term models like EIGEN-CG01C (Reigber et al. 2006), EIGEN-CG03C (F\\\"{o}rste et al. 2005), EIGEN-GRACE02S, GGM02S (Tapley et al. 2005). Moreover, the part of the systematic error due to them is still rather model-dependent ranging from $\\sim 4\\%$ to $\\sim 9\\%$. \\item Another source of aliasing for the combination of \\rfr{iorform} is represented by the secular variations  $\\dot J_4$ and $\\dot J_6$ whose  signal grows quadratically in time. Their bias on the measurement of the Lense-Thirring effect with the combination of \\rfr{iorform} was evaluated to be of the order of $\\sim 10\\%$ (Iorio 2005). They are, at present, known with modest accuracy and there are few hopes that the situation could become more favorable in the near future. Moreover, also interannual variations of $J_4$ and $J_6$ may turn out to occur \\end{itemize} Thus, it seems unlikely that relevant improvements in the reliability and accuracy of the tests conducted with the adopted node-node combination of the LAGEOS satellites will occur in the foreseeable future.   In Iorio and Doornbos (2005) the following combination \\eqi\\delta\\dot\\Omega^{\\rm LAGEOS}+k_1\\delta\\dot\\Omega^{\\rm LAGEOS\\ II}+k_2\\delta\\dot\\Omega^{\\rm Ajisai}+k_3\\delta\\dot\\Omega^{\\rm Jason-1}=\\mu_{\\rm LT}49.5,\\lb{jason}\\eqf with \\eqi k_1=0.347,\\ k_2=-0.005,\\ k_3=0.068,\\lb{jasoncoef}\\eqf was designed: it comes from \\rfr{syst} applied to the nodes of LAGEOS, LAGEOS II, Jason-1 and Ajisai. A similar proposal was put forth by Vespe and Rutigliano (2005): however, the less accurate CHAMP-only Earth gravity model EIGEN3p (Reigber et al. 2005b) was used in that exhaustive analysis. Such a combination involves the nodes of the geodetic Ajisai satellite and of the radar altimeter Jason-1 satellite. Their orbital parameters, together with those of the LAGEOS satellites, are listed in Table \\ref{sat_par}. The combination of \\rfr{jason} allows cancellation of the first three even zonal harmonics $J_2,J_4, J_6$ along with their temporal variations. The resulting systematic error of gravitational origin is of the order of $\\sim 1\\%$. The practical implementation of the proposed test would consist in the following three stages \\begin{itemize} \\item The best possible nodes from independent arcs of data (for example, weekly) will be assembled as a time-series for the four satellites \\item Correspondingly, an integrated long-term node time-series will be constructed for each satellite with the best available dynamical models not using such force models derived empirically from the same tracking data determining the observed nodes (otherwise, also the Lense-Thirring effect would be removed. See Section \\ref{cazzo}) \\item From such two time-series a residual time-series will be built up for each satellite, combined according to \\rfr{jason} and analyzed for both secular and periodic terms; the secular component will be used to extract the Lense-Thirring effect \\end{itemize}   The goal of the present paper is to analyze in detail some important critical aspects of the use of such a combination. They are \\begin{itemize} \\item The impact of the higher degree even zonal harmonics introduced by the lower orbiting satellites Ajisai and Jason-1 \\item The impact of the realistically obtainable accuracy of a truly dynamical orbital reconstruction for Ajisai and Jason-1 \\item The impact of the atmospheric drag and of the other non-gravitational perturbations on Ajisai and, especially, Jason-1 \\end{itemize} ",
        "conclusions": "In this paper the use of a suitable linear combination of the nodes of LAGEOS, LAGEOS II, Ajisai and Jason-1 to measure the Lense-Thirring effect in the gravitational field of the Earth is examined. Below we list the major sources of errors along with our evaluations of their impact on the proposed measurement. They are also summarized in Table \\ref{tabwag}. \\subsection{The gravitational error} It turns out that the systematic error of gravitational origin due to the even zonal harmonics can be presently evaluated to be $\\sim 1\\%$, according to the latest Earth gravity models based on the combined data of CHAMP, GRACE and ground-based measurements. Such an estimate is rather model-independent and will be likely further improved when the new, forthcoming solutions for the terrestrial gravitational potential will be available. The temporal variations of the even zonal harmonics do not represent a major concern because the secular and possible interannual variations of the first three even zonal harmonics are cancelled out, by construction, along with their static components. Moreover, the uncancelled tidal perturbations, like the solar $K_1$ tide, vary with relatively high frequencies, so that they could be fitted and removed from the time-series or averaged out over an observational time span of at least 3 years (the longest period is that of the LAGEOS node amounting to 2.84 years). \\subsection{The measurement errors} Our largely conservative evaluation for the measurement errors amounts to $\\sim 3\\%/N$, where $N$ is the number of years of the experiment duration, by assuming a really pessimistic 1 m error in a truly dynamical orbit reconstruction for Ajisai and Jason-1 over the adopted time span. \\subsection{The non-gravitational error} In regard to the non-gravitational perturbations, which especially affect Jason-1, it is worthwhile noting that no secular aliasing trends should occur, but only high-frequency harmonic perturbations. However, particular attention should be paid to an as accurate as possible truly dynamical modelling of the non-gravitational accelerations acting on the node of Jason-1. Also a careful choice of the observational time span of the analysis would be required in order to reduce the uncertainties related to the orbital maneuvers which are mainly in plane, although a small, unknown, part of them affects also the out-of-plane part of the orbit. We evaluate the error due to the non-gravitational accelerations as large as $\\sim 4\\%$ over 2 years. \\subsection{Final remarks} In conclusion, the use of the proposed combination, although undoubtedly  difficult and demanding, seems to be reasonable and feasible; we give a total root-sum-square uncertainty of $\\sim$ 4-5$\\%$ over at least 3 years required to average out the uncancelled tidal perturbations. Moreover, the efforts required to perform the outlined analysis should be rewarding not only for the relativists' community but also for people involved in  space geodesy, altimetry and oceanography.    \\newpage"
    },
    "0512/astro-ph0512044_arXiv.txt": {
        "abstract": "We use images taken with the Advanced Camera for Surveys (ACS) on board the Hubble Space Telescope (HST) to derive effective radii and effective surface brightnesses of 15 early-type lens galaxies identified by the Sloan Lens ACS (SLACS) Survey as described in paper I. The structural parameters are used in combination with stellar velocity dispersions from the Sloan Digital Sky Survey (SDSS) database to construct the Fundamental Plane (FP) of lens galaxies. The size of the sample and the relatively narrow redshift range ($z=0.06-0.33$) allows us to investigate for the first time the distribution of lens galaxies in the FP space. After correcting the effective surface brightnesses for evolution, we find that lens galaxies occupy a subset of the local FP. The edge-on projection (approximately effective mass $M$ vs effective mass-to-light ratio $M/L$) is indistinguishable from that of normal early-type galaxies.  However -- within the fundamental plane -- the lens galaxies appear to concentrate at the edge of the region populated by normal early-type galaxies. We show that this is a result of our selection procedure, which gives higher priority to the highest velocity dispersions ($\\sigma\\gtrsim$240\\kms). Accounting for selection effects, the distribution of our galaxies inside the FP is indistinguishable from that of the parent sample of SDSS galaxies. We conclude that SLACS lenses are a fair sample of massive (high velocity dispersion) early-type galaxies. By comparing the central stellar velocity dispersion ($\\sigma$) with the velocity dispersion that best fits the lensing models ($\\sigma_{\\rm SIE}$; from paper III) we find $\\langle f_{\\rm SIE} \\rangle \\equiv \\langle \\sigma/\\sigma_{\\rm SIE} \\rangle =1.01\\pm0.02$ with 0.065 rms scatter. We conclude that within the Einstein radii (typically $R_e/2$ or $\\sim 4$ kpc) the SLACS lenses are very well approximated by isothermal ellipsoids, requiring a fine tuning of the stellar and dark matter distribution (bulge-halo ``conspiracy'') in the transition regions of early-type galaxies. Interpreting the offset from the local FP in terms of evolution of the stellar mass-to-light ratio, we find for the SLACS lenses $d \\log (M/L_{\\rm B}) /dz=-0.69\\pm0.08$ (rms 0.11) consistent with the rate found for field early-type galaxies and with a scenario where most of the stars were formed at high redshift ($>2$) with secondary episodes of star formation providing less than $\\sim 10$ \\% of the stellar mass below $z=1$.  We discuss star formation history and structural homogeneity in the context of formation mechanisms such as collisionless (``dry'') mergers. ",
        "introduction": "The properties of early-type galaxies in the local universe obey several empirical scaling laws such as the correlation between host velocity dispersion and mass of the central supermassive black hole (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000), correlations between velocity dispersion and stellar ages and chemical composition (Bender, Burstein \\& Faber 1992, 1993), and the correlation between velocity dispersion, effective radius, and effective surface brightness known as the Fundamental Plane (FP; Djorgovski \\& Davis 1987; Dressler et al.\\ 1987). Understanding the origin of these scaling laws is a challenge for galaxy formation models, because they imply a degree of homogeneity difficult to explain without invoking a substantial amount of fine tuning or feedback, still unexplained in the hierarchical merging scenario.  In particular, the Fundamental Plane can be seen as a scaling relation between a galaxy's effective (dynamical) mass and effective (dynamical) mass to light ratio (Faber et al.\\ 1987; Bender, Burstein \\& Faber 1992; van Albada, Bertin \\& Stiavelli 1995), in the sense that mass-to-light ratio increases with effective mass (the ``tilt'' of the FP). This tilt could be due (Ciotti, Lanzoni \\& Renzini 1996; Pahre 1998; Trujillo, Burkert \\& Bell 2004; Lanzoni et al. 2004) to trends in stellar populations (more massive galaxies are older and more metal rich), in the distribution of dark matter (more massive galaxies contain more dark matter), or in structural properties (the distribution function depends on mass, e.g. Bertin, Ciotti \\& del Principe 2002). The distribution of galaxies along the FP is also non-trivial: galaxies do not occupy the whole plane but live in well defined zones, avoiding a well defined region of the plane (``zone of avoidance''). From a theoretical point of view it is difficult to reproduce both the tilt and the distribution of galaxies inside the FP (e.g. Nipoti, Londrillo \\& Ciotti 2003).  In recent years, several groups have measured the FP of early-type galaxies at cosmological redshift (van Dokkum \\& Franx 1996; Kelson et al. 1997; Pahre 1998; Bender et al.\\ 1998; van Dokkum et al.\\ 1998; Treu et al.\\ 1999; Kelson et al. 2000; van Dokkum et al. 2001; Treu et al.\\ 2001a,b, 2002; van Dokkum \\& Ellis 2003; Gebhardt et al.\\ 2003; van Dokkum \\& Stanford 2003; van der Wel et al. 2004; Wujits et al.\\ 2004; Fritz et al. 2005; Holden et al. 2005; Treu et al. 2005a,b; van der Wel. et al. 2005; di Serego Alighieri et al.\\ 2005; Moran et al.\\ 2005) to measure the evolution of their mass to light ratio and hence constrain their star formation history. The results of these studies can be summarized as follows: massive early-type galaxies (above 10$^{11.5}$ M$_{\\odot}$) appear to be evolving slowly below z$\\sim$1 consistent with passive evolution of an old stellar population formed at $z>2$. At smaller masses, signs of recent star formation start to appear first in the field (at $z\\sim0.5$; Treu et al.\\ 2002; 2005b, van der Wel et al.\\ 2005) and at higher redshifts possibly even in clusters (Holden et al.\\ 2005). Mass seems to be the dominant parameter determining the star formation history, while environment affects $M/L$ only to a lesser degree (Moran et al.\\ 2005; Yee et al.\\ 2005). This trend is consistent with the {\\it downsizing} (i.e. star formation activity moves to lower masses from high to low redshift; Cowie et al. 1996) scenario seen in a number of studies (e.g. McIntosh et al.\\ 2005; Treu et al.\\ 2005a; Juneau et al.\\ 2005; see de Lucia et al.\\ 2006 for a theoretical point of view; see Gavazzi 1993 for a discussion of downsizing {\\it ante-litteram} for disk galaxies, based on fossil evidence). In the case of early-type galaxies, the concept of downsizing refers to the fact that stars are oldest in the most massive systems.  Little is known about scaling relations of early-type lens galaxies. Do they follow the same scaling laws of normal early-type galaxies? In principle, anomalous structural properties -- such as an unusually high concentration/mass density -- or mass from a large scale structure (e.g. a group or a filament) projected along the line of sight could boost lensing efficiency and therefore a lensing-selected sample could be biased. Based on the limited amount of information available so far, no significant difference has been found between the structural properties (e.g. Treu \\& Koopmans 2004) of lens and non-lens galaxies, supporting the hypothesis that E/S0 lenses are representative of the whole E/S0 population.  In this paper, we exploit the large and homogeneously selected sample of lenses identified in a relatively narrow redshift range by the Sloan Lenses ACS Survey (SLACS; Bolton et al.\\ 2005; Bolton et al. 2006, hereafter paper I; www.slacs.org), to study in detail the Fundamental Plane of lens galaxies, both in terms of tilt and distribution of galaxies along the plane. We quantify the degree of homogeneity of the early-type galaxies by measuring the ratio between stellar velocity dispersion and velocity dispersion of the singular isothermal ellipsoid (SIE) mass model that best fits the geometry of the multiple images, using the results of the lens models derived by Koopmans et al.\\ (2006; hereafter paper III).  In addition, we use the sample to study the evolution of the FP of lens galaxies with redshift as a diagnostic of their stellar populations. Previous studies on this topic have reached discordant conclusions. In a pioneering work, Kochanek et al.\\ (2000; see also Rusin et al.\\ 2003) assumed isothermal mass density profiles (see Rusin \\& Kochanek 2005 for a more general approach) to convert image separations into velocity dispersion and construct the FP of early-type lens galaxies without stellar velocity dispersion information all the way out to $z\\sim1$, finding relatively slow evolution of their mass-to-light ratio ($d \\log (M/L_{\\rm B})/dz= -0.56\\pm0.04$) at variance with the faster evolution measured by direct determinations of the stellar velocity dispersions (e.g. $-0.72^{+0.07}_{-0.05}\\pm0.04$ Treu et al.\\ 2005b; $-0.76\\pm0.07$ van der Wel et al. 2005).  Other than a genuine intrinsic difference in the stellar populations, the origin of this discrepancy could be attributed to departures from isothermality, to environmental effects, to differences in the analysis and fitting techniques (e.g. van de Ven, van Dokkum \\& Franx 2003 reanalyzed the Rusin et al.\\ sample, obtaining faster evolution $d \\log (M/L_B)/dz=-0.62\\pm0.13$). Or else, the difference could be understood in terms of downsizing. In fact -- as a result of the shape of the luminosity function of early-type galaxies and the scaling of lensing cross section with mass -- lens samples tend include mostly massive early-type galaxies ($\\sigma\\sim250$ \\kms) as opposed to luminosity selected samples which tend to be dominated by galaxies close to the limiting magnitude.  Early results on the distribution of mass and light in E/S0 lenses were obtained by the LSD survey (Koopmans \\& Treu 2002,2003; Treu \\& Koopmans 2002,2003,2004; hereafter collectively KT), which published stellar velocity dispersions for 5 lenses in the redshift range $z\\approx0.5-1$. KT found that the image separation under isothermal assumptions provided a good approximation for stellar velocity dispersion and that the evolution of stellar mass to light ratio was in good agreement with that measured for non-lens samples, albeit with large error bars because of the small sample size. The SLACS sample is the largest sample of lenses to date with measured stellar velocity dispersions. It is therefore ideally suited to be combined with the LSD sample to investigate with higher precision any difference in the stellar populations of lens and non-lens early-type galaxies, bearing in mind the SLACS selection process (luminous red galaxies and/or quiescent spectra, H$\\alpha<1.5$\\AA; see paper I and Section 2.).  A plan of the paper follows.  Section~2 presents the sample and the data. Section~3 investigates the internal structure and homogeneity of SLACS lenses using the FP as a diagnostic tool. Section~4 describes the redshift evolution of the FP of the combined SLACS+LSD sample and compares it to that of non-lens early-type galaxies. Section~5 summarizes our results. As in the rest of this series, we adopt Vega magnitudes and a cosmological model with $\\Omega_{\\rm m} = 0.3$, $\\Omega_\\Lambda = 0.7$, and H$_0 = 70\\,h_{70}$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ (with $h_{70} = 1$ when needed). ",
        "conclusions": "\\label{sec:sum}  We have performed 2D surface photometry on the HST-ACS images of 15 lens early-type galaxies identified by SLACS. As described in paper I (Bolton et al.\\ 2006) the lenses are selected from the luminous red galaxy and main galaxy sample of the SDSS database with quiescent spectra, i.e. equivalent width of H$\\alpha<1.5$\\AA. In combination with stellar velocity dispersions from the SDSS database, we have constructed the Fundamental Plane of lens galaxies and measured its evolution with redshift. We have compared the measured stellar velocity dispersion with the velocity dispersion of the singular isothermal ellipsoid that best fits the lensing geometry (c.f. paper III) to study the homogeneity of lens galaxies and the accuracy of the isothermal approximation to measure the evolution of the FP of lens galaxies. The main results can be summarized as follows:  \\begin{enumerate}  \\item SLACS lenses define a Fundamental Plane correlation over almost a decade in effective radii. The lenses are typically brighter than local early-type galaxies for a given velocity dispersion and effective radius, consistent with lower mass to light ratios, i.e. younger stellar populations at $z=0.06-0.33$ than today.  \\item After correction for evolution of the stellar populations, the SLACS lenses fall on the FP of early-type galaxies in the local Universe.  The edge-on projection of the FP of SLACS lenses is consistent with that of local galaxies within the errors. In contrast, SLACS lenses occupy a relatively small portion of the plane, concentrating along the border of the so-called ``zone of avoidance'' of local early-type galaxies. We show that this is the result of our selection procedure focused on the systems with the highest velocity dispersion. Accounting for the selection procedure, the distribution of distances from the ``zone of avoidance'' is indistinguishable from that of the SDSS parent samples (MAIN and LRG). We conclude that the SLACS lenses are a fair sample of high velocity dispersion ($\\sigma \\gtrsim 240$ \\kms) early-type galaxies.  \\item The ratio between the central stellar velocity dispersion ($\\sigma$) and velocity dispersion of the singular isothermal ellipsoid ($\\sigma_{\\rm SIE}$) that best fits the lensing geometry is found to be $\\langle f_{\\rm SIE} \\rangle =1.010\\pm0.017$ with an rms scatter of 0.065.  The isothermal approximation for the SLACS lenses works better than that for the LSD sample of 5 galaxies at higher redshifts ($\\langle f_{\\rm SIE} \\rangle =0.87$ with rms scatter 0.08). If this redshift dependency is confirmed by larger and homogeneously selected samples at higher redshift, possible explanations include an intrinsic change in the properties of early-type galaxies with cosmic time or simply by an increased contribution of external convergence, resulting from group or clusters associated with the lens or large scale structures along the line of sight.  \\item Interpreting the evolution of the FP in terms of evolution of the stellar populations, the effective mass to light ratio of SLACS lenses evolves as $d \\log (M/L_{\\rm B}) / dz = -0.69\\pm0.08$ with an rms scatter of 0.11. Adding the 5 galaxies from the LSD sample, the best fit evolutionary rate changes to $-0.76\\pm0.03$ leaving the scatter unchanged. The evolutionary rate changes within the error if $\\sigma_{\\rm SIE}$ is used to construct the FP instead of $\\sigma$.  \\end{enumerate}  We now briefly discuss these results in terms of their implications for our understanding of the formation and evolution of early-type galaxies\\footnote{To avoid duplications within the series, the main discussion is left for paper III.}.  From the point of view of stellar populations, we find that the SLACS lenses have mostly old stellar population, with at most a small ($<10$\\%) contribution of stellar mass accreted at $z<1$. This is in agreement with what is found (e.g. Treu et al. 2005a,b) for non-lens early-type galaxies of comparable mass ($\\langle\\sigma \\rangle=279$, r.m.s. scatter 45 km s$^{-1}$). This conclusion is unlikely to be significantly biased by our selection criterion against H$\\alpha$ emission, since at the redshifts of the current SLACS sample ($\\langle z \\rangle $=0.19), emission lines are not frequent even in morphologically selected samples (e.g. Treu et al.\\ 2002).  Concluding that most stars are old, however, does not answer the question of how and when the stars and dark matter are assembled. An often-invoked mechanism for the assembly of massive early-type galaxies (e.g. Khochfar \\& Burkert 2003) are the so-called ``dry'' mergers, i.e. mergers that do not involve a significant amount of cold gas and hence not associated with star formation (see Bell et al.\\ 2006 and van Dokkum 2005 for recent results and discussions). Since the stars in merging galaxies are already old and the dynamical timescales for mergers are rather short (a few hundred million years), dry mergers provide an efficient way to assemble large amounts of old stars with little remnants.  However, the old stellar ages are only one of the observational tests of the dry mergers hypothesis. The tight scatter of empirical scaling laws such as the Fundamental Plane and the black-hole mass $\\sigma$ relation must also be explained under this scenario. Numerical simulations show that plausible configurations of major dry mergers preserve the tightness of the FP, i.e. if two galaxies start on the FP they also end up on the FP (Nipoti, Londrillo \\& Ciotti 2003; Gonzalez-Garcia \\& van Albada 2003; Boylan-Kolchin, Ma \\& Quataert 2005), with an edge-on thickness comparable to that observed. In contrast, other scaling laws such as the Kormendy and Faber-Jackson relationships and the black-hole mass $\\sigma$ relationship, are not naturally preserved by dry mergers (Nipoti, Londrillo \\& Ciotti 2003). A substantial amount of dissipation seems to be necessary to preserve those (e.g. Kazantzidis et al. 2005).  The observations present here provide a new series of tests for this hypothesis. First, we measure the distribution of early-type galaxies within the FP: this also should be reproduced by a successful formation mechanism. At the moment there is no good explanation for the ``zone of avoidance'', nor does it appear clear whether dry mergers alone can populate the right portion of the plane (as suggested perhaps by the difficulties in reproducing the Kormendy and Faber-Jackson relation). It is possible that dissipation may end up to be necessary to move the objects towards higher concentration, against the boundary of the ``zone of avoidance'' (see Bender, Burstein \\& Faber 1992).  Second, the proposed mechanism must also explain the distribution of mass {\\it inside} each galaxy. Our study shows very clearly that, at these scales (typically $R_{\\rm e}/2$), early-type lens galaxies are well approximated by singular isothermal ellipsoids (see also KT and paper III). Not only this mass density profiles differs significantly from cosmologically motivated ones (e.g. Navarro et al.\\ 1997; Moore et al.\\ 1998), but also it requires a significant amount of fine tuning between the distribution of luminous and dark matter (bulge-halo ``conspiracy''). It has been suggested that the total mass density profile can act as a ``dynamical attractor'' for collisionless particles (i.e. stars and dark matter; Loeb \\& Peebles 2003; Gao et al. 2004) in a similar way to the close-to-isothermal profiles obtained for (incomplete) violent relaxation scenarios (e.g. van Albada 1982). Our measurements provide a quantitative test for this idea. From the point of view of dry mergers, the close-to-isothermal mass density profile raises two problems. The first is whether such property is preserved during dry mergers. Assuming that this is the case, the second problem is how did the progenitors get their initial mass density profile. Simulations of dry mergers typically start with input galaxies already ``dense'' and close-to-isothermal, but -- as we argue in paper III -- this process cannot be traced back {\\it ad infinitum}. A different process -- perhaps gas rich mergers -- -- appear to be needed at some point to create these high concentration objects. This brings us back to the main unanswered question: what is the origin of the bulge-halo ``conspiracy''?"
    },
    "0512/astro-ph0512102_arXiv.txt": {
        "abstract": "The Advanced Camera for Surveys (ACS) enables low resolution slitless spectroscopic imaging in the three channels. The most-used modes are grism imaging with the WFC and the HRC at a resolution of $40$ and 24 \\AA/pixel, respectively. In the far UV there are two prisms for the SBC and a prism for the HRC in the near-UV. An overview of the slitless spectroscopic modes of the ACS is presented together with the advantages of slitless spectroscopy from space. The methods and strategies developed to establish and maintain the wavelength and flux calibration for the different channels are outlined. Since many slitless spectra are recorded on one deep exposure, pipeline science quality extraction of spectra is a necessity. To reduce ACS slitless data, the aXe spectral extraction software has been developed at the ST-ECF. aXe was designed to extract large numbers of ACS slitless spectra in an unsupervised way based on an input catalogue derived from a companion direct image. In order to handle dithered slitless spectra, drizzle, well-known in the imaging domain, has been applied. For ACS grism images, the aXedrizzle technique resamples 2D spectra from individual images to deep, rectified images before extracting the 1D spectra. aXe also provides tools for visual assessment of the extracted spectra and examples are presented. ",
        "introduction": "The ACS has three channels, the Wide Field Channel (WFC), the High resolution Channel (HRC) and the Solar Blind Channel (SBC), and each channel is capable of delivering slitless spectroscopic images by inserting a grism or prism into the optical beam. The five different combinations of ACS channel and dispersing element offer low resolution spectroscopy from the UV to the far-red wavelength regime. Table \\ref{spec_modes} lists important parameters such as spectral resolution, wavelength range and field of view (FoV) for all slitless modes of ACS.\\  \\begin{table} \\begin{tabular}{cccccc} \\tableline \\tableline Channel & Disperser & Wav. Range & Resolution & Pixel Size & FOV \\\\ &           &      [\\AA]       & [\\AA/pixel]& [mas/pixel]& [arcsecond]\\\\ \\tableline WFC     & G800L     & $5500-10500$     &     $38.5$   & $50\\times 50$& $202\\times 202$\\\\ HRC     & G800L     & $5500-10500$     &     $23.5$   & $28\\times 24$& $29\\times 26$\\\\ HRC     & PR200L     & $1700-3900$     & 20[@2500\\AA] & $28\\times 24$& $29\\times 26$\\\\ SBC     & PR130L     & $1250-1800$     & 7[@1500\\AA]   & $34\\times 30$& $35\\times 31$\\\\ HRC     & PR110L     & $1150-1800$     & 10[@1500\\AA] & $34\\times 30$& $35\\times 31$\\\\ \\tableline \\end{tabular} \\caption{The spectroscopic modes of the ACS.} \\label{spec_modes} \\end{table} Figure \\ref{example-fig-1} illustrates the differences of using grisms and prisms as dispersive elements by comparing the dispersion and the sensitivity of the HRC grism G800L (left) and the HRC prism PR200L (right). The G800L has an almost constant dispersion over the entire accessible wavelength range. The dispersion of the PR200L increases drastically towards longer wavelengths, and any value for the dispersion must be accompanied by the wavelength at which it is specified (see Tab.\\ \\ref{spec_modes} and the ACS Instrument Handbook). The different properties of the prisms and grisms require a flexible data reduction software to be able to reduce both slitless grism and prism data.\\\\  Following the demise of STIS in August 2004 the interest in using ACS grisms and prisms has increased substantially, since it is the only optical-UV spectral capability aboard HST. As a result, around 10\\% of all approved orbits in Cycle 14 are devoted to slitless spectroscopy mode with ACS (Macchetto et al.\\ 2005).  \\begin{figure} \\epsscale{0.30} \\includegraphics[angle=-90,width=.8\\textwidth]{kuemmelF1.ps} \\caption{A comparison between dispersion and sensitivity of the HRC grism G800L (left) and prism PR200L (right).} \\label{example-fig-1} \\end{figure} ",
        "conclusions": "\\begin{itemize} \\item All aspects of ACS slitless spectroscopy (calibration, software etc.,) are supported such that users can obtain and reduce slitless data in a pipeline way. \\item ACS slitless spectroscopy is successfully used in various science projects such as the HUDF HRC Parallels (Walsh, K\\\"ummel \\& Larsen 2004), high redshift supernovae research (Riess et al. 2004) and the search for high redshift galaxies (the GRAPES and PEARS programmes, see Pirzkal et al., 2004). \\item More information about ACS slitless spectroscopy, the calibration and the aXe software is given on the aXe webpages at http://www.stecf.org/software/aXe/ \\item User support concerning all topics related to ACS slitless spectroscopy is provided by the ACS group at the Space Telescope - European Coordinating Facility (ST-ECF). The centralised email address for requests is acsdesk@eso.org. \\end{itemize}"
    },
    "0512/astro-ph0512428_arXiv.txt": {
        "abstract": "% To test the existence of a possible radial gradient in oxygen abundances within the Local Group dwarf irregular galaxy NGC~6822, we have obtained optical spectra of 19 nebulae with the EFOSC2 spectrograph on the 3.6-m telescope at ESO La Silla. The extent of the measured nebulae spans galactocentric radii in the range between 0.05 kpc and 2 kpc (over four exponential scale lengths). In five \\hii\\ regions (Hubble~I, Hubble~V, K$\\alpha$, K$\\beta$, KD~28e), the temperature-sensitive \\othreea\\ emission line was detected, and direct oxygen abundances were derived. Oxygen abundances for the remaining \\hii\\ regions were derived using bright-line methods. The oxygen abundances for three A-type supergiant stars are slightly higher than nebular values at comparable radii. Linear least-square fits to various subsets of abundance data were obtained. When all of the measured nebulae are included, no clear signature is found for an abundance gradient. A fit to only newly observed \\hii\\ regions with \\othreea\\ detections yields an oxygen abundance gradient of $-0.14 \\pm 0.07$ dex~kpc$^{-1}$. The gradient becomes slightly more significant ($-0.16 \\pm 0.05$ dex~kpc$^{-1}$) when three additional \\hii\\ regions with \\othreea\\ measurements from the literature are added. Assuming no abundance gradient, we derive a mean nebular oxygen abundance 12$+$log(O/H) = $8.11 \\pm 0.10$ from \\othreea\\ detections in the five \\hii\\ regions from our present data; this mean value corresponds to [O/H] = $-0.55$. ",
        "introduction": "%  The evolution of element abundances with time provides important clues to the history of chemical enrichment and star formation in galaxies. Abundances in \\hii\\ regions provide information about the most recent period of metals enrichment in the interstellar medium, now illuminated by recently-formed O, B-type stars. Dwarf irregular galaxies contain few metals and are thought to be chemically homogeneous. These low-mass, metals-poor, gas-rich, and star-forming systems provide unique venues to examine detailed processes of star formation within environments which may be related to star-forming systems observed at early times. Thus far, no nearby ($D \\approx 5$ Mpc) dwarf irregular galaxy has ever exhibited a significant radial abundance gradient (i.e., \\citealp{ks96,ks97,ls04,lsv05}). A possible reason is that the most recent burst of star formation has expelled the most recent synthesis of metals into the hot phase of the interstellar medium, and that these metals have yet to cool and ``rain'' back down onto the cooler phases of the interstellar medium (e.g., \\citealp{tt96,ks97}). It would indeed be very interesting to find an example of a dwarf galaxy with localized oxygen enrichment. Recent developments in high-efficiency spectrographs on 8- and 10-m telescopes have made possible spectroscopy of bright blue A-type supergiants in nearby dwarf irregular galaxies (e.g., \\citealp{venn03,kaufer04}). These young hot massive stars allow for the simultaneous measurements of present-day $\\alpha$- and iron-group elements. These measurements also allow for the direct comparison of stellar $\\alpha$-element abundances with nebular measurements, as massive stars and nebulae are similar in age and have similar formation sites.  NGC~6822 is the nearest gas-rich dwarf irregular galaxy in the Local Group \\citep{mateo98}. Basic properties are listed in Table~\\ref{table_n6822}; see also \\cite{vdb00}. Studies of the stellar populations and the star formation history of NGC~6822 are described by \\cite{gallart96a,gallart96b,gallart96c}, \\cite{cb98}, \\cite{hutchings99}, \\cite{tolstoy01}, \\cite{wyder01}, and \\cite{clementini03}. A large \\hi\\ halo extending much farther out than the optical extent was discovered by \\cite{roberts72}, and confirmed by \\cite{debw00} and \\cite{komiyama03}. A population of young, blue stars was found, whose spatial distribution is farther than the accepted optical radius and is very similar to that of the \\hi\\ extent \\citep{battinelli03,debw03,komiyama03}. With the characterization of the spatial and color-magnitude distributions of red giant and asymptotic giant branch stars across the entire galaxy in the near-infrared, \\cite{ch05} found a metallicity spread of 1.56~dex from the ratio of carbon- to M-type stars. \\cite{leisy05} have recently reported 13 new candidate planetary nebulae (PNe), bringing to 17 total PNe in all. \\cite{musch99} obtained spectra of three B-type supergiant stars, and reported a mean iron abundance of [Fe/H] = $-0.5 \\pm 0.2$$\\,$\\footnote{ We use the notation: [X/Y] = log(X/Y) $-$ log(X/Y)$_{\\odot}$. }. \\cite{venn01} reported for the first time oxygen abundances for two A-type supergiant stars in NGC~6822, and showed that the mean stellar oxygen abundance was higher than the known \\hii\\ region nebular abundances by at least 0.1~dex. Because their stars were at low galactocentric radii, the authors suggested the possibility of a radial gradient in oxygen abundance.  The recent measurements of oxygen abundances in A-type supergiant stars have motivated the reevaluation of published nebular oxygen abundances in NGC~6822. Spectra of and abundances for the brightest \\hii\\ regions in NGC~6822 were reported by \\cite{ps70}, \\cite{alloin74}\\footnote{ \\cite{alloin74} did not derive an oxygen abundance as their \\othree\\ emission lines were saturated and not measured. }, \\cite{smith75}, \\cite{lequeux79}, \\cite{talent80}, and \\cite{pes80}. However, all of these spectra were obtained with inherently nonlinear detectors, and in many cases the character of the nonlinearities were not understood until well after publication (e.g., \\citealp{jenkins87}); so, subsequent corrections for nonlinearity were not possible. \\cite{stm89} obtained CCD spectroscopy to measure \\othreea\\ in Hubble~V, and derived an oxygen abundance of 12$+$log(O/H) = 8.20. In two planetary nebulae, \\cite{rm95} derived oxygen abundances (12$+$log(O/H) = 8.01, 8.10) in agreement with published \\hii\\ region oxygen abundances. \\cite{miller96} remeasured Hubble~V and Hubble~X, and derived oxygen abundances 12$+$log(O/H) = 8.32 and 8.36, respectively. In their program on open clusters, \\cite{chandar00} also obtained spectra for a few nebulae, and derived smaller nebular oxygen abundances than expected. We have reanalyzed a number of their \\hii\\ region spectra, which we discuss in Sect.~\\ref{sec_discuss}. \\cite{hgom01} also obtained spectra of Hubble~V and Hubble~X, and while there were no large differences in oxygen abundances between the two \\hii\\ regions, they claimed small-scale abundances variations on $<$ 10~pc length scales. \\cite{josh02} derived sulfur abundances from {\\em ISO\\/} measurements of [\\ion{S}{3}] and [\\ion{S}{4}] emission lines in the mid-infrared, and showed that the S$^{+3}$ ion is the largest contributor to the total sulfur abundance in extragalactic \\hii\\ regions. With VLT data, \\cite{peimbert05} derived recombination-line abundances for the \\hii\\ region Hubble~V, and showed that their derived oxygen abundances were in better agreement with the stellar oxygen abundances. Some of these results are discussed further in Sec.~\\ref{sec_recentwork} below.  This is the second of two papers of our study examining the spatial homogeneity of oxygen abundances in Local Group dwarf irregular galaxies; WLM was discussed previously in \\cite{lsv05}. The main goal here was to obtain a homogeneous set of nebular spectra for \\hii\\ regions in NGC~6822 over a large range in galactocentric radii. The outline of this paper is as follows. Descriptions of the observations, reductions, measurements and analysis are presented in Sect.~\\ref{sec_obsmeas}. Element abundances and abundance ratios are described in Sect.~\\ref{sec_abund}. In Sect.~\\ref{sec_discuss}, we compare present results with recent studies, and examine the presence of spatial inhomogeneities in oxygen abundances. A summary is given in Sect.~\\ref{sec_concl}. For the present discussion, we use the notation [O/H] = log(O/H) $-$ log(O/H)$_{\\odot}$, where the solar value of the oxygen abundance is 12$+$log(O/H) = 8.66 \\citep{asplund04,melendez04}. ",
        "conclusions": "% \\label{sec_concl}  Optical spectra were measured for 19 nebulae in NGC~6822 at galactocentric radii between 0.05 kpc and 2 kpc, or out to about four exponential scale lengths. \\othreea\\ was detected in five \\hii\\ regions, and subsequent direct oxygen abundances were derived. Oxygen abundances for the remaining \\hii\\ regions were derived using bright-line methods. Oxygen abundances for the A-type supergiant stars are consistent with \\hii\\ region abundances at comparable galactocentric radii. Linear least-square fits to various subsets of abundance data were obtained. When all of the measured nebulae are included, no clear signature is found for an abundance gradient out to over four exponential scale lengths. A fit to \\hii\\ regions with only \\othreea\\ detections is consistent with zero slope ($2\\sigma$). The abundance gradient becomes slightly more significant ($3.2\\sigma$) when three additional \\hii\\ regions with \\othreea\\ measurements from the literature are included. The resulting slope and extrapolated central abundance are $-0.16 \\pm 0.05$ dex~kpc$^{-1}$ and 12$+$log(O/H) = $8.23 \\pm 0.05$, respectively. Assuming zero abundance gradient, we take the five \\hii\\ regions with \\othreea\\ detections from our present work, and derive a mean nebular oxygen abundance 12$+$log(O/H) = $8.11 \\pm 0.10$, which corresponds to [O/H] = $-0.55$. Additional deep high-quality spectra of nebulae and stars are required to distinguish clearly between either a zero or a non-zero slope. The latter would be confirmation of an abundance gradient seen for the first time in a dwarf irregular galaxy."
    },
    "0512/nucl-th0512093_arXiv.txt": {
        "abstract": "An improvement in the treatment of the isovector channel of relativistic mean field (RMF) models based on effective field theory (E-RMF) is suggested, by adding an isovector scalar ($\\delta$) meson and using a similar procedure to the one used by Horowitz and Piekarewicz to adjust the isovector-vector channel in order to achieve a softer density dependent symmetry energy of the nuclear matter at high density. Their effects on the equation of state (EOS) at high density and on the neutrino mean free path (NMFP) in neutron stars are discussed. ",
        "introduction": "\\label{sec_intro} Neutrino transport in stellar matter plays an important role in some phenomena such as the mechanism of supernovae explosions, structure of protoneutron stars, etc. Theoretical input needed for understanding the neutrino transport is the neutrino mean free path (NMFP) and the equation of state (EOS). A unified matter  model used in calculating  both observables is a requirement for having a  consistent neutrino transport result. We note that considerable efforts have been devoted to investigate the neutrino interaction in matter at high density~\\cite{reddy1,niembro01,parada04,horo1,mornas1,mornas2,reddy2,horowitz91,yama,caiwan,margue,chandra,cowel,caroline05} with different matter models, variety of approximation  levels and purposes.  To this end, a kind of matter models that can be used is the relativistic mean field (RMF) model. The original or standard RMF models use the sigma, omega, and rho mesons with  additional cubic and quartic nonlinearities of sigma meson to effectively describe the interaction among nucleons. NL-3 parameter set~\\cite{lala} belongs to this model. The parameter set has been very successful in the description of a variety of ground state properties of nuclei~\\cite{todd}. Details of the models as well as their applications can be seen in Refs.~\\cite{pg,ring,serot}. The effects of including delta meson in standard RMF models  including the corresponding linear responses have been studied for the asymmetric nuclear matter at low densities in Refs.~\\cite{kubis97,liu02,greco03}, whereas for heavy ion collisions in Refs.~\\cite{greco2,gaitanos1,gaitanos2} and on the neutron star properties in Ref.~\\cite{liu04}. It was found that the $\\delta$ field leads to a larger repulsion in the dense neutron rich matter (stiffer symmetry energy that leads to a larger proton fraction at high density), as well as a definite splitting of proton and neutron effective masses. Both features are influencing the stability conditions of a neutron star~\\cite{liu04}.  Inspired by the concept of effective field and density functional theories for hadrons, Furnstahl, Serot and Tang~\\cite{Furnstahl96} constructed a new RMF model (from now on, will be denoted by E-RMF) that can be considered as an extension of the standard RMF models. One of the parameter sets in this model is G2. Besides yielding accurate predictions in finite nuclei and normal nuclear matter~\\cite{serot,Furnstahl96,sil}, G2 has the interesting features like a positive value of quartic sigma meson coupling constant that leads to the existing lower bound in energy spectrum of this model~\\cite{arumu,baym} and to the missing zero sound mode in the high density symmetric nuclear matter~\\cite{cailon}. Moreover, the agreement of the nuclear matter and the neutron matter EOS at high density of G2 with the Dirac Brueckner Hartree Fock (DBHF) calculation ~\\cite{sil,arumu} is better than those of NL-3, NL-1 and TM1 models (the standard RMF plus a quartic omega meson interaction). Nevertheless, from  the comparison between the neutron matter EOS of this model and that of the DBHF  result, the authors of Ref.~\\cite{sil} pointed out that the present form of the E-RMF model still needs a substantial  improvement in the treatment of the isovector sector. It has also been shown  that the G2 parameters set of the E-RMF model still predicts a too large proton fraction~\\cite{parada04}.  It is known that proton fraction correlates to the direct URCA cooling process~\\cite{lati}. It is also known that this problem is caused by the role of isovector terms.  So far, the effects of the  delta meson inclusion on this model has not been studied yet.  Therefore, in this work, first, we will study the effects of adding an isovector-scalar ($\\delta$) meson in the E-RMF model and the effects of the  isovector-vector channel adjustment by using a similar procedure to the one used in Ref.~\\cite{Horowitz01}. The aim of these adjustments is to achieve a softer density dependent symmetry energy of nuclear matter at high density. The symmetry energy has a wide range of effects, such as  from giant dipole resonances to heavy ion collisions in nuclear physics and from supernovae to neutron stars  in astrophysics. In spite of its diverse influences, its magnitude and density dependence are not well understood~\\cite{steiner}. More detail informations on the role of symmetry energy and related topics can be consulted to Ref.~\\cite{steiner} and references therein. Second, we will also extend the analysis of our previous report~\\cite{parada04} to give a more solid argument about the source of the possible appearance of the anomalous behavior in the NMFP of neutron stars predicted by RMF models. Here the anomalous behavior in the NMFP  means a contra intuitive NMFP results in a form of the decreasing of the mean free path with respect to the decreasing of the matter density~\\cite{niembro01}. It has been known that this anomalous behavior exists in the NMFP predicted by non relativistic nuclear models of the Skyrme type~\\cite{reddy1,niembro01}. This  anomalous behavior is attributed to the dominating term at high density that responsible to the appearance of the acausal behavior (the speed of sound exceeds the speed of light) of the model at high densities~\\cite{reddy1}. It has been reported in Refs.~\\cite{reddy1,reddy2,niembro01} that relativistic models alleviate this problem. But we have eventually found that not every parameter set of relativistic models is free from this problem~\\cite{parada04}.  In section~\\ref{sec_frml} we will briefly explain the  formalism used in this work. In section ~\\ref{sec_rslt} we discuss the results of our calculations. We will give the summary of our findings in section~\\ref{sec_sum}. ",
        "conclusions": "\\label{sec_rslt} In this section we study the effects of the isovector-vector channel adjustment and the addition of $\\delta$ meson in E-RMF models on the corresponding nuclear matter properties predictions. We also study  the role of every factor involved in the predicted NMFP in neutron stars of RMF models. We start with considering the effects of isovector-vector channel of the E-RMF model. Here we analyze the effects of  different g$_{\\rho}$ and $\\eta_{\\rho}$ combinations of the G2 parameter set. The value of various coupling constant combinations can be seen in Table~\\ref{tab:params2}. Their effects on matter properties are shown in Fig.~\\ref{G2fac2}.  \\begin{table} \\centering \\caption {Isovector-vector channel adjustment in the G2 parameter set of E-RMF model.}\\label{tab:params2} \\begin{tabular}{crrrr} \\hline\\hline \\raisebox{-0.8ex}[0cm][0cm]{Isovector} & \\multicolumn{4}{c}{Set} \\\\ \\cline{2-5} \\raisebox{0.8ex}[0cm][0cm]{~~parameter~~} &I ~~~~~~~&II~~~~~~ &III~~~~~~ &IV~~~~~~ \\\\\\hline $g_{\\rho}$     &9.358 ~ ~ & 9.483 ~ ~ & 11.786 ~ ~  &13.687 ~ ~ \\\\ $\\eta_{\\rho}$  &0.190 ~ ~ & 0.390 ~ ~ & 4.490 ~ ~ &8.490 ~ ~ \\\\\\hline\\hline \\end{tabular}\\\\ \\end{table}  \\begin{figure*} \\centering \\mbox{\\epsfig{file=g2fac2.eps,height=11.0cm}} \\caption{ Isovector-vector channel adjustment in the E-RMF model. Effects of different combinations of $g_{\\rho}$ and $\\eta_{\\rho}$ on the symmetry energy of the nuclear matter are shown in the upper left panel, on the pressure and $M^*$ of the PNM are in the lower left and right panels, respectively, and on the proton fraction predictions in the upper right panel. Shaded region in the upper right panel corresponds to the proton fraction  threshold of the direct URCA process~\\cite{lati}.\\label{G2fac2}} \\end{figure*}  The symmetry energies of sets I--IV given in Table~\\ref{tab:params2} are shown in the upper left panel. Since it has been pointed out by Lattimer $\\it{et~al.}$~\\cite{lati}, that the crucial role of the proton fraction  value for the onset of the direct URCA process is to enhance the neutron star cooling rate~\\cite{baldo97}, we plot the corresponding proton fractions in the upper right panel. To estimate their effects on the EOS of neutron star, we plot the pressure as a function of energy of the pure neutron matter (PNM) as the dominant contributor in the EOS of a  neutron star, in the lower left panel. It is clear from the figure in the lower right panel that the adjustments of $g_{\\rho}$ and $\\eta_{\\rho}$ have no effect on the PNM effective mass ($M^*$).  In conclusion, for the E-RMF model, proton fraction predictions depend strongly on the behavior of the density dependent of $E_{\\rm sym}$ at high density. On the other hand, the predicted neutron star EOS does not drastically depend on the behavior of $E_{\\rm sym}$ at high density.  \\begin{table} \\centering \\caption{Isovector-vector channel adjustment in the Z271 parameter set of the Horowitz and Piekarewicz model~\\cite{Horowitz01}.}\\label{tab:params3} \\begin{tabular}{crrrr} \\hline \\hline \\raisebox{-0.8ex}[0cm][0cm]{Isovector} & \\multicolumn{4}{c}{Set} \\\\\\cline{2-5} \\raisebox{0.8ex}[0cm][0cm]{~~parameter~~} &I ~~~~~~&II~~~~~~ &III~~~~~~ &IV~~~~~~ \\\\\\hline $g_{\\rho}$     &9.498 ~ ~ & 9.672 ~ ~ & 11.506 ~ ~ &12.145 ~ ~ \\\\ $\\Lambda_{V}$  &0 ~ ~ & 0.01 ~ ~ & 0.03 ~ ~ &0.035 ~ ~ \\\\\\hline\\hline \\end{tabular}\\\\ \\end{table}  \\begin{figure*} \\centering \\mbox{\\epsfig{file=eospiek.eps,height=11.0cm}} \\caption{ Isovector-vector channel adjustment in the Horowitz and Piekarewicz model (Z271)~\\cite{Horowitz01}. Effects of different combinations of $g_{\\rho}$ and $\\Lambda_{V}$ on the symmetry energy of SNM are shown in the upper left panel, on pressure and $M^*$ of PNM in the lower left and right panels, respectively, and on proton fraction predictions in the upper right panel. Shaded region in the upper right panel corresponds to   the proton fraction threshold for the direct URCA process~\\cite{lati}.\\label{eospiek}} \\end{figure*}  A similar analysis for the standard RMF plus an additional isovector-vector nonlinear term model of Horowitz and Piekarewicz~\\cite{Horowitz01} with Z271* parameter set has been also performed. The various sets of coupling constant combinations are shown in Table~\\ref{tab:params3}, whereas their effects on matter properties are shown in Fig~\\ref{eospiek}. Similar conclusion is obtained both for proton fraction and $M^*$ in the PNM. Significant dependency in the EOS of this model on  $E_{\\rm sym}$ is found. The different trend in  $E_{\\rm sym}$ and EOS of this model compared to the E-RMF one originates from the different form of the isovector-vector non linear terms used in both models. \\begin{figure*} \\centering \\mbox{\\epsfig{file=nmfp.eps,height=11.0cm}} \\caption{ Effects  of the isovector-vector channel adjustment on NMFP predictions for G2, NL-3 and Z271 models. The upper left panel is for the E-RMF model (G2) and the upper right panel is for the Horowitz and Piekarewicz model (Z271). For comparison, the result of NL-3 is given in the upper-right panel by the dots form. The pressure as a function of the energy density and $M^*$ as a function of the nucleon-saturation densities ratio are given in the lower left and lower right panels, respectively.\\label{nmfp}} \\end{figure*}  To investigate which factor dominantly controls the behavior of the NMFP based on RMF models, we compare the EOS and NMFP of both models with the one of the standard RMF model, i.e., the NL-3 parameter set. The results are shown in Fig~\\ref{nmfp}. In the upper left panel, we show the effects from the variation of $E_{\\rm sym}$ in the G2* parameter sets (represented by the variation of  $g_{\\rho}$ and $\\eta_{\\rho}$ values) on the NMFP. Similarly, for the Z271* and NL-3 parameter sets (represented by the dots form),  the results are shown in the upper right panel. The EOS and $M^*$ in the PNM of the standard G2, Z271 and NL-3 parameter sets are displayed in the lower left and right panels, respectively. Different from NL-3, which has an anomalous behavior in its NMFP, it is found that   G2* and  Z271* parameter sets predict the NMFP trends, which do not change with the variation of  $E_{\\rm sym}$ (in both models, the anomalous behavior in the NMFP does not appear in every parameter set even though they have different E$_{\\rm sym}$ values). This means that the appearance of the anomalous behavior in the NMFP seems to be insensitive to the value of the proton fraction. The NL-3 parameter set has PNM with a stiff EOS and a relatively small $M^*$ value but, on the contrary, Z271 and G2 have PNM with a soft EOS and large $M^*$ value at high density. This fact gives an indication that a soft EOS and a normal behavior of the NMFP are mostly determined by the relatively large $M^*$ value at high density. On the other hand, finite nuclei calculations using the standard RMF model~\\cite{pg,ring,serot,anto05} inform us that acceptable shell structure predictions in finite nuclei regions require a small $M^*$ value ($\\sim$ 0.6 $M$) in the saturation density which is fulfilled by G2~\\cite{sil} and NL-3~\\cite{pg,ring,serot,anto05} parameter sets.  \\begin{table} \\centering \\caption {$M^*$ variations in the NL-Z parameter set of the standard RMF model.~\\cite{anto05}}\\label{tab:params4} \\begin{tabular}{crrr} \\hline\\hline Parameter &NL-Z~~ & P-070~~& P-080~~ \\\\\\hline $g_S/(4 \\pi)$ &$0.801$~~& 0.673~~& 0.575~~ \\\\ $g_V/(4 \\pi)$ &$1.028$~~& 0.817~~& 0.632~~ \\\\ $\\kappa_3$     &$2.084$~~& 2.463~~& 2.899~~ \\\\ $\\kappa_4$     &$-8.804$~~& $-7.595$~~& 12.779~~ \\\\\\hline\\hline \\end{tabular}\\\\ \\end{table}  \\begin{figure*} \\centering \\mbox{\\epsfig{file=mstrmfp.eps,height=11.0cm}} \\caption{Effects of the $M^*$ variations on the proton fraction and on the  NMFP (upper left and upper right panels), and on the EOS (lower left panel).  Shaded region in the upper left panel corresponds to a threshold of the proton fraction for the direct URCA process~\\cite{lati}.  Variations of the $M^*$ as a function of  the nucleon-saturation densities ratio is shown in the lower right panel.  \\label{mstrmfp}} \\end{figure*}  Actually, the above results could be more clearly interpreted  by looking at the effects of different $M^*$ on the NMFP in the same model. The difference in each parameter set is only in the values of some adjusted coupling constants, while they should have acceptable SNM and PNM predictions at saturation density (more details about matter predictions of these models can be seen in Ref.~\\cite{anto05}), and for our purpose here, the proton fraction predictions should be similar. The parameter sets of Ref.~\\cite{anto05} are suitable for this task. The coupling constants variations of the models are shown in Table~\\ref{tab:params4}, whereas the results are shown in Fig.~\\ref{mstrmfp}. The effects of the $M^*$ variation on the proton fraction are shown in the upper left panel, while the effects on the NMFP  and EOS of PNM are shown in the upper right and lower left panels, respectively. Clearly, if $M^*$ becomes too low then the corresponding EOS becomes too stiff and the anomalous behavior in the NMFP appears. The P-070 and P-080 parameter sets have unacceptable shell structure predictions in some nuclei~\\cite{anto05} since these parameter sets have a too large $M^*$ in the saturation density which, as a consequence, leads to a too narrow spin-orbit splitting prediction. Recently, ``FSU GOLD'' parameter set has been introduced by Todd-Rutel and Piekarewicz~\\cite{todd} which yields a soft EOS, while still accurately reproducing experimental data of binding energies and charge radii of some magic nuclei and also centroid energies for breathing mode of $^{208}\\rm Pb$ and $^{90}\\rm Zr$. Unfortunately the shell structure prediction of this parameter set is not  reported in that paper. Therefore, before drawing any further conclusion, we should wait for their full calculation, including the predicted shell structure properties of some magic nuclei.  In conclusion, these results confirm previous findings~\\cite{sil,arumu} about the wide range of applications of the E-RMF model. In our view, the reason comes from the fact that the E-RMF model has a relatively small $M^*$ ($\\sim 0.6~M$) in saturation density (demanding feature for finite nuclei) but a relatively large $M^*$ at high density (demanding feature for the neutron star). Extra nonlinear and tensor terms of this model compared to the standard one seem to be the source of this behavior (see Table \\ref{tab:params1}). From the possibility that the density dependent  $E_{\\rm sym}$  can be adjusted, the claim that RMF models predict relatively lower threshold densities for the direct URCA process and this fact can be considered as a weak point of the models~\\cite{blaschke,kolo,migdal}, now can be re-explored. A precise density dependent $E_{\\rm sym}$ at high density experimentally determined and/or extracted from the properties of the neutron star are needed in this case. \\begin{table} \\centering \\caption {Effects of the $\\delta$ meson on the G2 parameter set of the E-RMF model. Case $\\eta_{\\rho}=0.39$.}\\label{tab:params5} \\begin{tabular}{crrr} \\hline\\hline  \\raisebox{-0.8ex}[0cm][0cm]{Isovector} & \\multicolumn{3}{c}{Set} \\\\\\cline{2-4} \\raisebox{0.8ex}[0cm][0cm]{~~parameter~~} &I ~~~~~~~&II~~~~~~ &III~~~~~~ \\\\\\hline $g_{\\rho}$ &9.483 ~ ~ ~& 12.313~ ~ ~& 15.937 ~ ~ ~ \\\\ $g_{\\delta}$ &0 ~ ~ ~& 5.026~ ~ ~& 7.540~ ~ ~ \\\\\\hline\\hline \\end{tabular}\\\\ \\end{table}  \\begin{table} \\centering \\caption {Effects of the $\\delta$ meson in the G2* parameter sets of the E-RMF model. Case $\\eta_{\\rho}=4.49$.}\\label{tab:params6} \\begin{tabular}{crrr} \\hline\\hline  \\raisebox{-0.8ex}[0cm][0cm]{Isovector} & \\multicolumn{3}{c}{Set} \\\\\\cline{2-4} \\raisebox{0.8ex}[0cm][0cm]{~~parameter~~} &I ~~~~~~~&II~~~~~~ &III~~~~~~ \\\\\\hline $g_{\\rho}$ &11.786 ~ ~& 15.304~ ~ ~& 18.784 ~ ~ ~\\\\ $g_{\\delta}$ & 0 ~ ~& 5.026~ ~ ~& 7.540~ ~ ~ \\\\\\hline\\hline \\end{tabular}\\\\ \\end{table}    \\begin{figure*} \\centering \\mbox{\\epsfig{file=delfrac.eps,height=11.0cm}} \\caption{ Effects of the $\\delta$ meson addition in the E-RMF model for $\\eta_{\\rho}=0.39$ on the symmetry energy and proton fraction (upper left and right panels) and on the EOS and $M^*$ in the PNM ( lower left and right panels).  Shaded region in the upper right panel corresponds to  the proton fraction  threshold for the direct URCA process~\\cite{lati}.\\label{delfrac}} \\end{figure*}  \\begin{figure*} \\centering \\mbox{\\epsfig{file=delfrac2.eps,height=11.0cm}} \\caption{ Same as in Fig.~\\ref{delfrac}, but for  $\\eta_{\\rho}=4.49$.\\label{delfrac2}} \\end{figure*}  To study the  effects of $\\delta$ meson on the E-RMF model, we start with the standard G2 with  $\\eta_{\\rho}=0.39$ and generate different  $g_{\\rho}$ and $g_{\\delta}$ parameters combinations but we keep the same $E_{\\rm sym}=26.57$ MeV  at $k_F=1.172$ fm$^{-1}$ (note: there would be  no change in the conclusion if we took  $E_{\\rm sym}=24.1$ MeV  at $k_F=1.14$ fm$^{-1}$). The combinations of the coupling constants of the models can be seen in Table~\\ref{tab:params5}. The matter properties predictions are shown in Fig.~\\ref{delfrac}. It is clearly seen that the presence of the $\\delta$ meson results in a higher $E_{\\rm sym}$ at high density. This fact leads to a higher proton fraction. But, in the region $\\rho_B$=(1.5-2)$\\rho_0$, the difference  between the $E_{\\rm sym}$ value of the E-RMF plus a $\\delta$ and that without a $\\delta$ meson is not so significant. The presence of the $\\delta$ meson also makes the  PNM EOS stiffer and reduces the value of the PNM $M^*$. The reduction magnitude depends on the magnitude of $g_{\\delta}$. A similar trend is also found  in the  case of $\\eta_{\\rho}=4.49$. The combinations of $g_{\\rho}$ and $g_{\\delta}$ coupling constant are shown in Table~\\ref{tab:params6} and the corresponding results can be seen in Fig.~\\ref{delfrac2}.  \\begin{figure*} \\centering \\mbox{\\epsfig{file=nmfpdel.eps,height=5.5cm}} \\caption{  Effects of the $\\delta$ meson addition in the E-RMF model on the corresponding NMFP prediction. Left panel is for $\\eta_{\\rho}=0.39$, while right panel is for  $\\eta_{\\rho}=4.49$.\\label{nmfpdel}} \\end{figure*}  In Fig.~\\ref{nmfpdel}, it can be seen that the presence of the $\\delta$ meson removes the anomalous behavior in the predicted NMFP. The effect appears more pronounce in the case  of $\\eta_{\\rho}=4.49$, rather than $\\eta_{\\rho}=0.39$. This fact is clearly depicted  in Fig.~\\ref{nmfpdel}.  \\begin{figure*} \\centering \\mbox{\\epsfig{file=eos1.eps,height=11.0cm}} \\caption{ Energy per nucleon ($E/A$) of the SNM (upper left) and PNM (upper right) of G2, Z271, NL3 and $G2^*+\\delta$ parameter sets. The corresponding pressures are given in the lower left and lower right panels. For comparison, we also  show the results from variational calculation of Akmal {\\it et al.}~\\cite{akmal98}, DBHF calculation of Li {\\it et al.}~\\cite{li92}, and BHF with AV14 potential plus 3BF of Baldo {\\it et ~al.}~\\cite{baldo97}. Shaded regions correspond to experimental data from Danielewicz {\\it et al.} \\cite{daniel02}.\\label{eos1}} \\end{figure*}  In Fig.~\\ref{eos1}, we show the $E/A$ ratio and the pressure, for the SNM as well as the PNM of the four RMF parameter sets. NL-3 is a parameter set with good predictions for observables of finite nuclei and  has a stiff EOS at high density. Z271 is a parameter set that is specially constructed  for the neutron star and has a soft EOS at high density. G2 is a parameter set with acceptable predictions for observables of finite nuclei and has a relatively soft EOS at high density. Clearly, parameter sets with soft EOS are consistent with experimental data of Danielewicz $\\it{et~al.}$~\\cite{daniel02} and close to the results of the variational calculation by Akmal  $\\it{et~al.}$~\\cite{akmal98}, BHF calculation with AV14 potential plus 3BF of Baldo $et~al.$~\\cite{baldo97}, and DBHF calculation of Li  $\\it{et~al.}$~\\cite{li92}. It seems also from the results of $G2^*+ \\delta$, that the enhancement in the isovector channel of G2 shifts the $E/A$ ratio and pressure predictions of that parameter set closer to the result of variational calculation from Akmal  $\\it{et~al.}$~\\cite{akmal98}."
    },
    "0512/astro-ph0512224_arXiv.txt": {
        "abstract": "We investigate the general relativistic dynamics of Robertson-Walker models with a non-linear equation of state (EoS), focusing on  the  quadratic case $P=P_0+\\alpha \\rho +\\beta \\rho^2$. This may be taken to represent the Taylor expansion of any arbitrary barotropic EoS, $P(\\rho)$. With the right combination of $P_0$, $\\alpha$ and $\\beta$, it  serves as a simple phenomenological model for dark energy, or even unified dark matter. Indeed we show that this simple model for the EoS can produce a large variety of qualitatively different dynamical behaviors that we classify using dynamical systems theory. An almost universal feature is that accelerated expansion phases are mostly natural for these non-linear EoS's. These are often asymptotically de~Sitter thanks to the appearance of an {\\it effective cosmological constant}. Other interesting possibilities that arise from the quadratic EoS  are closed models that can oscillate with no singularity, models that bounce between infinite contraction/expansion and models which evolve from a phantom phase, asymptotically approaching a de Sitter phase instead of evolving to a ``Big Rip''. In a second paper we investigate the effects of the quadratic EoS in inhomogeneous and anisotropic models, focusing in particular on singularities. ",
        "introduction": "Model building in cosmology requires two main ingredients: a theory of gravity and a description of the matter content of the universe.  In general relativity (GR) the gravity sector of the theory is completely fixed, there are no free parameters. The matter sector is represented in the field equations by the energy-momentum tensor, and for a fluid the further specification of an equation of state (EoS) is required. Apart from scalar fields, typical cosmological fluids such as radiation or cold dark matter (CDM) are represented by a {\\it linear} EoS, $P={\\it w}\\rho $.  The combination of cosmic microwave background radiation (CMBR)~\\cite{CMBR,spergel}, large scale structure (LSS)~\\cite{LSS} and supernova type Ia (SNIa)~\\cite{SNI} observations provides support for a flat universe presently dominated by a component, dubbed in general ``dark energy'',  causing an accelerated expansion. The simplest form of dark energy is an {\\it ad hoc} cosmological constant $\\Lambda$ term in the field equations, what Einstein called his ``biggest blunder''. However,  although  the standard  $\\Lambda$CDM ``concordance'' model provides a rather robust framework for the interpretation of present observations (see e.g.~\\cite{spergel,turner}), it requires a $\\Lambda$ term that is at odds by many order of magnitudes with theoretical predictions~\\cite{weinberg}. This has prompted theorists to explore possible dark energy sources for the acceleration that go beyond the standard but unsatisfactory $\\Lambda$. With various motivations, many authors have attempted to describe dark energy as quintessence, {\\it k}-essence or a ghost field, i.e. with scalar fields with various properties. There have also been attempts to describe dark energy by a fluid with a specific non-linear EoS like the Chaplygin gas~\\cite{KMP}, generalized Chaplygin gas~\\cite{GCG}, van der Waals fluid~\\cite{GK}, wet dark fluid~\\cite{HN} and other specific gas EoS's~\\cite{CTTC}. Recently, various ``phantom models'' (${\\it w}=P/\\rho<-1 $) have also been considered~\\cite{RRC, sahni}. More simply, but also with a higher degree of generality, many authors have focused on phenomenological models where dark energy is parameterized by assuming a $w=P/\\rho=w(a)$, where $a=a(t)$ is the expansion scale factor (see e.g.~\\cite{bruce,will}).   Another possibility is to advocate a modified theory of gravity. At high energies, modification of gravity beyond general relativity could come from extra dimensions, as required in string theory. In the brane world~\\cite{SMS,BMW,DL,RM} scenario the extra dimensions produce a term quadratic in the energy density in the effective 4-dimensional energy-momentum tensor. Under the reasonable assumption of neglecting 5-dimensional Weyl tensor contributions on the brane, this quadratic term has the very interesting effect of suppressing anisotropy at early enough times. In the case of a Bianchi I brane-world cosmology containing a scalar field with a large kinetic term the initial expansion is quasi-isotropic~\\cite{MSS}. Under the same assumptions, Bianchi I and Bianchi V brane-world cosmological models containing standard cosmological fluids with linear EoS  also behave in a similar fashion\\footnote{This only requires $P/ \\rho={\\it w}>0 $, as opposed to ${\\it w}>1 $ in the GR case. In the case of ekpyrotic/cyclic and pre-big bang models the initial expansion is only isotropic if ${\\it w}>1 $ as in the case of GR ~\\cite{EWST}.}~\\cite{CS}, and the same remains true for more general homogeneous models~\\cite{coley1,coley2} and even some inhomogeneous exact solutions~\\cite{coley3}. Finally, within the limitations of a perturbative treatment, the quadratic-term-dominated isotropic brane-world models have been shown to be local past attractors in the larger phase space of inhomogeneous and anisotropic models~\\cite{DGBC, GDCB}. More precisely, again assuming that the 5-d Weyl tensor contribution to the brane can be neglected, perturbations of the isotropic models decay in the past. Thus in the brane scenario the observed high isotropy of the universe is the natural outcome of {\\it generic initial conditions}, unlike in GR where in general cosmological models with a standard energy momentum tensor are highly anisotropic in the past (see e.g. \\cite{LL}).  Recently it has been shown that loop quantum gravity corrections result in a modified Friedmann equation~\\cite{KV}, with the modification appearing as a negative term which is quadratic in the energy density. Further motivation for considering a quadratic equation of state comes from recent studies of {\\it k}-essence fields as unified dark matter (UDM) models\\footnote{These attempt to provide a unified model for both the dark matter and the dark energy components necessary to make sense of observations.}~\\cite{GH,RS}. The general {\\it k}-essence field can be described by a fluid with a closed-form barotropic equation of state. The UDM fluid discussed in~\\cite{GH} has a non-linear EoS of the form $P\\propto \\rho^2$ at late times. More recently, it has been shown that any purely kinetic {\\it k}-essence field can be interpreted as an isentropic  perfect fluid with an EoS of the form $P=P(\\rho)$~\\cite{DTF}. Also, low energy dynamics of the Higgs phase for gravity have been shown to be equivalent to the irrotational flow of a perfect fluid with equation of state $P=\\rho^2$~\\cite{ACLMW}.  Given the isotropizing effect that the quadratic energy density term has at early times in the brane scenario this then prompts the question: can a term quadratic in the energy density have the same effect in general relativity. This question is non-trivial as the form of the equations in the two cases is quite different. On the brane,  for a given EoS the effective 4-dimensional Friedmann and Raychaudhuri equations are modified, while the continuity equation is identical to that of  GR. With the introduction of a quadratic EoS in GR, the Friedman equation remains the same, while the continuity and Raychaudhuri equations are modified\\footnote{With Respect to the case of the same EoS with vanishing quadratic term.}.  Taking into account this question (to be explored in detail in Paper II~\\cite{AB}), the diverse motivations  for a quadratic energy density term mentioned above and with the dark energy problem in mind, in this paper we explore the  GR dynamics of homogeneous isotropic Robertson-Walker models with a quadratic EoS, $P=P_0+\\alpha \\rho +\\beta \\rho^2$. This is the simplest model we can consider without making any more specific assumptions on the EoS~\\cite{MV}. It represents the first terms of the Taylor expansion of {\\it  any} EoS function $P=P(\\rho)$ about $\\rho=0$. It can also be taken to represent (after re-grouping of terms) the Taylor expansion about the present energy density $\\rho_0$, see \\cite{MV}. In this sense therefore the out-coming dynamics is very general. Indeed it turns out  that this simple model can produce a large variety of qualitatively different dynamical behaviors that we classify using dynamical systems theory~\\cite{WE, AP}. An outcome of our analysis is that accelerated expansion phases are mostly natural for non-linear EoS's. These are {\\it in general} asymptotically de Sitter thanks to the appearance of an {\\it effective cosmological constant}. This suggests that an EoS with the right combination of $P_0$, $\\alpha$ and $\\beta$ may provide a good and simple phenomenological model for  UDM, or at least for a dark energy component. Other interesting possibilities that arise from the quadratic EoS  are closed models that can oscillate with no singularity, models that bounce between infinite contraction/expansion and models which evolve from a phantom phase, asymptotically approaching a de Sitter phase instead of evolving to a ``big rip'' or other pathological future states \\cite{RRC,BLJM,NOT}.  As mentioned before, the question of the dynamical effects  the quadratic energy density term  has on the anisotropy  in GR is explored in Paper II~\\cite{AB}. There we analyze Bianchi I and V models with the EoS $P=\\alpha\\rho +\\beta \\rho^2$, as well as perturbations of the isotropic past attractor of those  models that are singular in the past. We anticipate that Bianchi I and V non-phantom models with $\\beta>0$ have an isotropic singularity, i.e. they are asymptotic in the past to a certain isotropic model, and that perturbations of this model decay in the past. Phantom anisotropic models with $\\beta>0$ are necessarily asymptotically de~Sitter in the future, but the shear anisotropy  dominates in the past. For $\\beta<0$ all models are anisotropic in the past, while their specific future evolution depends on the value of $\\alpha$.   The paper is organized as follows. In section~\\ref{sec2} we outline the setup and the three main cases we will investigate. In section~\\ref{sec3}, we study the dynamics of isotropic cosmological models in the high energy limit (neglecting the $P_0$ term). We find the critical points, their stability nature and the occurrence of bifurcations of the dynamical system. In section~\\ref{sec4}, we consider the low energy limit (neglecting the $\\rho^2$ term). The full system is then analyzed in section~\\ref{sec5}, showing the qualitatively different behavior with respect to the previous cases. We then finish with some concluding remarks and an outline of work in progress in section~\\ref{sec6}. Units are such that $8\\pi G/c^4=1$. ",
        "conclusions": "\\label{sec6}   In this paper we have systematically studied the dynamics of homogeneous and isotropic cosmological models containing a fluid with a quadratic EoS. This has it's own specific interest (see Section I for a variety of motivations) and serves as a simple example of more general EoS's. It can also be taken to represent the truncated Taylor expansion of any barotropic EoS, and as such it serves (with the right choice of parameters) as a useful phenomenological model for dark energy, or even UDM. Indeed, we have shown the dynamics to be very different and much richer than the standard linear EoS case, finding  that an almost generic feature of the evolution is the existence of an accelerated phase, most often asymptotically de Sitter,  thanks to the appearance of an {\\it effective cosmological constant}. Of course to properly build physical cosmological models would require to consider the quadratic EoS for dark energy or UDM together with standard matter and radiation. Our analysis was aimed instead to derive and classify the large variety of different dynamical effects that the quadratic EoS fluid has when is the dominant component. In this respect, it should be noticed that a positive quadratic term in the EoS allows, in presence of another fluid such as radiation, equi-density between the two fluid to occur twice, i.e. the quadratic EoS fluid can be dominant at early and late times, and subdominant in an intermediate era.  In Section II we have made some general remarks, mostly based on conservation of energy only and as such valid independently of any specific theory of gravity. We have also given the various possible functional forms of the energy density as a function of the scale factor, $\\rho(a)$, and listed the many subcases, grouped in three main cases, what we call: {\\it i)}  the high energy models (no constant $P_o$ term); {\\it ii)} the low energy  affine EoS with no quadratic term; {\\it iii)} the complete quadratic EoS.  The quadratic term in the EoS affects the high energy behavior as expected but can additionally affect the dynamics at relatively low energies. First, in Section III, we have concentrated on the high energy models. The specific choice of parameters fixes the behavior of the fluid, it can behave in a phantom or standard manner. In the case of phantom behavior, $\\rho$ can tend to zero at early times and either tend to an effective cosmological constant ({\\bf C1}) or a Type III singularity ({\\bf A2}) at late times. Alternatively $\\rho$ can also tend to an effective cosmological constant in the past ({\\bf B2}) and a Type III singularity at late times. When the fluid behaves in a standard manner, it can tend to a Type III singularity at early times, with $\\rho$ either tending to zero ({\\bf A1}) or to an effective cosmological constant ({\\bf B1}) at late times. The fluid can also behave as an effective cosmological constant at early times with $\\rho$ decaying away at late times ({\\bf C2}). The effective cosmological constant allows for the existence of  generalized Einstein static ($E$) and flat de Sitter fixed ($dS_{\\pm}$) points which modify the late time behavior. The main new feature is the existence of models which evolve from a Type III singularity and asymptotically approach a flat de Sitter model ($dS_{+}$).  Of specific interest are the closed models of this type, which can also evolve through an intermediate  loitering phase.  Neglecting the quadratic term, in Section IV we have considered the low energy models with affine EoS. As expected, the constant term in the quadratic EoS affects the relatively low energy behavior. It can result in a variety of qualitatively different dynamics with respect to those of the linear EoS case. Again, the fluid can have a phantom or standard behavior. When the fluid behaves in a phantom manner, $\\rho$ can tend to an effective cosmological constant ({\\bf F2}), or  can tend to a Type I (``Big Rip\") singularity ({\\bf D2}) at late times. Alternatively, $\\rho$ can also tend to an effective cosmological constant in the past and a Big Rip in the future({\\bf E1}). When the fluid behaves in a standard manner, we recover the linear EoS at early times and $\\rho$ can either tend to zero ({\\bf D1}) or to an effective cosmological constant ({\\bf E2}) at late times. The fluid can also behave as an effective cosmological constant at early times, with $\\rho$ decaying away at late times ({\\bf F1}). The effective cosmological constant allows for the existence of new fixed points($E$ and $dS_{\\pm}$). Comparing with standard linear EoS cosmology, the most interesting differences are new closed models which oscillate indefinitely and new closed models which exhibit phantom behavior which do not terminate in a ``Big Rip\", but asymptotically approach an expanding flat de Sitter model (flat and closed models where the fluid behaves as case {\\bf F2}).  When we study the dynamics of the system with the complete quadratic EoS, Section V, we see the appearance of new fixed points representing generalized Einstein and de Sitter models which are not present in the high/low energy systems. The various models of the simplified systems are present in the full system (but with differing $\\rho(a)$), but there are also models with qualitatively new behavior. As with the previous cases, in the case of phantom behavior, $\\rho$ can tend to zero at early times and either tend to an effective cosmological constant ({\\bf H3} and {\\bf I4}) or a Type III singularity ({\\bf G2}) at late times. Alternatively $\\rho$ can also tend to an effective cosmological constant in the past ({\\bf H4} and {\\bf I6}) and a Type III singularity at late times. Finally, in the phantom case $\\rho$ can also tend to an effective cosmological constant both in the past and future ({\\bf I2}). In the case of standard behavior the fluid can tend to a Type III singularity at early times, with $\\rho$ either tending to zero ({\\bf G1}) or to an effective cosmological constant ({\\bf H2} and {\\bf I3}) at late times. The fluid can also behave as an effective cosmological constant at early times with $\\rho$ decaying away at late times ({\\bf H1} and {\\bf I1}). Finally, in the standard fluid case $\\rho$ can also tend to an effective cosmological constant both in the past and future ({\\bf I5}). There are  models which evolve from a Type III singularity, reach a maximum $a$ (minimum $x$) and then evolve to Type III singularity. These also enter a loitering phase before and after the turn around point. We also see bounce models which enter a loitering phase and asymptotically tend to generalized expanding (contracting) de Sitter models at late (early) times.  Of specific interest are models which evolve from a Type III singularity as opposed to the standard ``Big Bang\" ({\\bf A1, B1}). The simplest models of this type correspond to the high energy EoS with a positive quadratic term (is possible to recover standard behavior at late times). For these models the positive quadratic energy density term has the potential to force the initial singularity to be isotropic. The effects of such a fluid on anisotropic Bianchi I and V models is investigated in Paper II~\\cite{AB}. This is  achieved by carrying out a dynamical systems analysis of these models. Additionally, using a linearized perturbative treatment we study the behavior of inhomogeneous and anisotropic perturbations at the singularity. The singularity is itself represented by an isotropic model and, If the perturbations of the latter decay in the past,  this model represents the  local past attractor in the larger phase space of inhomogeneous and anisotropic models (within the validity of the  perturbative treatment). This would mean that  in  inhomogeneous anisotropic models  with a  positive non-linear term (at least quadratic)   in the EoS isotropy is   a natural outcome of {\\it generic initial conditions}, unlike in the standard linear EoS case where generic cosmological models are, in GR, highly anisotropic in the past."
    },
    "0512/astro-ph0512012_arXiv.txt": {
        "abstract": "This paper presents a brief overview of the accomplishments of the Chandra satellite that are shedding light on the origin of high energy particles in astrophysical sources, with the emphasis on clusters of galaxies. It also discusses the prospects for the new data to be collected with instruments recently launched - such as Suzaku - or those to be deployed in the near future, and this includes GLAST and NuSTAR. ",
        "introduction": "The last several years can be truly called the ``golden era'' of high energy astrophysics.  At the time of this meeting, Chandra and XMM-Newton were conducting imaging observations in the soft X-ray band;  RXTE was measuring the timing properties of variable celestial X-ray sources;  Integral covered the hard X-ray and soft gamma-ray regime;  and Swift, just launched, was effectively discovering and measuring properties of gamma-ray bursts.  All this resulted in tremendous advancements of our understanding of sources of high energy radiation - shedding light on the physical processes responsible for the particle acceleration in the Universe.  This paper focuses on selected results derived from Chandra data, but also presents the prospects for the near future - and this includes the recently launched Suzaku, and the approved satellite missions GLAST and NuSTAR. ",
        "conclusions": ""
    },
    "0512/astro-ph0512538_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:introduction} Observations of an astrophysical signal in the sky are usually corrupted by some level of contamination (called {\\it noise} or {\\it background}), due to other astrophysical emissions and/or to the detector itself. A common situation is that the signals of interest are spatially well-localised, i.e. each of them covers only a small fraction of the image, but we do not know a priori its position and/or its amplitude. Some examples are the detection of extragalactic sources in cosmic microwave background (CMB) observations (see Fig.~\\ref{fig:example_planck}), the identification of local features (emission or absorption lines) in noisy one-dimensional spectra or the detection of objects in X-ray images. It is clear that our ability to extract all the useful information from the image will critically depend on our capacity to disentangle the signal(s) of interest from the background. \\begin{figure} \\centering \\includegraphics[angle=270,width=8cm]{example_planck.ps} \\caption{Simulation of the 44 GHz Planck frequency channel in a small patch of the sky, containing CMB, Galactic foregrounds, extragalactic point sources and instrumental noise. The point sources can be seen as localised objects embedded in the background. The Planck Mission~\\cite{bb:planck} is a satellite of the European Space Agency to be launched in 2007 that will provide with multifrequency observations of the whole sky with unprecedented resolution and sensitivity.} \\label{fig:example_planck} \\end{figure}  The process to detect a localised signal in a given image usually involves three different steps, which are not necessarily independent:  1.- Processing: some processing of the data (commonly linear filtering) is usually performed in order to amplify the searched signal over the background. This is an important step because in many cases the signals are relatively weak with respect to the background and it becomes very difficult to detect them in the original image. This is illustrated in Figure~\\ref{fig:white_noise}: the top panel shows a simulation of white noise where a source with a Gaussian profile has been added in the centre of the map; the bottom panel gives the same simulation after filtering with the so-called matched filter. It becomes apparent that the source was hidden in the original image whereas it has been enhanced in the filtered image. \\begin{figure} \\centering \\includegraphics[height=\\hsize]{white_noise.ps} \\caption{This example illustrates the importance of filtering. A source with a Gaussian profile has been placed in the centre of the image in a background of white noise. The source can not be distinguished in the original image (left panel), however, after filtering (right panel), the source is enhanced over the background fluctuations.} \\label{fig:white_noise} \\end{figure}  2.- Detection: we need a detection criterion, the {\\it detector}, to decide if some structure in the image is actually a real signal or if it is due to the background. A very simple and widely used detector in Astronomy is {\\it thresholding}: if the intensity of the image is above a given value (e.g. 5$\\sigma$, where $\\sigma$ is the dispersion of the map), a detection of the signal is accepted, otherwise one assumes that only background is present. In the example of Fig.~\\ref{fig:white_noise}, we see that several peaks appear in the filtered image (right panel) but only one of them is above the considered threshold $\\nu \\sigma$. Therefore, in this case, we would accept only the highest peak as a true signal. Note that thresholding uses only the intensity of the data to make the decision, however other useful information could also be included in order to improve the detector (e.g. curvature, size, etc.).  3.- Estimation: a procedure must be established to estimate the parameters (amplitude, size, position...) characterising the detected signal. For instance, a simple possibility is to estimate the required parameters by fitting the signal to its theoretical profile.  The aim of these lecture notes is to present the problem of the extraction of localised signals (compact sources) in the context of CMB Astronomy and to review some of the methods developed to deal with it. In section~\\ref{sec:microwave_sky}, we outline the problem of component separation in CMB observations. Section~\\ref{sec:sources} reviews some of the techniques developed for extraction of point sources, including, among others, the matched filter and the Mexican Hat Wavelet. Sections~\\ref{sec:tsz} and~\\ref{sec:ksz} deal with the extraction of the thermal and kinematic Sunyaev-Zeldovich effects in multifrequency microwave observations, respectively. Section~\\ref{sec:statistical} briefly discusses some techniques for the extraction of statistical information from undetected sources. Finally, in section~\\ref{sec:conclusions} we present our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  The development of techniques for the extraction of compact sources in CMB observations has become a very relevant and active topic. This is due to the necessity of cleaning the CMB maps from astrophysical contaminants that would impair our ability to extract all the valuable information encoded in the cosmological signal but also because the recovered catalogues of point sources and/or SZ clusters would contain themselves extremely relevant astrophysical and cosmological information.  An important effort has been done in the last years towards the development of more powerful and sophisticated tools to extract compact sources. Many of them have been tested on simulated Planck observations, showing their potentiality. However some important work still remains to be done. First of all, in some cases, quite ideal conditions have been assumed. For instance, it is commonly assumed that the cluster profile is known but, in general, this will not be the case for real data. Other methods have been applied to simulations that do not include foreground emissions. Therefore, these and other problematics -- beam asymmetry, extension to the sphere, relativistic effects, anisotropic noise, etc. -- present on real data should be taken into account to establish the true performance of the developed methods. Also, the methods do not always use all the available information present in the data. For instance, if multifrequency observations are available, it would be useful to include this multifrequency information in the detection of extragalactic point sources even if they do not follow a simple well-known spectral law. The final and most important step would be to apply these techniques to real CMB data (e.g. WMAP) as they become available.  We would also like to point out that many polarisation experiments are currently planned (or already operating) which will provide with a wealth of information about our universe \\cite{bb:cha}. Given the weakness of the cosmological signal in polarisation and the current lack of knowledge regarding the foreground emissions, a careful process of cleaning of the CMB polarisation maps is even more critical than for the intensity case. However no techniques have been yet specifically developed to extract compact sources from polarisation CMB observations. Therefore it is crucial to extend some of the current methods -- or to develop new ones -- to deal with this type of maps.  Finally, a very critical issue is to assess which is the impact of possible residuals left in the CMB data after applying these techniques~\\cite{bb:bar05}. In particular, it is very important to control the effect of undetected sources, or even possible artifacts introduced in the image after subtracting the signals, on the estimation of the power spectrum of the CMB. In addition, this process should not modify the underlying CMB temperature distribution, since it would impair our ability to perform Gaussianity analyses of the CMB (or even lead us to wrong conclusions), which are of great importance to learn about the structure formation of our universe."
    },
    "0512/astro-ph0512362_arXiv.txt": {
        "abstract": "{We present a JHK$_s$L survey of the massive star forming region RCW 57 (NGC 3576) based on L-band data at 3.5 $\\mu$m taken with SPIREX (South Pole Infrared Explorer), and 2MASS JHK$_s$ data at 1.25-2.2 $\\mu$m. This is the second of two papers, the first one concerning a similar JHK$_s$L survey of 30 Doradus. Colour-colour and colour-magnitude diagrams are used to detect sources with infrared excess. This excess emission is interpreted as coming from circumstellar disks, and hence gives the cluster disk fraction (CDF). Based on the CDF and the age of RCW 57, it is possible to draw conclusions on the formation and early evolution of massive stars. The infrared excess is detected by comparing the locations of sources in JHK$_s$L colour-colour and L vs. (K$_s$-L) colour-magnitude diagrams to the reddening band due to interstellar extinction. A total of 251 sources were detected. More than 50\\% of the 209 sources included in the diagrams have an infrared excess. Comparison with other JHK$_s$L surveys, including the results on 30 Doradus from the first paper, support a very high initial disk fraction ($>$80\\%) even for massive stars, although there is an indication of a possible faster evolution of circumstellar disks around high mass stars. 33 sources only found in the L-band indicate the presence of heavily embedded, massive Class I protostars. We also report the detection of diffuse PAHs emission throughout the RCW 57 region. ",
        "introduction": "\\label{intro} \\subsection{IR-excess as a measure of circumstellar disks} \\label{intro1} This paper is the second of two papers using IR-excess in the JHK$_s$L plane (1.2 - 3.5 $\\mu$m) to measure the fraction of sources with circumstellar disks in high mass star forming regions. The first paper concerned 30 Doradus in the LMC (Maercker \\& Burton \\cite{maerckerburton}, from now referred to as Paper I). IR-excess can be detected using near infrared colour-colour diagrams by comparing the position of sources relative to the reddening vectors due to interstellar extinction. The excess radiation above that of a blackbody can be explained by models of circumstellar disks around a young stellar object (eg. Lada \\& Adams \\cite{ladaadams}). Although this excess radiation can be detected using JHK$_s$ data (1.2 - 2.2 $\\mu$m) alone, the nature of IR-excess is not always clear. JHK$_s$L observations give a larger separation to the IR-excess sources in colour-colour diagrams, whereas JHK$_s$ observations tend to underestimate the fraction of stars with IR-excess. On account of the difficulties of ground based observations at longer wavelengths, the L-band (3.5 $\\mu$m) proves to be the best wavelength for detecting circumstellar disks, although longer wavelengths are preferable, if space observations are available. Kenyon \\& Hartmann (\\cite{kenyonhartmann}) show the advantage of (K$_s$-L) as a measure of IR-excess by comparing the frequency distributions in young stellar clusters for (H-K$_s$) and (K$_s$-L). Whereas the (H-K$_s$) distribution has one clear peak at (H-K$_s$)$\\sim$0.2-0.4 and a long tail, the (K$_s$-L) distribution has a clear second peak at (K$_s$-L)$\\sim$0.8-1.0 made up mostly of class II sources with optically thick, circumstellar disks.  \\subsection{RCW 57} \\label{rcw57} This young massive star forming region, also known as NGC 3576, is one of  the brightest HII regions in the infrared in our Galaxy. The kinematic distance is 3.0 $\\pm$ 0.3 kpc, adopted from De Pree et al. (\\cite{depreeetal}). An asymmetrical structure in the region can be seen in the 21 cm map by Retallack \\& Goss (\\cite{retallackgoss}), which extends to the northeast but has a sharp cut off in the southwest. The spectral energy distributions of five objects detected using a 10 $\\mu$m map (Frogel \\& Persson~\\cite{frogelpersson}) suggest that these are protostellar objects with silicate absorption features, therefore indicating Class I objects (Persi et al.~\\cite{persietal}). Near infrared photometry by Persi et al., together with an 8 - 13 $\\mu$m CVF spectrum of IRS 1 (IRAS 11097-6102 in the IRAS Point Source Catalogue), show that the majority of stars ($>$ 70\\%) have an infrared excess in the JHK$_s$ plane. 19 of these sources could be matched with the present data, 15 of which we show can also be classified as having an infrared excess in the JHK$_s$L plane. The sources discussed by Persi et al. are confined to the central region, confirming the youth of the cluster. Radio recombination lines were detected by McGee \\& Gardner (\\cite{mcgeegardner}), Wilson et al. (\\cite{wilsonetal}) and De Pree et al. (\\cite{depreeetal}). The detection of maser sources in CH$_3$OH and H$_2$O (Caswell et al.~\\cite{caswelletal95} and~\\cite{caswelletal89} respectively) are indications of early stages of star formation in a dense circumstellar environment. Thorough investigations of the central region have been undertaken by Figuer\\^edo et al. (\\cite{figueredoetal}) in the near-infrared (NIR), and Barbosa et al. (\\cite{barbosaetal}) in the mid-infrared (MIR). Nine of the MIR sources match sources seen in the L-band image presented here. In the NIR, JHK$_s$ colour-colour and colour-magnitude diagrams show sources affected by excess emission, indicating the presence of circumstellar disks around the less massive members of the cluster (Figuer\\^edo et al.~\\cite{figueredoetal}). Eight spectra of the brightest sources show rising continua towards the IR. Three of these have a clear infrared excess. The detection of CO bandheads (2.2935$\\mu$m) in emission and absorption indicates the presence of several sources still heavily embedded in their stellar birthclouds (Figuer\\^edo et al.~\\cite{figueredoetal}, Barbosa et al. \\cite{barbosaetal}). Based on the radio data (Goss \\& Shaver~\\cite{gossshaver}), RCW 57 can be classified as a Giant HII (GHII) region, with ${1.6}\\times{10}^{50}$ photons $s^{-1}$ in the UV (defining sources brighter than $10^{50}$ Lyman continuum photons per second as GHII regions (Figuer\\^edo et al.~\\cite{figueredoetal})). A possible ionizing source has been found at the peak emission of the 3.4 cm map (DePree et al.~\\cite{depreeetal}), the source being a cluster of stars that have broken out of their natal cocoons but remain hidden behind dark clouds along the line of sight (Barbosa et al.~\\cite{barbosaetal}). The radio peak emission (at $\\sim$ RA 11h11m51s, Dec -61\\degr18\\arcmin45\\arcsec (J2000)) is also hidden behind clouds in the SPIREX image. This is further confirmed by Walsh et al (\\cite{walshetal}), who find the peak of the 8.64 GHz continuum emission to lie at approximately the same position, behind clouds in the N-band image (their Fig. 2). In their follow up survey, Barbosa et al. for the first time resolved IRS 1 into four sources in the 10 $\\mu$m band, approximately 1.5\\arcsec apart from each other (Nos. 48, 50, 60 and 60b in their paper. Source numbers labelled in Table~\\ref{photres} with  `No' and a number are from Barbosa et al.). One of these shows evidence for a UC HII region and they conclude that the sources in the central region of RCW 57 are in the UC HII region phase. The position of IRS 1 coincides with the brightest L-band source in our study (m$_L$=4.1, \\#88, Table~\\ref{photres}) and is also found in the L-band by Moneti (\\cite{moneti}) with magnitude m$_L$=4.05. Barbosa et al. also found a new MIR source, without a counterpart in the NIR, which is possibly a hot core. In the central region a strong CO $J=2-1$ line at 230 GHz was observed by White and Phillips (\\cite{whitephillips}). Shock-excited H$_2$ line emission may also indicate the presence of gas outflows (Figuer\\^edo et al.~\\cite{figueredoetal}). ",
        "conclusions": "As the second of two papers on IR-excess in massive star forming regions measured by the SPIREX telescope at the South Pole, L-band photometry for the giant HII region RCW 57 has been presented. The L-band photometry from SPIREX was combined with JHK$_s$ data from 2MASS to determine the fraction of sources with IR-excess in JHK$_s$L colour-colour and colour magnitude diagrams. As for 30 Doradus, it is apparent that the JHK$_s$ data alone would considerably underestimate the fraction of IR-excess sources, with only 25 sources classified as having an IR-excess. Using the JHK$_s$L data, 75 are counted towards the total fraction of stars with an IR-excess in the JHK$_s$L diagram. More than 50\\% of detected L-band sources have an IR-excess. This is, however, likely to still be a lower estimate, since foreground contamination has not fully been taken into account. Limiting the analysis to the inner $\\sim$6\\arcmin~of the source, the disk fraction increases to $\\sim$65\\%. The results were compared to earlier surveys (Haisch et al.~\\cite{haischetal}) of clusters of varying ages and masses in the Galaxy (ages between 0.3 Myr for NGC 2024 to 4.5 Myr for NGC 2362; masses down to 0.13 M$_{\\sun}$ for NGC 2024). Although the CDF for RCW 57 lies at the lower end of what is predicted by Haisch et al., it is still consistent with their data, confirming a very high initial disk fraction ($>$80\\%) and a lifetime of $\\lesssim$6 Myr. However, our results for both RCW 57 and 30 Doradus (Paper I) suggest a faster evolution of circumstellar disks around high mass stars, since the disk fractions appear to be slightly lower. This could be caused by photoevaporation of the disks due to the intense radiation environment generated by high mass stars."
    },
    "0512/astro-ph0512648_arXiv.txt": {
        "abstract": "A galaxy-sized halo may contain a large number of intermediate mass $10^{2-4}\\ms$ compact objects (IMCOs), which can be intermediate mass black holes (IMBHs) or the CDM subhalos. We propose to directly detect the IMBHs by observing multiply imaged QSO-galaxy lens systems with a high angular resolution ($\\sim 0.03$mas), which would be achieved by the next-VLBI space missions. The silhouette of the IMBHs would appear as an either monopole-like or dipole-like variation at the scale of the Einstein radius against the QSO jets. As a byproduct, we can also directly detect the $10^{4-5}\\ms$ CDM subhalos. From a measurement of the local distortion in the surface brightness of the QSO jet, we can make a distinction between a point mass (corresponding to an IMBH) and an extended structure (corresponding to a CDM subhalo). It would be a unique probe of the IMCOs whose nature has been under the veil of mistery. ",
        "introduction": "Gravitational lensing is a powerful tool for constraining the amount of the dark matter in the form of compact objects. However, there is no stringent constraint on the intermediate mass compact objects (IMCOs) with a mass scale of  $10^{2-4}\\ms$ and on the sub-lunar mass compact objects (SULCOs) with a mass scale of $<10^{-7}\\ms$ (Inoue \\& Tanaka 2003).  In particular, a method of directly detecting the IMCOs has been highly desired because the IMCOs are important baryonic dark matter candidates as well. The IMCO can be an intermediate mass black hole or an intermediate mass cold dark matter (CDM) subhalo. The nature of both types of IMCOs has not been understood well.  In fact, the abundance of these IMCOs can be much larger than expected. Recent observation of the ultra-luminous X-ray sources suggests a presence of IMBHs not only in the neighborhood of galaxy nucleus but also in the star clusters in the galactic halo far from the nucleus (Matsumoto et al 2001, Roberts et al 2004). Furthermore, the early reionization ($z\\sim 20$) (Bennett et al. 2003) suggested by the WMAP observation of the temperature-polarization correlation may indicate strong UV radiation from massive first stars at the early period ($z \\sim 20$) with a top-heavy IMF (Cen, 2003) or that from a large number of micro-QSOs (Madau, et al. 2004), which imply the existence of a large number of IMCOs. The IMBHs are certainly important building blocks for making the super massive black holes (SMBHs). Unfortunately, their evolution has been under the veil of mistery. Measurements of the abundance and the spatial distribution of the IMBHs inside the galactic halo would shed a new light on the evolution history of the SMBHs.  Similarly, we do not know anything about the abundance of the intermediate mass  $10^{2-5}\\ms$ CDM subhalos. Although they may suffer the tidal breaking owing to the gradient in the gravitational potential of the galaxy halo, any compact objects whose size is smaller than the tidal radius can survive. Because a calculation of the survival probability during the major merger of galaxies is an intractable problem(see also Taylor \\& Babul 2005), we need to observationally constrain the abundance of the CDM subhalos.  To do so, we propose to observe radio-loud QSO-galaxy strong lens systems with a submilli-arcsecond resolution, which will be achieved by the next space VLBI missions sush as the VSOP2 (Hirabayashi et al).  Then the submillilensing effects by IMCO perturbers can be directly measured (Inoue \\& Chiba 2003). ",
        "conclusions": "If there are a sufficent number of IMCOs inside the macrolens galactic halo, then we will be able to directly detect the silhouette of the IMCOs using the next space interferometers with a submilli-arcsec resolution in the radio band. By measuring the local distortion of the QSO jet, we will be able to determine the density profile of the IMCOs. The direct detection of the IMBHs or the CDM subhalos will shed a new light on the formation process of the SMBHs and the reionization process which have been under the veil of mistery. \\vspace{-0.2cm}"
    },
    "0512/astro-ph0512154_arXiv.txt": {
        "abstract": "The gamma-ray burst (GRB) 050904 at $z = 6.3$ provides the first opportunity of probing the intergalactic medium (IGM) by GRBs at the epoch of the reionization. Here we present a spectral modeling analysis of the optical afterglow spectrum taken by the Subaru Telescope, aiming to constrain the reionization history. The spectrum shows a clear damping wing at wavelengths redward of the Lyman break, and the wing shape can be fit either by a damped Ly$\\alpha$ system with a column density of $\\log (N_{\\rm HI}/{\\rm cm^{-2}}) \\sim 21.6$ at a redshift close to the detected metal absorption lines ($z_{\\rm metal} = 6.295$), or by almost neutral IGM extending to a slightly higher redshift of $z_{\\rm IGM,u} \\sim 6.36$. In the latter case, the difference from $z_{\\rm metal}$ may be explained by acceleration of metal absorbing shells by the activities of the GRB or its progenitor. However, we exclude this possibility by using the light transmission feature around the Ly$\\beta$ resonance, leading to a firm upper limit of $z_{\\rm IGM,u} \\leq 6.314$. We then show an evidence that the IGM was largely ionized already at $z=6.3$, with the best-fit neutral fraction of IGM, $x_{\\rm HI} = 0.00$, and upper limits of $x_{\\rm HI} < 0.17$ and 0.60 at 68 and 95\\% C.L., respectively.  This is the first direct and quantitative upper limit on $x_{\\rm HI}$ at $z > 6$. Various systematic uncertainties are examined, but none of them appears large enough to change this conclusion. To get further information on the reionization, it is important to increase the sample size of $z \\gtrsim 6$ GRBs, in order to find GRBs with low column densities ($\\log N_{\\rm HI} \\lesssim 20$) within their host galaxies, and for statistical studies of Ly$\\alpha$ line emission from host galaxies. ",
        "introduction": "\\label{section:introduction}  Although more than 30 years have passed since the discovery of gamma-ray bursts (GRBs) (Klebesadel, Strong, \\& Olson 1973, see, e.g., M\\'esz\\'aros 2002; Piran 2004 for recent reviews), it was rather recent that GRBs became widely recognized as a unique tool of cosmological studies and exploring the early universe.  It was just before GRBs were proven to have the cosmological origin (Metzger et al. 1997) that the first attempt to study the cosmic star formation history by using GRBs was made (Totani 1997; Wijers et al. 1998). The potential use of GRBs as a probe of the reionization history of the intergalactic medium (IGM) was also pointed out (Miralda-Escud\\'e 1998), but at that time, it was thought that GRBs can probe the universe at most up to modest redshifts of $z \\sim $ 3 by instruments available at that time or in the near future.  Such notion was, however, soon discarded by the discoveries of extremely luminous GRBs like GRB 971214 (Kulkarni et al. 1998) and 990123 (Kulkarni et al. 1999), which could be detected even at redshifts beyond 10. Then it did not take long time before astrophysicists started to discuss GRBs as a promising lighthouse to study the extremely high redshift universe, potentially giving cosmologically important information including the population III star formation and the reionization history (Lamb \\& Reichart 2000; Ciardi \\& Loeb 2000).  During 2000--2004, satellites such as the BeppoSAX and the HETE-2 continued to discover more and more GRBs, and there was important progress including the establishment of the firm connection between long duration GRBs and energetic supernovae (Stanek et al. 2003; Hjorth et al. 2003a).  However, the distance of GRBs did not extend to very high redshift, with the highest record of $z = 4.5$ (Andersen et al. 2000). The launch of the Swift satellite in 2004 then allowed the GRB community to search for fainter and more distant GRBs with improved detection rate. The GRB 050904 was discovered by the Swift on 2005 September 4 at 01:51:44 UT (Cusumano et al. 2005), and follow-up photometric observations of the afterglow found a strong spectral break between optical and near-infrared (NIR) bands, indicating a very high redshift of $z \\sim 6$ (Haislip et al. 2005; Price et al. 2005; Tagliaferri et al. 2005). This suggestion was confirmed by the subsequent spectroscopic observation by the Subaru Telescope, which found metal absorption lines at $z = 6.295$ and the corresponding Lyman break and red damping wing (Kawai et al. 2005, hereafter Paper I).  This opens a new era of GRB observations at redshifts that are close to the cosmic reionization and comparable to those of the most distant galaxies (Taniguchi et al. 2005) and quasars (White et al. 2003; Fan et al. 2006).  Here we report a detailed analysis and interpretation of the optical afterglow spectrum of GRB 050904 presented in Paper I, to derive implications for the reionization.  The famous Gunn-Peterson (GP) test tells us that the IGM is highly ionized at $z \\lesssim 5$ (Gunn \\& Peterson 1965), while the observations of the cosmic microwave background radiation (CMB) indicates that the universe became neutral at the recombination epoch of $z \\sim 1100$ (Spergel et al. 2003). The reionization of the IGM is believed to have occurred during $z \\sim$ 6--20 by the first stars and/or quasars, and the precise epoch and nature of the reionization is one of the central topics in the modern cosmology (see Loeb \\& Barkana 2001; Barkana \\& Loeb 2001; Miralda-Escud\\'e 2003; Haiman 2004 for recent reviews).  The dramatic increase of the optical depth of the Ly$\\alpha$ forest with increasing redshift at $z \\gtrsim 5.2$ (Djorgovski et al. 2001) and the subsequent discovery of broad and black troughs of Ly$\\alpha$ absorption (the GP troughs) in the spectra of $z \\gtrsim 6$ quasars (Becker et al. 2001; White et al. 2003; Fan et al. 2003, 2006) indicate that we are beginning to probe the epoch of reionization (Fan et al. 2002; Cen \\& McDonald 2002; but see also Songaila \\& Cowie 2002).  On the other hand, the polarization observation of the CMB by the Wilkinson Microwave Anisotropy Probe (WMAP) indicates a much higher redshift of reionization, $z = 17 \\pm 5$ (Kogut et al. 2003). Some theorists have argued that the hydrogen in the IGM could have been reionized twice (Wyithe \\& Loeb 2003; Cen 2003).  Because the cross section of the Ly$\\alpha$ resonance absorption is so large, the light blueward of the Ly$\\alpha$ wavelength at the source is completely attenuated if the IGM neutral fraction $x_{\\rm HI} \\equiv n_{\\rm HI}/ n_{\\rm H}$ is larger than $\\sim 10^{-3}$, and hence the Ly$\\alpha$ trough of $z\\sim 6$ quasars gives a constraint of only $x_{\\rm HI} \\gtrsim 10^{-3}$.  The cross section becomes much smaller for longer wavelength photons than the Ly$\\alpha$ resonance, and the spectral shape of the red damping wing of the GP trough can potentially be used to measure $x_{\\rm HI}$ more precisely (Miralda-Escud\\'e 1998). However, applying this method to quasars is problematic because of the uncertainties in the original unabsorbed quasar spectra and the proximity effect (Bajtlik, Duncan, \\& Ostriker 1988), i.e., the ionization of surrounding IGM by strong ionizing flux from quasars (Cen \\& Haiman 2000; Madau \\& Rees 2000). Therefore, though some authors suggested $x_{\\rm HI} \\gtrsim 0.1$ using $z \\sim 6$ quasar spectra (Mesinger \\& Haiman 2004; Wyithe \\& Loeb 2004; Wyithe, Loeb, \\& Carilli 2005), these estimates are generally model dependent.  The Ly$\\alpha$ line emission is seriously attenuated if it is embedded in the neutral IGM, and hence the Ly$\\alpha$ line emissivity of galaxies at $z \\gtrsim$ 6 is another probe of the reionization. Therefore the detection of Ly$\\alpha$ emission from many galaxies at $z \\gtrsim 6$ (Hu et al. 2002; Kodaira et al. 2003; Taniguchi et al 2005) may indicate that the universe had already been largely ionized at that time (Malhotra \\& Rhoads 2004; Stern et al. 2005; Haiman \\& Cen 2005). However it should not be naively interpreted as implying small $x_{\\rm HI}$, since these Ly$\\alpha$ emitters (LAEs) are selected by strong Ly$\\alpha$ emission and hence they may be biased to those in ionized bubbles created by themselves or clusters of undetected sources (Haiman 2002; Wyithe \\& Loeb 2005).  On the other hand, the Lyman-break galaxies (LBGs) selected by broad-band colors are free from the selection bias about Ly$\\alpha$ emission, but the statistical nature of Ly$\\alpha$ emission from LBGs at $z \\gtrsim 5$ is not yet well understood (Ando et al. 2004; Bouwens et al. 2004; Dickinson et al. 2004; Giavalisco et al. 2004; Stanway et al. 2004), compared with those at $z \\sim 3$ (Shapley et al. 2003). Furthermore, the Ly$\\alpha$ line emission from lower-redshift starburst galaxies is often redshifted with respect to the systemic velocity of galaxies (e.g., Pettini et al. 2000), and such a relative redshift will lead to a higher detectability of LAEs at $z \\gtrsim 6$, indicating a possible systematic uncertainty in the reionization study by Ly$\\alpha$ emission.  GRBs have a few advantages as a probe of the cosmic reionization, compared with quasars or LAEs/LBGs. GRB afterglows are much brighter than LAEs/LBGs and comparable to or even brighter than quasars if they are observed quickly enough after the explosion. The Ly$\\alpha$ or ultraviolet luminosity of host galaxies is irrelevant to the detectability of GRBs, and hence GRBs can probe less biased regions in the early universe, while quasars and bright LAEs/LBGs are likely biased to regions of rapid structure formation with strong clustering.  In most cases it is expected that the IGM ionization state around GRB host galaxies had not yet been altered by strong ionizing flux from quasars. Finally, the spectrum of GRB afterglows has a much simpler power-law shape than complicated lines and continuum of quasars and LAEs/LBGs, and hence model uncertainty can greatly be reduced. Especially, a detailed fitting analysis of the damping wing of the GP trough in an afterglow spectrum may lead to a precise determination of $x_{\\rm HI}$.  In this paper we take these advantages for the first time and derive the first implications for the reionization from GRBs by using the spectrum of GRB 050904. In \\S \\ref{section:spectrum}, we briefly describe the overall features of the observed spectrum. Formulations of the model fitting will be given in \\S \\ref{section:model}.  Before the fitting, theoretically possible ranges of some model parameters will be defined (\\S \\ref{section:z-relations}). Following the description of the fitting procedure (\\S \\ref{section:procedure}), the fitting results to the red damping wing will be given in \\S \\ref{section:separate-fit} and \\S \\ref{section:joint-fit}. The additional constraint from the Ly$\\beta$ feature is discussed in \\S \\ref{section:ly-beta}. We then derive a quantitative constraint on the IGM neutral fraction, taking into account all constraints and identified uncertainties in \\S \\ref{section:x_HI}. In \\S \\ref{section:discussion} we discuss the prospects of the future GRB data as the reionization probe, based on the lessons from this first data set; \\S \\ref{section:discussion-DLA} is for the neutral hydrogen column density within host galaxies, and \\S \\ref{section:discussion-lya-emission} is for the Ly$\\alpha$ line emission from host galaxies. A summary and conclusions will be presented in \\S \\ref{section:summary}.  Throughout this paper, we use the WMAP values of the cosmological parameters in the flat universe: $H_0 = 71 \\ \\rm km \\ s^{-1} Mpc^{-1}$, $\\Omega_B = 0.044$, and $\\Omega_M = 0.27$ (Spergel et al. 2003), and the primordial helium mass fraction in the total cosmic baryon, $Y_p = 0.25$ (Kawasaki, Kohri, \\& Moroi 2005). ",
        "conclusions": "\\label{section:summary}  We presented a comprehensive theoretical modeling of the red damping wing of the Ly$\\alpha$ absorption found in the optical afterglow spectrum of GRB 050904 at $z \\sim 6.3$, which provides the first opportunity of studying the cosmic reionization by using GRBs.  We tried to model the observed damping wing shape by the two components of absorbers: one is by the DLA associated to the GRB host galaxy, and the other is by neutral hydrogen in the IGM. The redshift of the metal absorption lines in the spectrum is $z_{\\rm metal} = 6.295 \\pm 0.002$, but we allowed different values of $z_{\\rm DLA}$ (the DLA redshift) and $z_{\\rm IGM,u}$ (the upper extension bound of neutral hydrogens in the IGM), and discussed various theoretical possibilities for the deviation of these parameters from $z_{\\rm metal}$.  The shape of the red damping wing can be explained either by the DLA at $z_{\\rm DLA} \\sim z_{\\rm metal}$ with $\\log N_{\\rm HI} \\sim $ 21.6, or by almost neutral IGM extending to a higher redshift of $z_{\\rm IGM,u} \\sim $ 6.36. Though the DLA seems a more straightforward interpretation, we cannot exclude the latter possibility simply by the redshift difference, since blueshift of metal absorption lines up to a few thousands km/s with respect to the restframe of the host galaxy has been observed in a few GRBs.  However, we found that the Ly$\\beta$ feature can be used to break this degeneracy, since the two different solutions predict different wavelengths at which the Ly$\\beta$ GP trough ends.  Then we concluded that the damping wing is mostly contributed from the DLA at $z_{\\rm DLA} \\sim z_{\\rm metal}$, and derived a firm upper bound of $z_{\\rm IGM,u} \\leq 6.314$.  We argued that the DLA is likely to be associated physically with the metal absorption lines, and the inferred column densities, metallicities, and depletion of silicon are all reasonable as a DLA found in a GRB afterglow.  Next we examined the preferred value of the IGM neutral fraction, $x_{\\rm HI}$, in the viable model of $z_{\\rm IGM,u} = z_{\\rm DLA} = 6.295$. Treating $N_{\\rm HI}$ of the DLA as a free parameter, we found that a smaller value of $x_{\\rm HI}$ is favored with the best-fit value of $x_{\\rm HI} = 0.00$, and upper limits of $x_{\\rm HI} < 0.17$ and 0.60 (68 and 95\\% C.L., respectively) were derived.  We examined various possible systematic uncertainties that could affect this result, including the afterglow spectral index, dust extinction at the host galaxy, the redshift parameters of the DLA and IGM absorptions, weak unidentified absorption lines, and time variability of the DLA column density. We found that none of these effects is large enough to change the above result. Hence we conclude that the universe was largely ionized already at $z \\sim 6.3$, excluding the completely neutral IGM at $\\sim$ 99\\% C.L.  This is the first quantitative {\\it upper} limit on $x_{\\rm HI}$ at $z \\gtrsim 6$ by a direct method,\\footnote{Though there are some ``data points'' of the IGM optical depth at $z \\gtrsim 6$ derived directly from the GP trough of quasar spectra (Fan et al. 2006), these are by averaging sharp spikes of transmission in a wavelength range, which are probably corresponding to small regions of ionized bubbles. Therefore there is no upper limit on the mass-weighted or volume-weighted optical depth.} being consistent with the recent results by an indirect approach using the number density evolution of LAEs (Malhotra \\& Rhoads 2004; Stern et al. 2005; Haiman \\& Cen 2005). Since the IGM is optically thin for photons in the damping wing region, all the IGM neutral hydrogens at $z \\sim 6.1$--6.3 (see Fig. \\ref{fig:tau_IGM_z_IGMl}) contribute to the damping wing, allowing us to derive a robust constraint on the mass-weighted $x_{\\rm HI}$ which is insensitive to any clumpiness of IGM within this redshift interval. Combined with the suggestions of $x_{\\rm HI} \\gtrsim 0.1$ from quasar spectra (Mesinger \\& Haiman 2004; Wyithe, Loeb, \\& Carilli 2005), a plausible value of $x_{\\rm HI} \\sim 0.1$ is suggested for the IGM at $z \\sim 6$--6.3.  The large DLA column density of $\\log N_{\\rm HI} \\gtrsim 21$ dominates the IGM absorption extending to the same redshift, making it difficult to derive a stronger constraint on $x_{\\rm HI}$ than derived here. However, some GRBs have low column densities of $\\log N_{\\rm HI} \\lesssim 20$, and detection of such GRBs at $z \\gtrsim 6$ will be a promising chance to get better information for the reionization history of the universe.  We did not detect Ly$\\alpha$ emission from the host galaxy, leading to an upper limit for extinction-uncorrected star formation rate as SFR $\\lesssim 0.79 M_\\odot$/yr.  We discussed the potential of Ly$\\alpha$ emission from GRB host galaxies as a reionization probe.  Statistically smaller equivalent width and transmission only at the red tail of Ly$\\alpha$ emission are expected for GRB host galaxies before the reionization, compared with those at lower redshifts. This may be tested by using future large samples of GRBs at $z \\gtrsim 6$.  As a reionization probe using Ly$\\alpha$ emission, GRB host galaxies have an advantage of being free from the selection bias compared with LAEs. Another advantage compared with both LAEs and LBGs is that an accurate redshift determination is possible by absorption lines in bright afterglows. On the other hand, a possible disadvantage is that typical GRB host galaxies and their absolute Ly$\\alpha$ luminosity may not be as bright as LAEs and LBGs found in deep surveys.  We conclude that the GRB 050904 has opened a new era of cosmological study by GRBs, and future data will give us even more unique and important information about the epoch when the early-generation luminous objects changed the physical state of almost all the baryonic matter in the cosmos.  We would like to thank the Subaru Telescope staff for their warm assistance in taking this invaluable data. We would also like to thank the referee for useful comments. This work was supported in part by the Grant-in-Aid for the 21st Century COE ``Center for Diversity and Universality in Physics'' from the Ministry of Education, Culture, Sports, Science, and Technology (MEXT) of Japan.  T.T. and N.K. were also supported by the Grant-in-Aid for Scientific Research from the MEXT, 16740109 and 14GS0211, respectively."
    },
    "0512/hep-th0512033_arXiv.txt": {
        "abstract": "Taking seriously the interpretation of black hole entropy as the logarithm of the number of microstates, we argue that thermal gravitons may undergo a phase transition to a kind of black hole condensate. The phase transition proceeds via nucleation of black holes at a rate governed by a saddlepoint configuration whose free energy is of order the inverse temperature in Planck units. Whether the universe remains in a low entropy state as opposed to the high entropy black hole condensate depends sensitively on its thermal history. Our results may clarify an old observation of Penrose regarding the very low entropy state of the universe. ",
        "introduction": "\\label{intro}  Years ago Penrose noticed that the universe must have begun in a very low entropy state \\cite{penrose}. By considering the entropy of black holes, he argued that the current state of the universe has significantly lower entropy than the maximum possible entropy state. For example, while holding the number of baryons fixed one could increase the total entropy tremendously by letting matter collapse into black holes \\cite{grentropy}. Indeed, it seems that while the matter degrees of freedom were born hot, i.e., in a maximum entropy thermal state, the gravitational degrees of freedom were born in a very special low entropy state. Interpreting entropy as the logarithm of phase space volume, a low entropy state is an exponentially unlikely state and hence can only result from fine-tuned initial conditions \\cite{ergodic}. Reasoning along these lines suggests that spacetimes with numerous horizons, perhaps resembling a dense agglomeration of black holes, occupy an exponentially larger fraction of gravitational phase space than smooth spacetimes like the usual Friedmann-Robertson-Walker (FRW) cosmologies.  For related discussions, see, e.g., \\cite{carroll,wald} and references contained therein.  One may ask whether special initial conditions at the Planck scale are sufficient to produce the low entropy universe we see today. It might be the case that interactions with thermal matter in the early universe inevitably cause the gravitational degrees of freedom to thermalize as well. Such a thermal state, assuming ergodicity of gravity, would likely evolve to a configuration of much higher entropy, and hence a cosmology very different from the one we observe.  Since black holes are our only hint at the highly entropic configurations of gravity \\cite{horizon}, they should play a prominent role in the transition from low entropy to high entropy spacetimes. In this paper we suggest a specific mechanism involving the nucleation of black holes from a thermal graviton state. We note that the corresponding nucleation rate from a thermal matter state is much smaller, and probably irrelevant cosmologically. The mechanism we describe provides a plausible means by which Penrose's ergodic evolution could proceed. We examine whether the transition to a new, highly entropic, phase of condensed black holes can occur in standard big bang cosmology. The result depends sensitively on the thermal history of the universe at early times. Moreover, the relevant energy scales are all higher than the energy scale at which an inflationary epoch is usually assumed to take place. Therefore, we are considering a phase transition which only may take place before and not after inflation. Presumably, the probability for a given patch to inflate would be affected by whether or not that patch has undergone a phase transition to the high entropy phase.   We should note that there is not a consensus on the issue of whether gravity is ergodic nor on the interpretation of the gap between the maximum allowed and the actual entropy in an FRW spacetime. Tipler \\cite{tipler} showed that under a reasonable set of assumptions, closed universes are technically not ergodic, i.e., there is no Poincar\\'e recurrence. Moreover, Barrow \\cite{barrow} has pointed out that in a spacetime restricted to be FRW there is necessarily an entropy gap, i.e., the entropy in thermal radiation is much less than the entropy associated with a black hole of horizon size. However, the phenomena we described in the previous paragraph, which are investigated in this paper, are independent of these larger questions about general relativity. That is, the mechanism by which black holes are nucleated occurs on sub-horizon time and length scales. The statistical approach we take below is justified by the presence of a thermal bath of gravitons or other particles, whose existence is not in dispute. In this sense we do not require any assumption of ergodicity, {\\it except in some small sub-horizon patch.}  In Sec. \\ref{stat} we consider black hole nucleation in a system of thermal gravitons and compare to a thermal system of matter. In Sec. \\ref{therm} we show that gravitons may thermalize in the early universe even if they started out cold. We determine the conditions necessary for a phase transition to a black hole condensate via percolation in Sec. \\ref{perc}. Finally, in Sec. \\ref{disc} we relate these results to Penrose's observation.  We use Planck units throughout, i.e., $\\hbar = c = G = k_{\\rm B} = 1$. ",
        "conclusions": "\\label{disc}  We examined a possible first order phase transition of spacetime to a black hole phase with high entropy.  Percolation of the high entropy phase occurs if gravitons are ever in a thermal state with temperature above $T_0^c$, either because they were born hot at the Planck epoch or because they were thermalized due to interactions with thermal matter. It seems possible, as suggested by entropic arguments, that almost all of gravitational phase space is accounted for by the nonperturbative phase. However, we find that the low entropy phase is metastable over timescales which are exponentially sensitive to the temperature, and potentially quite long.  One may wonder how inflation changes this conclusion. We note that $T_0^c$ is higher than the energy scale at which inflation is usually assumed to take place. If gravitons are never thermalized above a temperature of $T_0^c$ then presumably inflation would simply take place as originally envisioned. However, if gravitons are ever thermal with a temperature above $T_0^c$ then we would speculate that it may be less probable for a given patch to inflate, although depending on the details of the model some non-zero probability may remain, even if a phase transition to the high entropy black hole phase does occur.  Hot gravitons with temperature slightly below $T_0^c$ will not lead to a phase transition; they will simply be red-shifted away. Both gravitons and matter may be born hot, as long as the temperature of the universe (either initially or after a period of inflation) is never greater than $T_0^c$.  This does not require fine tuning because $T_0^c$ is of the same order as $q_{\\rm pert}$, the energy scale below which quantum gravity effects are small. It may still be the case that the initial conditions represent a subset of measure zero in the total phase space, which is dominated by the nonperturbative black hole phase \\cite{anthropic}. However, our analysis does show that once the initial choice of the low entropy phase is made, no transition to the high entropy phase need occur.  These conclusions remain unchanged in a spacetime of arbitrary dimension $d$. For hot gravitons in $d$ dimensions, the exponent governing the nucleation rate of Eq. (\\ref{lambda}) goes as $F_*/T \\sim T^{-(d-2)}$, while for matter $F_*/T \\sim T^{-(d-1)(d-2)/2}$. For $d > 3$, black hole nucleation is suppressed more strongly in the matter system than in the gravitational one.  Moreover, as $d \\rightarrow \\infty$, $T_0^c$ increases, meaning a transition to the black hole condensate phase is less likely in a universe with a large number of spacetime dimensions.  {\\bf Note added}: After this paper was completed we became aware of earlier work using Euclidean path integral methods in which Eqs. (\\ref{free_energy}), (\\ref{f_star}), and (\\ref{lambda}) were independently derived \\cite{gross}. In these calculations imaginary time boundary conditions are applied to the gravitational field. Therefore, those authors were also studying thermal gravity and not only thermal matter. For discussion of black hole phase transitions, see \\cite{hu}.          \\bigskip   \\begin{center} \\textbf"
    },
    "0512/astro-ph0512632_arXiv.txt": {
        "abstract": "Many of the known extrasolar planets are ``hot Jupiters,'' giant planets with orbital periods of just a few days. We use the observed distribution of hot Jupiters to constrain the location of its inner edge in the mass--period diagram.  If we assume a slope corresponding to the classical Roche limit, then we find that the edge corresponds to a separation close to {\\it twice\\/} the Roche limit, as expected if the planets started on highly eccentric orbits that were later circularized. In contrast, any migration scenario would predict an inner edge right at the Roche limit, which applies to planets approaching on nearly circular orbits. However, the current sample of hot Jupiters is not sufficient to provide a precise constraint simultaneously on both the location and slope of the inner edge. ",
        "introduction": "Early discoveries of hot Jupiters hinted at a pile-up near a 3-day period, but recent transit surveys and more sensitive radial velocity observations have discovered planets with even shorter periods. The data now suggest that the inner limit for hot Jupiters is not defined by an orbital period, but rather by a tidal limit, which depends on both the separation and the planet-star mass ratio (Fig.\\ 1).  This would arise naturally if the inner edge were related to the Roche limit, the critical distance within which a planet would start losing mass (Faber et al.\\ 2005).  The Roche limit separation, $a_R$, is given by $ R_P = 0.462\\, a_R \\,\\mu^{1/3}$, where $R_P$ is the radius of the planet, and $\\mu=m/M_*$ is the planet-star mass ratio.  The many formation scenarios proposed for hot Jupiters can be divided into two broad categories.  The first involves slow migration on quasi-circular orbits, perhaps due to interaction with a gaseous disk or planetesimal scattering (Murray et al.\\ 1998; Trilling \\etal~1998). This would result in an inner edge precisely at the Roche limit. The second category invokes tidal circularization of highly eccentric orbits with very small pericenter distances, following planet-planet scattering (Rasio \\& Ford 1996; Weidenschilling \\& Marzari 1996; Ford et al.\\ 2001; Papaloizou \\& Terquem 2001; Marzari \\& Weidenschilling 2002), secular perturbations from a wide binary companion (Holman et al.\\ 1997; Wu \\& Murray 2003), or tidal capture of free-floating planets (Gaudi 2003).  These would result in a limiting separation of {\\it twice\\/} the Roche limit, assuming that circularization can take place without significant mass loss from the planet\\footnote{This is very easy to show: consider a  planet on an initially eccentric orbit,  with initial  eccentricity  $e$  and pericenter  distance $r_p$.  Circularizing  this   orbit  under  ideal conditions leads to dissipation of energy but conservation of mass and angular momentum. Simply equating the angular momentum of the initial and final orbits gives a final circularized radius $a  = r_p (1+e) \\simeq 2 r_p$  for $e \\simeq 1$.}  (Faber et al.\\ 2005; Gu et al.\\ 2003; Rasio et al.\\ 1996). ",
        "conclusions": "The current distribution of hot Jupiters shows a cutoff that is a function of orbital period and planet mass. Under the assumption that the slope of this cutoff follows the Roche limit, our Bayesian analysis solidly rejects the hypothesis that the cutoff occurs inside or at the present Roche limit. This is in constrast to what would be expected if these planets had slowly migrated inwards on  quasi-circular orbits and with radii close to the presently measured values around $1.2\\,R_{J}$.  If confirmed by future analyses of a more extensive data set, this result would be highly significant, as it would eliminate a broad class of popular migration scenarios for the formation of hot Jupiters.  Instead, our analysis shows that this cutoff occurs at a distance nearly twice that of the Roche limit, as expected if the planets had been circularized from a highly eccentric orbit. These findings suggest that hot Jupiters may have formed via planet-planet scattering (Rasio \\& Ford 1996), tidal capture of free floating planets (Gaudi 2003), or secular perturbations from a highly inclined binary companion (Holman et al.\\ 1997).  Regardless of the exact mechanism, our model would require that the hot Jupiters all started on highly eccentric orbits and survived the strong tidal dissipation needed to circularize their orbits. A few caveats are worth mentioning. Our study addresses the statistical properties of the population of hot-Jupiters  and does not attempt to advance the state of  knowledge of any specific planet. In particular, we adopt average properties of an assumed distribution that is analogous to---and derived from---the presently known distribution of hot Jupiters, but we do not consider or solve for the specific properties of any individual extrasolar planet. Moreover, strongly non-random or non-gaussian effects would be poorly modeled with the technique developed here.  An alternative explanation is that the planets migrated inwards at an early time and arrived at their Roche limit on a quasi-circular orbit when their radii were still $\\ge 2\\,R_{J}$ (Burrows et al.\\ 2000). The dissipation process causing the migration must then have stopped immediately afterwards to avoid further decay of the orbit as the planets continued to cool and contract. We find this scenario unattractive, especially since there is no natural explanation for the factor of 2 in this case.  Yet another alternative is that short-period giant planets are destroyed by some process {\\em before\\/} they reach the Roche limit.  HST observations of HD 209458 indicate absorption by matter presently beyond the Roche lobe of the planet and have been interpreted as evidence for a wind leaving the planet and powered by stellar irradiation (Vidal-Madjar et al.\\ 2003, 2004).  Further theoretical work will help determine under what conditions these processes can cause significant mass loss (e.g., Hubbard et al.\\ 2005) and whether complete destruction could occur rather suddenly when the orbital radius decreases below $\\sim 2 a_{R}$.  Future planet discoveries will either tighten the constraints on the model parameters or provide evidence for the existence of planets definitely closer than twice the Roche limit. Additionally, future discoveries of transiting hot-Jupiters around young stars could help discriminate between the above alternatives. Moreover, new detections of lower-mass planets with very short periods could help better constrain the shape of the inner cutoff as a function of mass.  In the future, an improved statistical analysis could also include such low- mass planets, where surveys are not yet complete."
    },
    "0512/astro-ph0512404_arXiv.txt": {
        "abstract": "Archival observations of 18 starburst galaxies that span a wide range in metallicity reveal for the first time a correlation between the ratio of emission line fluxes of [FeII] at 26 $\\mu$m and [NeII] at 12.8 $\\mu$m and the 7.7 $\\mu$m PAH strength, with the [FeII]/[NeII] flux ratio decreasing with increasing PAH strength. We also find a strong correlation between the [FeII]/[NeII] flux ratio and the host galaxy metallicity, with the flux ratio decreasing with increasing metallicity.  Since [FeII] emission  has been  linked primarily  to supernova shocks, we attribute the  high [FeII]/[NeII] ratios in low-metallicity galaxies to enhanced supernova activity. We consider this to be a dominant mechanism for  PAH  destruction,  rather than grain destruction in photoionized regions surrounding young massive stars. We also consider whether  the   extreme  youth  of  the low-metallicity galaxies  is responsible for the lack of PAH emission. ",
        "introduction": "Recent  surveys  in  the  UV and  IR \\citep{gal05,eng05,hogg05,mad05,ros05} have highlighted  the importance of  low-metallicity  galaxies for  the investigation of vigourous episodes of star formation.  In particular, nearby  blue compact dwarf  (BCD)  galaxies  provide evidence  that isolated, small, low-metallicity  galaxies may experience  high star formation levels in the present epoch \\citep{hop02}.  Since the advent of {\\it ISO} it  has been  known that very  low metallicity  ($\\it Z$ $\\ll$ 1) galaxies  show very little in the  way of polycyclic aromatic hydrocarbonate   (PAH) \\citep{pug89}    emission   in   the   mid-IR \\citep{mad00,stu00}.  PAH emission  is thought to be linked  to star forming activity \\citep{rou01}  - UV photons  from massive young stars excite the PAH carriers, and therefore trace the extent of star formation.   Low PAH  luminosity in strongly  star forming galaxies is  contrary to what has  been seen previously by  {\\it ISO} surveys  - in starbursts, ubiquitous  strong PAH  emission was  detected (see review by Genzel \\& Cesarsky 2000). However from  {\\it  ISO}  data it has been tentatively noted that as  the metallicity of the galaxy increases, so too does the strength of the PAH  emission \\citep{stu00,haas02}. This begs the question - what causes the PAH emission deficiency in low-metallicity galaxies?  Is it a case of PAH carriers being destroyed or is the PAH deficiency intrinsic?  Hirashita et al. (2002) note that dust destruction is the dominant ISM process when the metallicity of  a low metallicity galaxy such as blue compact   dwarfs    (BCDs)   reaches    12   +   log (O/H) $\\leq$ $\\sim$8. Furthermore,  recent work  by \\citet{gal05} suggest that the small size of ISM grains emitting  in the mid and far-IR may be  due to destruction by shock waves from supernovae (SN) in these galaxies.  In  low-metallicity galaxies, star formation  is dominated by  massive stars (M$_{init}$ $\\geq$ 35M$_{\\odot}$) and high  SN rates are expected in galaxies experiencing  very recent massive  star formation \\citep{mas99}. An extremely  high SN  rate and/or extremely  energetic  SN  (with up to  10$^{52}$  ergs  being released  into the ISM \\citep{gal98})  from very massive stars could very well be responsible for the  lack of PAH in low-metallicity galaxies. In this scenario, strong SN shocks propagating into the ISM may be the culprit responsible for the PAH deficiency.   In  order to investigate  this question,  a tracer  of SN  activity is required.    The    mid-IR     [FeII]    lines    offer    an ideal extinction-insensitive tool to probe  the SN activity in actively star forming galaxies.  Iron is  an abundant refractory element; however it is highly  depleted from the gas  phase of the interstellar medium in galaxies  as  a  result   of  condensation onto  dust  grains  (e.g., de Boer, Jura \\& Shull 1987).  Forbidden line emission from low ionization states of Fe is  greatly enhanced behind hydrodynamic shock fronts  where grain processing  can result in near  solar gas phase Fe  abundances.  The near-infrared [FeII] lines at 1.257 $\\mu$m and 1.644 $\\mu$m have been widely    used to trace    the SN content    in    galaxies (e.g., Greenhouse et al. 1991, 1997; Calzetti 1997). Likewise, the $\\it a^{6}D$(7/2-9/2) transition of [FeII] near 26 $\\mu$m can also be used as a relatively pure tracer of SN activity. The low excitation potential corresponding to the transition  (550 K) may in fact imply that emission from  the mid-IR lines  is longer lasting than emission from  the higher excitation potential near-IR lines, which are thought  to persist only up  to 10$^{4}$ years \\citep{oli89} after a SN  explosion. This  feature, coupled with emission from [NeII] at 12.8 $\\mu$m which traces emission from very massive young stars, could be used to determine the SN shock emission  relative to the present ionizing photon production rate. [NeII]  is  a robust indicator   of  the ionizing photon rate  in dust-enshrouded starburst  galaxies.  For galaxies dominated by typical stellar UV fields, a simple linear proportionality between the [NeII] luminosity and the  Lyman continuum luminosity is expected from photoionization  models \\citep{tho00}. Measurements  of the [FeII]/[NeII] emission line ratio in starburst galaxies  could thus provide a sensitive measure of the supernova content of galaxies \\citep{gre91,gre96,moo88} and hence an  estimate of the  degree of grain  processing  and  destruction. Furthermore, an increased SN rate or shock  wave intensity, indicated by a higher [FeII]/[NeII] would be  highly suggestive that in low metallicity galaxies destruction by SN shocks is responsible for the absence of PAHs.  Given the availability of high spatial and spectral resolution data in the mid-IR from  {\\it Spitzer} of nuclear regions  of starbursts (where massive  stars are  located), we  can search  for [FeII] emission from nearby  starbursts, and  compare  the [FeII]/[NeII] ratio as a function  the strength of the PAH emission to  see if any correlation exists. Such a diagnostic, using  data from  galaxies of known metallicity, will us to  determine whether  SN driven shocks play an important role in the PAH  emission deficit in low-metallicity galaxies. ",
        "conclusions": "Using archival spectral data from {\\it Spitzer}, we undertook a survey of 18 starburst  galaxies in order to determine  whether SN shocks are responsible for the PAH deficiency in low-metallicity galaxies. Using the emission line ratio of [FeII] at 26 $\\mu$m to [NeII] at 12.8 $\\mu$m  plotted   against  the  PAH strength,  we found   a  strong anti-correlation,  with   the [FeII]/[NeII]  ratio decreasing  with increasing  PAH emission. Similar correlations were found   when  comparing the [FeII]/[NeII] ratio  to   the  PAH/IR luminosities and metallicity. Since  [FeII] emission has  been linked primarily to SN shocks, we  attribute the high [FeII]/[NeII] ratios in low-metallicity galaxies to enhanced SN activity, and consider this to be  the  dominant  mechanism  for  PAH destruction, rather  than grain destruction   in  photoionized regions surrounding   young massive stars.   We  also   consider whether the extreme   youth of  the low-metallicity  galaxies is responsible for the  lack of  PAH emission.  While  the  age of the known  AGB populations in  our low-metallicity  sample galaxies is less than the time required for full PAH  enrichment of the ISM, we argue that SN shocks (as evidenced by from the high [FeII]/[NeII] ratios and SN rates) further depress the PAH emission. We conclude that while SN shock destruction  of the  PAH carriers may not be fully responsible for the lack of PAH emission, it remains  a  prime culprit  for  the lack  of  PAH  emission  in  low metallicity galaxies."
    },
    "0512/astro-ph0512297_arXiv.txt": {
        "abstract": "{We introduce a method to relate a possible truncation of the star cluster mass function at the high mass end to the shape of the cluster luminosity function (LF). We compare the observed LFs of five galaxies containing young star clusters with synthetic cluster population models with varying initial conditions. The LF of the SMC, the LMC and NGC~5236 are characterized by a power-law behavior $N\\,\\dr L \\propto L^{-\\alpha}\\,\\dr L$, with a mean exponent of $<\\alpha> = 2.0 \\pm 0.2$.  This can be explained by a cluster population  formed with a constant cluster formation rate, in which the maximum cluster mass per logarithmic age bin is determined by the size-of-sample effect and therefore increases with log(age/yr).  The LFs of NGC~6946 and M51 are better described by a double power-law distribution or a Schechter function. When a cluster population has a mass function that is truncated below the limit given by the size-of-sample effect, the total LF shows a bend at the magnitude of the maximum mass, with the age of the oldest cluster in the population, typically a few Gyr due to disruption.  For NGC~6946 and M51 this  suggests a maximum mass of $\\mmax = 0.5-1\\times10^6 \\msun$, although the bend is only a 1-2 $\\sigma$ detection. Faint-ward of the bend the LF has the same slope as the underlying initial cluster mass function and bright-ward of the bend it is steeper. This behavior can be well explained by our population model. We compare our results with the only other galaxy for which a bend in the LF has been observed, the ``Antennae'' galaxies (NGC~4038/4039). There the bend occurs brighter than in NGC~6946 and M51, corresponding to a maximum cluster mass of $\\mmax = 1.3-2.5\\times10^6\\,\\msun$. Hence,  if the maximum cluster mass has a physical limit, then it can vary between different galaxies. The fact that we only observe this bend in the LF in the ``Antennae'' galaxies, NGC~6946 and M51 is because there are enough clusters available to reach the limit. In other galaxies there might be a physical limit as well, but the number of clusters formed or observed is so low, that the LF is not sampled up to the luminosity of the bend. The LF can then be approximated with a single power-law distribution, with an index similar to the initial mass function index. ",
        "introduction": "\\label{sec:introduction}  The study of young extra-galactic star clusters has become a whole new field of research since the discovery of young massive clusters. The {\\it Hubble Space Telescope (HST)} has made it possible to resolve these objects and to undertake systematic studies of the nature of these objects. Young clusters with masses in the range of our Milky Way globular clusters ($10^4 - 10^6 \\msun$) have been found in merging galaxies like the ``Antennae'' \\citep{1995AJ....109..960W}, interacting galaxies like M51 (\\citealt{2000MNRAS.319..893L}; \\citealt{2005A&A...431..905B}), starburst galaxies (\\citealt{1995AJ....110.2665M}; \\citealt{2003MNRAS.343.1285D}) and even non-interacting spiral galaxies \\citep{1999A&A...345...59L}.  Recently, a relatively young and very massive cluster was discovered in the merger remnant NGC~7252 \\citep{1993AJ....106.1354W}. Its (dynamical) mass was confirmed to be as high as $\\sim 10^8 \\msun$ \\citep{2004A&A...416..467M}. This suggests that there are clusters that fill the gap between the mass range of star clusters and that of dwarf galaxies. However, it remains to be seen if every galaxy is able to produce such massive cluster or if there are physical limitations to the maximum mass of star clusters in different environments.  The LF of star clusters is a powerful tool when studying populations of star clusters. It indirectly gives us information about the underlying mass function (MF). \\citet{1999ApJ...527L..81Z} showed for the young clusters in the ``Antennae'' galaxies that the cluster initial mass function (CIMF) can be well approximated by a power-law distribution: $N(M)\\,\\dr M\\propto\\!~M^{-\\alpha^{\\prime}}\\,\\dr M$, with an exponent\\footnote{Through the remainder of this work we will use $\\alpha^{\\prime}$ for the exponent of the cluster initial mass function and $\\alpha$ for the exponent of the total cluster luminosity function.} of $-\\alpha^{\\prime} = -2$. They derived the ages of the clusters using reddening-free parameters, which is necessary to extract the initial mass function from the total luminosity function. The exponent of the CIMF ($-\\alpha^{\\prime}$) is found to be close to $-2$ in a wide range of galaxy environments down to masses of $\\sim 10^3 \\msun$ \\citep{2003MNRAS.342..259D}. This resembles the mass function of molecular clouds \\citep{1987ApJ...319..730S}, consistent with clusters forming from molecular clouds. The LF can also be approximated with a power-law distribution: $N(L)\\,\\dr L\\propto\\!~L^{-\\alpha}\\,\\dr L$, where values for the exponent $-\\alpha$ are found in a range of $-2.4$ up till $-1.7$ (\\citealt{2002AJ....124.1393L}, hereafter L02; \\citealt{2003dhst.symp..153W}). In general the indices of the bright LFs are smaller (i.e. steeper) than the index of the CIMF (L02). One of the unanswered questions is whether the difference in slope between the mass function and LF have a physical meaning or if it is a result of different measurement techniques and artifacts like the contamination by stars at lower luminosities.  \\citet{1999AJ....118.1551W} observe a distinct bend in the LF of young star clusters in the ``Antennae'' galaxies, where the slope faint-ward is shallower than the slope bright-ward of the bend. The exact slopes depend on the different ways of correcting for stellar contamination. They argue that this could be the progenitor turn-over of the peak that is observed in the luminosity function of old globular clusters which appears at $M_V~\\simeq~-7.2$, corresponding to a mass of $2\\times10^5 \\msun$. The difficulty with directly relating the LF to a CIMF is that the LF consists of clusters of different ages, and the mass-to-light ratio of clusters changes drastically when clusters age. Between $10^7$ and $10^9$ year, a star cluster of constant mass fades about 4 magnitudes in the $V$-band.  Recently, \\citet{mengel05} observed the clusters in the ``Antennae'' galaxies in the Ks-band. They also find a double power-law behavior in the LF and argue that the mass function has a turn-over.  When there is no physical limit to the cluster mass or if there are not enough clusters to sample the CIMF up to any such limit, the mass of the most massive cluster will be determined by the total number of clusters and the slope of the CIMF (\\citealt{2003AJ....126.1836H}, hereafter H03).  A similar idea was posed by \\citet{2004MNRAS.350.1503W}, who suggest that the maximum cluster mass in a galaxy is determined by the star formation rate in that galaxy. It has not been shown yet in a large sample of galaxies that the most massive cluster is a result of size-of-sample effects. H03 only showed that this is the case in the SMC and the LMC. In this study we add M51 to the sample in order to see if the most massive object in this galaxy is also a result of size-of-sample effect or if there is a physical limit above which clusters cannot form in this galaxy. In addition, we use a method to relate a possible truncation of the mass function at the high mass end with the shape of cluster LF. To this end we introduce an analytical model to generate a cluster population and derive the LF for different cluster formation histories, disruption mechanism and CIMFs. We will show that even if the initial mass function of clusters is truncated, the brightest clusters in the sample may will still be determined by sampling statistics as found by \\citet{2003dhst.symp..153W} and L02.  The paper is organized as follows: In \\S~\\ref{sec:data} we introduce the data used for this study. In \\S~\\ref{sec:agemass} we present the masses and luminosities as a function of log(age/yr) for three galaxies in the sample. The LFs of cluster populations in five different galaxies are presented in \\S~\\ref{sec:lf}. We introduce in \\S~\\ref{sec:model} a cluster population model with which we can reproduce LFs based on various variables. The relation between maximum cluster mass and maximum cluster luminosity is discussed in \\S~\\ref{sec:maxlum}. A discussion in presented in \\S~\\ref{sec:discussion}. The conclusions are presented in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have compared observed star cluster luminosity function in five galaxies with analytical cluster population models. Our main results can be summarized as follows:  \\begin{itemize} \\item there are no clusters in M51 more massive than $M \\simeq 1\\times10^6 \\msun$, although they are predicted by the size-of-sample effect. When comparing the maximum cluster mass in increasing log(age/yr) bins, the LMC and SMC cluster population show an increase consistent with the size-of-sample of effect. The cluster population of M51, however, shows a much shallower increase. This suggests a physical upper limit to the masses of clusters M51, although the shallow increase can also be reproduced by a combined effect of cluster disruption, infant mortality and an increasing cluster formation rate. \\item when comparing the $\\redchisq$ results of different function fits to the five galaxies in our sample, we find that the LF of the SMC, the LMC and NGC~5236 can be well approximated by a power-law ($N\\,\\dr L \\propto L^{-\\alpha}\\,\\dr L$), with $1.9 < \\alpha < 2.1$, while the LF of NGC~6946 and M51 are slightly better approximated with a double power-law or Schechter function. \\item when fitting a double power-law function to the LF of NGC~6946 we find a break point at $M_V$ = $\\breakngcs$ mag. Faint-ward of the bend a power-law with exponent $\\slopengcsone$ can be fitted. Bright-ward of the bend an exponent of $\\slopengcstwo$ is found. The LF can also be well fitted by fit a Schechter function with a bend at $M_V$ = $\\mstarmf$ mag. \\item the LF of M51 shows a  break at $M_V$ = $\\breakmf$ mag. Faint-ward of the bend a power-law with exponent $\\slopemfone$ can be fitted. Bright-ward of the bend an exponent of $\\slopemftwo$ is found. The LF of M51 is also well fitted by a Schechter function with a bend at $M_V$ = $\\mstarmf$ mag. \\item the cluster LFs can be reproduced with a synthetic cluster population model. The bend in the LF of NGC~6946, M51 and ``Antennae'' galaxies can be explained with a truncation of the cluster mass function at $\\mmax = 0.5-1\\times10^6 \\msun$ (M51/NGC~6946) and $1.3-2.5\\times10^6 \\msun$ (``Antennae'').  \\end{itemize}  In a follow-up study \\citep{gieles05b} we present an improved LF of star clusters in M51 based on {\\it HST/ACS} data, taken as part of the Hubble Heritage project."
    },
    "0512/astro-ph0512542_arXiv.txt": {
        "abstract": "We present the absolute measurement of the unresolved 0.5--8 keV cosmic X-ray background (CXB) in the \\chandra\\ Deep Fields (CDFs) North and South, the longest observations with \\chandra\\ (2 Ms and 1 Ms, respectively).  We measure the unresolved CXB intensity by extracting spectra of the sky, removing all point and extended sources detected in the CDF. To model and subtract the instrumental background, we use observations obtained with ACIS in stowed position, not exposed to the sky.  The unresolved signal in the 0.5--1 keV band is dominated by diffuse Galactic and local thermal-like emission.  We find unresolved intensites in the 0.5--1 keV band of $(4.1\\pm0.3)\\times10^{-12}$ \\intens\\ for CDF-N and $(5.0\\pm0.4)\\times10^{-12}$ for CDF-S.  In the 1--8 keV band, the unresolved spectrum is adequately described by a power law with a photon index $\\Gamma=1.5^{+0.5}_{-0.4}$ and normalization $2.6\\pm0.3$ photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ sr$^{-1}$ at 1 keV.  We find unresolved CXB intensities of $(1.04 \\pm 0.14) \\times10^{-12}$ \\intens\\ for the 1--2 keV band and $(3.4 \\pm 1.7) \\times10^{-12}$ \\intens\\ for the 2--8 keV band.  Our detected unresolved intensities in these bands significantly exceed the expected flux from sources below the CDF detection limits, if one extrapolates the \\logn\\ curve to zero flux.  Thus these background intensities imply either a genuine diffuse component, or a steepening of the \\logn\\ curve at low fluxes, most significantly for energies $<$2 keV.  Adding the unresolved intensity to the total contribution from sources detected in these fields and wider-field surveys, we obtain a total intensity of the extragalactic CXB of $(4.6 \\pm 0.3) \\times10^{-12}$ \\intens\\ for 1--2 keV and $(1.7 \\pm 0.2) \\times10^{-11}$ \\intens\\ for 2--8 keV.  These totals correspond to a CXB power law normalization (for $\\Gamma=1.4$) of 10.9 photons cm$^{-2}$ s$^{-1}$ keV$^{-1}$ sr$^{-1}$ at 1 keV.  This corresponds to resolved fracations of $77\\pm3$\\% and $80\\pm8$\\% for 1--2 and 2--8 keV, respectively. ",
        "introduction": "Measurement of the cosmic X-ray background (CXB) has been a major effort in X-ray astronomy since it was first discovered in rocket flights in the 1960's \\citep{giac62}.  The total spectrum of the CXB has been studied at energies up to 50 keV by {\\it HEAO--1} and rocket experiments \\citep[e.g.,][]{mars80, mcca83, garm92, revn04}; for a review of pre-\\rosat\\ results, see \\citet{mcca90}. It was later studied in different parts of the 0.5--10 keV interval by \\rosat\\ \\citep[e.g.,][]{snow95, geor96, kunt01}, \\asca\\ \\citep[e.g.,][]{gend95,chen97,miya98,ueda99,kush02}, \\bepposax\\ \\citep[e.g.,][]{parm99,vecc99}, \\xmmnewton\\ \\citep[][hereafter DM04]{lumb02,delu04} and \\rxte\\ \\citep{revn03}. Deep observations with \\chandra\\ were used to study the component of the 0.5--8 keV CXB that resolves into point sources \\citep[e.g.,][and later works]{bran01a, giac02}. \\citet{mark03} used \\chandra\\ \\mbox{ACIS-S} to study the diffuse components at 0.5--2 keV.  For $E<1$ keV, the CXB is due to extragalactic and local discrete sources, as well as diffuse (Galactic and possibly Solar System) components \\citep[e.g.,][]{kunt00, crav00}. Above 1 keV, the CXB is primarily extragalactic in origin, and is well fit by a power law with a photon index of 1.4.  \\defcitealias{more03}{M03} \\defcitealias{delu04}{DM04}  As X-ray telescopes have improved in angular resolution, more and more of the CXB above 1 keV has been resolved into point sources, mostly active galactic nuclei \\citep[see][for a review]{bran05}.  However there still remains unresolved CXB flux of unknown origin and uncertain intensity. \\citet[][hereafter M03]{more03} added the contributions of detected point sources from a variety of narrow and wide-field X-ray surveys, and found resolved fractions of the extragalactic CXB of $94\\pm7$\\% for the 0.5--2 keV band and $89^{+8}_{-7}$\\% for the 2--10 keV band. \\citet{wors04} performed a similar study to find the fraction of the CXB that is resolved as a function of energy, using the detected sources from \\chandra\\ and \\xmmnewton\\ observations, and total CXB estimates from the \\xmmnewton\\ study of \\citetalias{delu04}.  They found that the resolved fraction of the extragalactic background decreases significantly at higher energies, from $\\sim$80\\% at $\\sim$1 keV to $\\sim$60\\% at 7 keV.  This leaves room for a possible population of faint sources that have yet to be detected, as well as truly diffuse components, including exotic ones such as emission from dark matter particle decay \\citep[e.g.,][]{abaz01}. There remains a great deal of uncertainty in the resolved fraction, largely due to uncertainty in the absolute flux of the CXB, and to cross-calibration uncertainties between different measurements.    Due to its angular resolution, \\chandra\\ is by far the best instrument for detecting point sources to very low fluxes and resolving the CXB. This study uses \\chandra\\ \\mbox{ACIS-I} for an absolute measurement of the intensity of the {\\em unresolved}\\/ X-ray background in the energy range 0.5--8 keV, after the exclusion of sources down to the lowest fluxes detectable in the deepest current exposures.  We use data from the \\chandra\\ Deep Fields \\citep[CDFs, e.g.,][]{bran01a,giac02}, which were designed specifically to resolve as much of the extragalactic CXB as possible.  As a byproduct of our measurement, we add our unresolved flux to the contributions from known sources (from the CDFs and other observations) to obtain the total intensity of the CXB.  This measurement has not previously been performed with \\chandra\\ because of difficulties in determining the instrumental ACIS backgrounds. Recent calibration using the ACIS detectors stowed out of the focal plane have made this study possible.  Throughout this paper we will define the power law photon index as $\\Gamma$, where the photon flux $F\\propto E^{-{\\Gamma}}$, and we will use 68\\% errors. ",
        "conclusions": ""
    },
    "0512/astro-ph0512556_arXiv.txt": {
        "abstract": "%  A growing number of early Be stars discovered in X-ray surveys exhibit X-ray luminosities intermediate between those of normal stars and those of most Be/X-ray binaries in quiescence. Their X-ray spectra are also much harder than those of shocked wind OB stars and can be best fitted by a thin thermal plasma with T $\\sim$ 10$^{8}$\\,K, added to a cooler and much fainter thermal component. An iron line complex including a fluorescence component is detected in many cases. There is no evidence for coherent pulsations in any of these systems but strong variability on time scales as short as 100\\,s is usually observed. Large oscillations with quasi-periods of the order of one hour or more are detected in the X-ray light curves of several sources, but have so far failed to prove to be strictly periodic. The optical and X-ray properties of these new objects strikingly resemble those of the so far unique and enigmatic Be star \\gcas\\ and define a new class of X-ray emitters. We discuss the possible origin of the X-ray emission in the light of the models proposed for \\gcas, magnetic disc-star interaction or accretion onto a compact companion object -- neutron star or white dwarf. ",
        "introduction": "X-ray surveys are the seed of many important discoveries in the field of high energy astrophysics. Population properties can be studied from the large source samples collected and interesting outliers can be discovered. The {\\it Einstein} satellite was instrumental in discovering stellar X-ray emission in most regions of the HR diagram. Early type stars can emit large amounts of X-rays (up to 10$^{33}$ \\ergs\\ in the 0.2-4.5 keV range for the most extreme cases) with a rather soft thermal spectral energy distribution with kT $\\sim$ 0.5\\,keV (Cassinelli et al. 1994). X-ray emission is believed to arise in thermalized shocks produced by instabilities in the high velocity wind, which is driven by the strong radiation pressure of the star. Recent high energy resolution X-ray spectra obtained by Chandra and XMM-Newton have confirmed and precised this picture. They indicate that the most active shocks are likely produced in the deep wind and suggest a highly clumped wind structure.  A number of relatively faint hard X-ray sources have been identified with Be stars in recent X-ray surveys. Their optical properties are those of early type stars with well developed circumstellar discs. Their X-ray properties are however distinct from those of normal OB or Be stars in showing much harder X-ray spectra and slightly enhanced X-ray luminosities.  HD 110432 is the first Be star found to exhibit these properties. Proposed as the possible counterpart of a faint X-ray source in the HEAO-1 all sky survey the identification was confirmed by Torrejon \\& Orr (2001) from BeppoSAX observations. HD 110432 was serendipitously in the field of three XMM-Newton observations (Obsids 0109480101, 0109480201 and 0109480401 for a total exposure of $\\sim$ 150\\,ksec). Two other objects, HD 161103 and SAO 49725 were also suspected of being low luminosity hard X-ray emitters from the ROSAT all sky survey (Motch et al. 1997) and confirmed by dedicated XMM-Newton observations (Lopes de Oliveira et al. 2005). Finally, the XMM-Newton galactic plane survey conducted by the Survey Science Center yielded at least two more examples of such objects, the Be star SS 397 and another Be star in the 50 Myr old open cluster NGC 6649 (Motch et al. 2005). ",
        "conclusions": "In spite of being among the first X-ray sources discovered, the mechanism explaining the X-ray emission of \\gcas\\ remains uncertain. The discovery of many more twins of this so far unique case will certainly revivify the interest in studying the puzzling origin of this very hard X-ray emission. In addition, these stars could significantly contribute to the population of galactic hard X-ray emitters since they account for nearly half of the Be/X-ray candidates found in the XMM-Newton SSC survey of the galactic plane (Motch et al. 2005). Members of this new class display very different X-ray properties from those of classical Be/X-ray binaries and ``normal'' OB stars. The narrowness of the observed range of X-ray and optical characteristics is intriguing and suggests that rather specific conditions must be met to generate the hard X-ray luminosity observed. Whatever is the true X-ray emitting mechanism, the discovery of Be + WD binaries or of a population of magnetically active early type stars is an exciting prospect."
    },
    "0512/astro-ph0512410_arXiv.txt": {
        "abstract": "We provide improved atomic calculation of wavelengths, oscillator strengths, and autoionization rates relevant to the $2\\to 3$ inner-shell transitions of Fe~VI--XVI, the so-called Fe~M-shell unresolved transition array (UTA). A second order many-body perturbation theory is employed to obtain accurate transition wavelengths, which are systematically larger than previous theoretical results by 15--45~m{\\AA}. For a few transitions of Fe~XVI and Fe~XV where laboratory measurements exist, our new wavelengths are accurate to within a few m{\\AA}. Using these new calculations, the apparent discrepancy in the velocities between the Fe~M-shell UTA and other highly ionized absorption lines in the outflow of NGC~3783 disappears. The oscillator strengths in our new calculation agree well with the previous theoretical data, while the new autoionization rates are significantly larger, especially for lower charge states. We attribute this discrepancy to the missing autoionization channels in the previous calculation. The increased autoionization rates may slightly affect the column density analysis of the Fe~M-shell UTA for sources with high column density and very low turbulent broadening. The complete set of atomic data is provided as an electronic table. ",
        "introduction": "Since the first astrophysical detection of $2\\to 3$ innershell absorption lines of Fe~VII--XII in the X-ray spectrum of IRAS 13349+2438 obtained with the Reflection Grating Spectrometer (RGS) on board \\textit{XMM-Newton} \\citep{sako01}, this so-called Fe~M-shell unresolved transition array (UTA) has been identified in many soft X-ray sources. The UTA is mainly comprised of a cluster of lines originating from $2p$--$3d$ transitions of M-shell iron ions that are located between 15--17~{\\AA}. These lines, when properly modeled, provide important information on the ionization structure, column density, and outflow kinematics of the absorbing materials. \\citet{behar01} calculated a complete set of atomic data for such modeling using the Hebrew University Lawrence Livermore Atomic Code \\citep[HULLAC,]{barshalom01}, and provided an abbreviated list of transition wavelengths, oscillator strengths, and autoionization rates. That complete data set has also been incorporated in some commonly used plasma modeling codes.  In most observations where the Fe~M-shell UTA has been identified, individual features from different ionization states are not resolved, mostly due to significant turbulent velocity broadening ($>700$~\\kms) at the source. Low statistics and insufficient spectral resolution may also hamper the identification of individual charge states. Therefore, one has to depend on global photoionized plasma models to constrain the physical conditions of the absorbing gas. In global fits, where individual UTAs are unresolved, the exact rest-frame position of each UTA is particularly important, especially since current photoionization balance calculations fail to predict the formation of Fe~M-shell ions consistently with highly ionized species of other elements \\citep{netzer03,netzer04,kraemer04}.  Recently, \\citet{holczer05} have reexamined the \\textit{Chandra} High Energy Transmission Grating Spectrometer (HETGS) spectrum of the active galactic nucleus (AGN) of NGC~3783, by coadding all 900~ks available observations. The low turbulent velocity of the outflow of NGC~3783, the high spectral resolution of the HETGS, and the very long exposure time enabled the individual absorption features from different Fe~M-shell ions to be resolved. The analysis of \\citet{holczer05} showed that the outflow velocity associated with the Fe~M-shell UTA appears to be different from those of other highly-ionized species, namely, Fe~XVII, O~VII, and O~VIII. In fact, the data prefer a zero velocity for the gas associated with Fe~M-shell ions, and a $-590$~km\\,s$^{-1}$ outflow velocity for the gas associated with more highly ionized ions. Therefore, a two component model was proposed where the Fe~M-shell ions are not part of the AGN outflow, with an important implication of much lower mass loss rate.  As stated in \\citet{holczer05}, the conclusions of that paper depend strongly on the calculated HULLAC wavelengths. The velocities of the Fe~M-shell UTA can be distinguished from those of other ions only if the HULLAC wavelengths are accurate to better than 30~m{\\AA}, which corresponds to a velocity of $\\sim$590~\\kms\\ at 16~m{\\AA}. However, it has been shown that the theoretical method employed by HULLAC can generally produce uncertainties of up to 20--50~m{\\AA} for lines in the 10--20~{\\AA} band \\citep{brown02, gu05}. Unfortunately, no laboratory measurements exist for the wavelengths of relevant ions except for a few strong lines of Fe~XVI and Fe~XV \\citep{brown01}. For these few lines, the HULLAC wavelengths are indeed underestimated by 20--30~m{\\AA}. However, the strongest and cleanest absorption features in the observed spectrum of NGC~3783 are from Fe IX--XI. If the HULLAC wavelengths for these ions are also underestimated by similar amounts, the implied outflow velocity would be similar to that derived from Fe~XVII and O ions. \\citet{holczer05} argued that errors in the calculated wavelengths would be distributed randomly in both directions, and therefore, the systematic shift observed in the data is unlikely a result of theoretical wavelength errors. Furthermore, it was shown that the HULLAC errors for the strongest 2p--3d transitions (those which determine the UTA centroid) of the Fe L-shell ions are generally much smaller than 30~m\\AA\\ \\citep{brown02, holczer05}. The present work, however, shows that these indirect indications concerning the HULLAC accuracy for the Fe~M-shell UTA wavelengths were misleading and that the assumption that these wavelengths could be used for velocity measurements was premature.  Recently, \\citet{gu05} showed that a combined configuration interaction and second order many-body perturbation theory (MBPT) is capable of predicting wavelengths of $2\\to 3$ transitions of Fe~XVII--XXIV to within a few m{\\AA}. The $2\\to 3$ lines of the Fe~M-shell UTA are similar to those of Fe XVII-XXIV, except for the addition of M-shell spectator electrons. Following the unexpected result of \\citet{holczer05}, in this paper, we apply an improved version of the MBPT method to the $2\\to 3$ transitions of Fe~VI--XVI. The resulting wavelengths of the UTA lines are found to be systematically larger than HULLAC ones by 15--45~m{\\AA}. We also give improved autoionization rates, which are significantly larger than the previous HULLAC results. The new data are used to model the NGC~3783 spectrum, and the outflow velocity derived from Fe~M-shell UTA is found to be consistent with that associated with Fe~XVII and O~VII--VIII ions.  In \\S\\ref{theory}, we describe the detail of MBPT and its implementation. \\S\\ref{result} presents the results of the calculation, and their comparison with the previous HULLAC data set. A brief summary is given in \\S\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We have developed a second order many-body perturbation theory with multi-reference model space. We apply the method to the calculation of wavelengths, oscillator strengths, and autoionization rates of Fe~M-shell UTA arising from $2p$--$3d$ transitions originating from the ground state of Fe VI-XVI. The wavelengths obtained in the present work are systematically larger than the HULLAC calculation of \\citet{behar01}; the present oscillator strengths of the strong absorption lines agree well with the HULLAC results; and the autoionization rates of \\citet{behar01} are found to be missing important autoionization channels, especially for lower charge states. Using the present data for Fe~VI--XVI, we find no evidence that the outflow velocity of the gas associated with Fe~M-shell ions is different from that derived from Fe~XVII and O~VII--VIII ions, as claimed in the recent analysis of \\citet{holczer05} using the HULLAC data. A complete list of Fe~M-shell UTA lines are given, which include the present MBPT calculation for Fe~VI--XVI, and a simple configuration interaction calculation for Fe~I--V. We recommend that this new database should be preferred over the earlier HULLAC calculation of \\citet{behar01} in the analyses of future absorption spectroscopy where Fe~M-shell UTA is prominent.  M.F. Gu and S.M. Kahn acknowledges the partial support of NASA grants NAG5-5419 and NNGG04GL76G. The research at the Technion was supported by The Israel Science Foundation (grant no. 28/03). E. Behar thanks the Stanford group for their hospitality during a visit in July 2005. We thank Shai Kaspi for making the 900~ks HETGS integrated spectrum available to us."
    },
    "0512/astro-ph0512626_arXiv.txt": {
        "abstract": "The early reionization of the intergalactic medium, which is favored from the WMAP temperature-polarization cross-correlations, contests the validity of the standard scenario of structure formation in the cold dark matter cosmogony.  It is difficult to achieve early enough star formation without rather extreme assumptions such as very high escape fraction of ionizing photons from proto-galaxies or a top-heavy initial mass function. Here we propose an alternative scenario that is additional fluctuations on small scales induced by primordial magnetic fields trigger the early structure formation.  We found that ionizing photons from Population III stars formed in dark haloes can easily reionize the universe by $z \\simeq 15$ if the strength of primordial magnetic fields is between $0.7$--$1.5 \\times 10^{-9}$ Gauss. ",
        "introduction": "The reionization process of the intergalactic medium (IGM) is one of the major remaining problems in modern cosmology.  From the Gunn-Peterson test of QSO absorption lines, it is known that the vast majority of IGM is ionized by $z \\sim 6$ \\citep{becker, fan}. The recent measurement of the cosmic microwave background (CMB) temperature and polarization cross correlations by WMAP implies that the optical depth of the universe is about $0.17$ \\citep{spergel, kogut}.  This result favors the early reionization scenario: the reionization process occurs at $z=15$ -- $20$ and the reionization sources are first stars, unlike quasars and galaxies which are known as the reionization sources previously (for the details see \\citealt{early}).  The early reionization process by the stellar sources has been studied in detail after WMAP~\\citep{cen, fuku-kawa, c-f-w, s-l, h-h}.  In these works, cold dark matter cosmogony with WMAP parameters is employed.  What they found was it is difficult to get $\\tau=0.17$ if the standard Salpeter initial mass function (IMF) is adapted.  To have early enough reionization, one needs to assume almost $100\\%$ escape fraction of ionizing photons from proto-galaxies, or introduce a top-heavy IMF.  Heavy stars may form in the early universe induced by the $\\rm H_2$ molecular cooling while it is still little known about the IMF in the early universe.  An alternative scenario to realize early reionization is to enhance the amplitude of the CDM power spectrum on very small scales. Such enhancement makes the dark haloes form earlier. Accordingly the star formation process starts early enough to achieve $\\tau=0.17$. The observational data from redshift surveys of galaxies such as 2dF and SDSS, and Ly-$\\alpha$ clouds \\citep{spergel, seljak} strongly constrain the amplitude of the density fluctuations in the scales larger than 1 Mpc.  However, the amplitude in the scales which are relevant to the first star formation, $0.01$ -- $0.1$ Mpc, is still unclear.  Therefore there is still room for considering additional power in the power spectrum on very small scales.  For example, models with initially running spectral index $n > 1$ on small scales, additional fluctuations  from the isocurvature modes, or non-Gaussian statistics are considered in the context of early reionization~\\citep{chen-cooray, avelino, s-z}.  If there exist strong enough primordial magnetic fields at the reionization epoch, these magnetic fields produce additional density fluctuations of baryons by the Lorentz force. % The magnetic tension is more effective on small scales where the entanglements of magnetic fields are larger.  Therefore we expect to have additional power in the density power spectrum on small scales as is the case of isocurvature perturbations.  Structure formation by magnetic fields was first discussed by \\citet{wasserman}. More detailed analysis was carried by \\citet{k-o} and the influence on the formation of large scale structure was recently estimated by \\citet{gopal}. \\citet{sethi} pointed out that nanoGauss magnetic fields can induce early structure formation and may have the potential to achieve the early reionization implied by the WMAP team. They also studied reionization of IGM induced by the dissipation of magnetic energy. In this paper, we thoroughly investigate the role of primordial magnetic fields on the early reionization process. We concentrate on the early structure formation due to the additional power spectrum generated by magnetic fields.         It is known that there exist magnetic fields of several $\\mu$Gauss in most of the galaxies and the clusters of galaxies while the origin of these magnetic fields is still uncertain. Coherence lengths of these magnetic fields are typically $100$ kpc --$1$ Mpc \\citep{magobs}.  Perhaps small seeds of the magnetic fields are produced inside astronomical objects such as stars and AGNs due to the Biermann battery mechanism.  Although the resultant magnetic fields are very weak, those may be amplified by the dynamo process (\\citealt{dynamo}; for a comprehensive review see \\citealt{magreview}).  Eventually these magnetic fields are spread by Supernova winds or AGN jets into IGM. However, for achieving observed values in clusters of galaxies and high redshift galaxies~\\citep{maghighz, clustermag1, clustermag2}, there are difficulties in dynamo theory \\citep{branden-subramanian}.  An alternative to the dynamo scenario is the generation of magnetic fields in the early universe, which we consider in this paper. Magnetic fields can be formed either due to the bubble collisions during the cosmic phase transition such as QCD or electroweak phase transitions, or due to the break of the conformal invariance in the Maxwell theory during the inflation.  For a detailed review, see \\citealt{giova}.  These magnetic fields are formed early enough to make influence on the first structure formation in the universe and the reionization process.  The primordial magnetic fields are constrained by Big-Bang nucleosynthesis and CMB \\citep{bbnconstraint1,bbnconstraint2,diss-j-k-o,mack,lewis,y-i,tashiro}.  The upper limit of the comoving amplitude of magnetic fields is $\\sim 10$ nGauss at the $1$ Mpc scale.    Throughout this paper we use the cosmological parameters measured by the WMAP teams,: the Hubble constant (in the unit of $100 ~{\\rm km s}^{-1} {\\rm Mpc}^{-1}$) $h= 0.71$, the matter density ratio $\\Omega_{\\rm m}= 0.27$ and the baryon density ratio $\\Omega_{\\rm b}= 0.044$ \\citep{spergel}. ",
        "conclusions": "In this paper we investigate the role of the additional density perturbations generated by the primordial magnetic fields on the reionization process in the early universe.  These additional density perturbations may trigger the early structure and star formation. Employing a simple analytic recipe, we estimate the number of ionizing photons emitted from the Population III stars.  We found that the reionization process almost completes by $z \\sim 15$ if the strength of primordial magnetic fields is larger than $0.7$ nGauss and less than $1.5$ nGauss.  Note that we adopt the Gaussian window function to calculate the mass dispersion~(see Eq.~(\\ref{dispersion})).  Different choice of the window function alters the mass dispersion at the magnetic Jeans scale.   Accordingly the magnetic field strength to be required for the early enough reionization is also changed. If we employ the sharp-$k$ window function, for example, the strength should be between $0.5$ and $0.7$ nGauss.  Such magnetic fields are not yet ruled out from current observations, i.e., BBN, CMB temperature anisotropies and polarization, and Faraday rotation of polarized lights from radio sources.  Although the formation process of such primordial magnetic fields is still uncertain, magnetic fields may be naturally generated during the cosmological phase transition \\citep{b-j-paper}.  The extra-power of the matter power spectrum will be directly probed by future observations such as the fluctuations of the Hydrogen 21cm line \\citep{loeb-zaldarriaga} and the substructure of lensing haloes \\citep{dalal-kochanek}.  Moreover, the thermal diffusion process of primordial magnetic fields may cause ionization of IGM even before Population III stars.  A measurement of fluctuations of the 21cm line will be a powerful tool to investigate such pre-reionization~\\citep{h-s-21}."
    },
    "0512/astro-ph0512140_arXiv.txt": {
        "abstract": "Propagation of radio waves in the ultrarelativistic magnetized electron-positron plasma of pulsar magnetosphere is considered. Polarization state of the original natural waves is found to vary markedly on account of the wave mode coupling and cyclotron absorption. The change is most pronounced when the regions of mode coupling and cyclotron resonance approximately coincide. In cases when the wave mode coupling occurs above and below the resonance region, the resultant polarization appears essentially distinct. The main result of the paper is that in the former case the polarization modes become non-orthogonal. The analytical treatment of the equations of polarization transfer is accompanied by the numerical calculations. The observational consequences of polarization evolution in pulsar plasma are discussed as well. ",
        "introduction": "\\subsection{Empirical model of pulsar polarization} The radio emission observed from pulsars is typically characterized by a high percentage of linear polarization. Within the framework of a well-known rotating vector model \\citep{RC69}, the orientation of pulsar polarization reflects the magnetic field geometry in the emission region and therefore the position angle (PA) changes monotonically as the pulsar beam rotates with respect to an observer. The characteristic {\\it S}-shaped swing of PA across the pulse is indeed observed in a number of pulsars. In addition, PA may show abrupt jumps by approximately $90^\\circ$ \\citep[e.g.][]{M75}, testifying to the presence of the two orthogonally polarized modes (OPMs). The early studies of this phenomenon have revealed that for each of the OPMs PA roughly follows the predictions of the rotating vector model \\citep*{B76} and mode changing is a stochastic process \\citep*{C78}. Further on the OPMs have been recognized as a fundamental feature of pulsar radio emission \\citep{BR80,S84a,S84b}. A comprehensive analysis of the observational data has proved an idea of superposed OPMs: At any pulse longitude the radio emission is believed to present an incoherent mixture of the two OPMs, whose intensities vary randomly from pulse to pulse \\citep{Mc98,Mc00}.  The plasma of pulsar magnetosphere does allow two types of non-damping natural waves, the ordinary and extraordinary ones. They propagate in a superstrong magnetic field, generally at not a small angle to the field lines, and, correspondingly, are linearly polarized in orthogonal directions. The electric vector of the ordinary wave lies in the same plane as the wavevector and the ambient magnetic field, while the extraordinary wave is polarized perpendicularly to this plane. The origin of the two types of natural waves is attributed either to the two distinct radio emission mechanisms \\citep[e.g.][]{Mc97} or to the partial conversion of the ordinary mode into the extraordinary one \\citep{P01}. A recently discovered anticorrelation of the OPM intensities \\citep{ES04} favours the latter scenario.  It should be noted, however, that direct identification of the observed superposed OPMs with the natural modes of pulsar plasma faces serious difficulties. First of all, some circular polarization is always present in pulsar radiation. Usually it is much lower than its linear counterpart but not negligible. It has been noticed that the sign of circular polarization is well correlated with PA \\citep{C78}, that is the two OPMs have circular polarization of opposite signs and can be regarded as purely orthogonal elliptical modes. As for the theoretical interpretation, immediate switching to the case of elliptical natural waves \\citep*{Mel77,Mel79,Kunzl98,Mel04a,Mel04b} seems problematic. The ellipticity of the natural waves may result from the gyrotropy of pulsar plasma (caused by difference in the distributions of electrons and positrons), if only the waves propagate quasi-longitudinally with respect to the magnetic field. Although the plasma gyrotropy is very probable, the regime of quasi-longitudinal propagation is not characteristic of pulsar magnetosphere. It can be the case only at some specific locations (most likely, close to the magnetic axis, where the divergence of magnetic field lines is less significant) and cannot account for the elliptically polarized natural waves observed throughout the pulse and over a wide frequency range.  Recent thorough studies of the single-pulse data have introduced further complications into the picture of pulsar polarization. It has been found that the observed fluctuations of the Stokes parameters cannot be explained solely by the pulse-to-pulse variation of the OPM intensities \\citep{Mc04,ES04}. To account for the observations it has been suggested to complement the OPMs with a randomly polarized component of unknown nature \\citep{Mc04}. Alternatively, the same results can be interpreted as a consequence of pulse-to-pulse jitter of both the ellipticity and PA \\citep*{KarIII,ES04}. However, even this generalized picture based on the OPMs with the randomly varying vector of the Stokes parameters appears incomplete. In some cases, the modes are clearly non-orthogonal \\citep{KarIII,Mc04,ES04}. It should be mentioned that the non-orthogonality of polarization modes manifests itself not only in PA, but also in the circular polarization. In particular, the same PA may be accompanied by the circular polarization of any sense, the correlation between PA and $V$ being less perfect at higher frequencies \\citep{KarI,KarIII}.  Diverse and complicated behaviour of the single-pulse polarization as well as its strong frequency dependence motivate our study of the propagation effects in pulsar magnetosphere.  \\subsection{Wave mode coupling in pulsar magnetosphere} Pulsar radio emission is believed to be generated deep inside the tube of open magnetic lines. Then it should propagate through the flow of an ultrarelativistic highly magnetized electron-positron plasma. As a result of propagation effects, polarization of the radio waves may evolve significantly. In the vicinity of the emission region, the ordinary and extraordinary waves are linearly polarized in orthogonal directions and the plasma number density is so large that the geometrical optics regime holds: the natural waves propagate independently, with the electric vectors being adjusted to the orientation of the ambient magnetic field. As the plasma number density decreases with distance from the neutron star, the difference in the refractive indices of the waves decreases as well, and finally the scale length for beating between the modes becomes comparable to the scale length for change in the plasma parameters. Then the polarization planes of the waves have no time to follow the local magnetic field direction, geometrical optics approximation is broken and wave mode coupling starts. Typically this occurs in the outer magnetosphere, at distances of a few tenth of the light cylinder radius. In the region of wave mode coupling, each of the incident natural waves becomes a coherent sum of both natural waves peculiar to the ambient plasma, with the amplitude ratio and phase difference varying along the trajectory. Correspondingly, the ellipticity of the waves increases with distance and the major axis of polarization ellipse is monotonically shifted, so that it no longer reflects the orientation of the ambient magnetic field. Further on, as the plasma density decreases considerably, the waves decouple from the plasma and propagate just as in vacuum, preserving their elliptical polarization. Therefore the process of wave mode coupling is usually called polarization-limiting effect.  This effect has long been used to explain the origin of circular polarization in pulsar radio emission \\citep{CR79,RR90,LP99}. The numerical tracings of polarization evolution in pulsar plasma have demonstrated that the mode coupling effect is strong enough to have marked observational consequences \\citep{PL00,P01,P03a}.  Polarization evolution of radio waves in pulsar plasma differs significantly from the evolution in the interstellar medium. In contrast to the case of Faraday rotation, within the pulsar magnetosphere the natural waves propagate quasi-transversely with respect to the magnetic field and have linear polarization. This rather corresponds to the Cotton-Mouton birefringence (or so called generalized Faraday rotation, in terms of the paper by \\citealt{KM98}). Another important distinctive feature of polarization evolution in pulsar magnetosphere is that it takes place in an essentially inhomogeneous medium. The magnetic field of a pulsar has approximately dipolar structure, and furthermore, because of continuity of the plasma flow in the tube of open magnetic lines, the plasma number density decreases rapidly with distance. Thus, the character of polarization evolution in pulsar magnetosphere is quite specific, though the underlying physics of birefringence is certainly the same.  As a result of the mode coupling effect, the outgoing waves acquire purely orthogonal elliptical polarization, matching the empirical representation of superposed elliptical OPMs. It is important to note that the degree of circular polarization of the modes and their shift in PA are related to each other, both being determined by the parameters of the plasma flow in the region of wave mode coupling \\citep{P03a}. Hence, the observed pulse-to-pulse variations in the ellipticity and PA of the OPMs can be attributed to the fluctuations in pulsar plasma. Besides that, the propagation origin of pulsar polarization should imply a correlation between the values of the ellipticity and PA of the OPMs at a given pulse longitude. An evidence for such a correlation has recently been found in \\citet{E04} \\citep[for more details see][]{P06}. Thus, the mode coupling effect can account for a number of important features of the observed single-pulse polarization.  At the same time, the question as to the origin of non-orthogonality of polarization modes still remains open. The observational manifestations of this phenomenon have recently been reported in a number of papers \\citep[e.g.][]{ES04,Mc04,Rank03,KarIII,KarI}. In the present paper, we concentrate on a more detailed treatment of polarization evolution in pulsar magnetosphere, which, in particular, explains non-orthogonality of the modes.  \\subsection{Statement of the problem} Polarization evolution in pulsar magnetosphere has previously been considered in the approximation of a superstrong magnetic field. It means that in the rest frame of the plasma flow the radio frequency, $\\omega^\\prime$, is much less than the electron gyrofrequency, $\\omega_H$. In other words, the radius of cyclotron resonance, where $\\omega^\\prime =\\omega_H$, has been assumed to be infinitely large. At the conditions relevant to pulsar magnetosphere, the radius of cyclotron resonance is often somewhat larger than that of the mode coupling region, but generally these quantities are of the same order of magnitude \\citep[][see also equation (15) below]{B86}. Therefore it is reasonable to inspect the role of the cyclotron resonance in the evolution of pulsar polarization.  In application to pulsars, the cyclotron absorption has been considered in a number of papers \\citep{BS76,LP98,P02,P03b} and found efficient, especially in case of small pitch-angles of the absorbing particles. In these papers, it is assumed that the resonant photons interact with a system of absorbing particles rather than with the plasma. The main motivation for such an assumption is that the resonance region lies in the outer magnetosphere, where the plasma number density is small enough. Indeed, in case of pulsars, taking into account the plasma effect on the process of cyclotron absorption introduces only small corrections to the absorption coefficients of the natural modes and does not change the total intensities of outgoing radiation considerably \\citep{LP98}. At the same time, cyclotron absorption in the plasma may markedly affect polarization evolution of radio waves. The contribution of cyclotron absorption, though not large quantitatively, may appear comparable to that of the mode coupling effect, modifying the final polarization of a pulsar drastically.  Given that the regime of geometrical optics is still valid in the region of cyclotron resonance, the ordinary and extraordinary waves are absorbed independently, with the absorption coefficients being slightly different. Note that this difference is purely the plasma effect, not characteristic of a simple system of absorbing particles. Beyond the resonance region, the waves propagate in the weakly magnetized plasma with rapidly decreasing number density and finally suffer the mode coupling. It should be noted that in case of weakly magnetized plasma the limiting polarization differs substantially from that for the superstrong magnetic field.  Given that the natural waves pass through the coupling region before the cyclotron resonance, in the resonance region each of the waves presents a coherent mixture of the two natural waves. Since for these constituents absorption is not identical, the wave polarization changes considerably. It is important to note that for the two incident waves polarization evolution is not the same, since they contain different portions of the ordinary and extraordinary waves. As a result, polarization states of the outgoing waves are non-orthogonal.  In the present paper, we concentrate on the analytical consideration of polarization transfer in pulsar plasma, taking into account the effect of cyclotron resonance. The plan of the paper is as follows. In Sect. 2, the main equations are derived, which describe the evolution of the wave fields in the inhomogeneous hot plasma embedded in the magnetic field. The basic numerical estimates are also given there. In Sect. 3, we solve the equations of polarization transfer in the two limiting cases, when the resonance region is well below and well above the mode coupling region, respectively. Section 4 contains the results of numerical tracing of polarization evolution. In Sect. 5, the observational consequences of polarization transfer in pulsar magnetosphere are discussed. Section 6 contains a brief summary. A statistical model of single-pulse polarization based on the propagation effects studied will be developed in the forthcoming paper \\citep{P06}. ",
        "conclusions": "We have studied polarization transfer in the hot magnetized plasma of pulsars. In the present consideration, we have assumed the non-gyrotropic plasma, with the identical distributions of the electrons and positrons, and the small pitch-angles of the particles. Proceeding from the Maxwell's equations, we have derived the set of equations describing the evolution of the Stokes parameters of the original linearly polarized natural waves. These equations have been solved analytically and the results have been confirmed by numerical calculations.  The polarization evolution of the waves has been found significant. The polarization characteristics change on account of the wave mode coupling and cyclotron absorption. In cases when the region of coupling lies well above and well below the resonance region, the resultant polarization is qualitatively distinct. If the waves pass through the region of wave mode coupling first, they acquire the elliptical polarizations purely orthogonal at the Poincare sphere. Further on, in the resonance region, they become non-orthogonal. If the waves enter the resonance region before coupling, cyclotron absorption does not affect their polarization states and suppresses the total intensities only, the extraordinary wave being absorbed somewhat more efficiently. Further on the waves suffer mode coupling in the limit of weak magnetic field, which differs from that in the strong field. Firstly, the total change of polarization parameters is much less. Besides that, the resultant circular polarization has the opposite sense.  The observational consequences of polarization transfer in pulsar magnetosphere can be summarized as follows. Because of cyclotron absorption, at large enough $\\eta$ one mode can markedly dominate another one, in which case one can expect strongly polarized profiles, with the same mode dominating throughout the average pulse and in most of the individual pulses. This is indeed characteristic of several pulsars. In some cases, the total-intensity profiles of the strongly polarized pulsars are classified as partial cones, with one of the conal components being absent. This seems to be an additional argument in favour of strong cyclotron absorption in these pulsars. At present it is not known exactly what of the natural waves actually dominates in the strongly polarized pulsars. Our consideration favours the dominance of the ordinary mode.  If one consider $\\eta$ as a function of radio frequency, with $\\eta\\approx 1$ at the intermediate frequencies, in the low-frequency range $\\eta\\geq 1$ and one can expect strong polarization of pulsar profiles, whereas at high enough frequencies there may be an increase of the degree of circular polarization because of non-orthogonality of the modes. The non-orthogonality of the outgoing waves is of a crucial importance for the statistical model of the individual-pulse polarization. This will be studied in detail in the forthcoming paper."
    },
    "0512/astro-ph0512189_arXiv.txt": {
        "abstract": "{ Spherically symmetric (1D) and two-dimensional (2D) supernova simulations for progenitor stars between 11$\\,M_\\odot$ and 25$\\,M_\\odot$ are presented, making use of the \\textsc{Prometheus/Vertex} neutrino-hydrodynamics code, which employs a full spectral treatment of neutrino transport and neutrino-matter interactions with a variable Eddington factor closure of the ${\\cal O}(v/c)$ moments equations of neutrino number, energy, and momentum. Multi-dimensional transport aspects are treated by the ``ray-by-ray plus'' approximation described in Paper~I. We discuss in detail the variation of the supernova evolution with the progenitor models, including one calculation for a 15$\\,M_\\odot$ progenitor whose iron core is assumed to rotate rigidly with an angular frequency of 0.5 rad$\\,$s$^{-1}$ before collapse. We also test the sensitivity of our 2D calculations to the angular grid resolution, the lateral wedge size of the computational domain, and to the perturbations which seed convective instabilities in the post-bounce core. In particular, we do not find any important differences depending on whether random perturbations are included already during core collapse or whether such perturbations are imposed on a 1D collapse model shortly after core bounce. Convection below the neutrinosphere sets in 30--40$\\,$ms after bounce at a density well above $10^{12}\\,$g$\\,$cm$^{-3}$ in all 2D models, and encompasses a layer of growing mass as time goes on. It leads to a more extended proto-neutron star structure with reduced mean energies of the radiated neutrinos, but accelerated lepton number and energy loss and significantly higher muon and tau neutrino luminosities at times later than about 100$\\,$ms after bounce. While convection inside the nascent neutron star turns out to be insensitive to our variations of the angular cell and grid size, the convective activity in the neutrino-heated postshock layer gains more strength in better resolved models. We find that low ($l = 1,\\,2$) convective modes, which require the use of a full 180 degree grid and are excluded in simulations with smaller angular wedges, can qualitatively change the evolution of a model. In case of an $11.2\\,M_\\odot$ star, the lowest-mass progenitor we investigate, a probably rather weak explosion by the convectively supported neutrino-heating mechanism develops after about 150$\\,$ms post-bounce evolution in a 2D simulation with 180 degrees, whereas the same model with 90 degree wedge fails to explode like all other models. This sensitivity demonstrates the proximity of our 2D calculations to the borderline between success and failure, and stresses the need to strive for simulations in 3D, ultimately without the constraints connected with the axis singularity of a polar coordinate grid. ",
        "introduction": "The mechanism by which massive stars explode is still unclear. State-of-the-art models with a spectral treatment of the neutrino transport by solving the Boltzmann equation or/and its moments equations agree in the finding that in spherical symmetry (1D) neither the prompt bounce-shock mechanism nor the delayed neutrino-driven mechanism lead to explosions for progenitors more massive than about 10$\\,M_\\odot$ (e.g., \\citealp{ramjan02}; \\citealp{liemez01,liemes04}; \\citealp{thobur03}; \\citealp{sumyam05}). Previous two-dimensional (2D) simulations (e.g., \\citealp{herben94}, \\citealp{burhay95}, \\citealp{janmue96}, \\citealp{fry99}, \\citealp{fryheg00}) and three-dimensional (3D) models (\\citealp{frywar02,frywar04}) show the importance of convective overturn in the neutrino-heating layer behind the stalled supernova shock, which can enhance the energy transfer from neutrinos to the stellar matter and thus cause ``convectively supported neutrino-driven explosions''. These multi-D models, however, employed radical simplifications of the treatment of neutrinos, mostly by grey diffusion or in a parametric way as heating terms. Concerns about the reliability of such approximations of crucial physics in studies of the supernova explosion mechanism were expressed by \\cite{mezcal98:ndconv}.  Also the influence of convective activity inside the nascent neutron star, i.e. below the neutrinosphere, on the explosion mechanism has long been a matter of debate and requires further studies. The Livermore group (\\citealp{wilmay88,wilmay93}) obtained explosions in their basically 1D models by assuming that so-called neutron-finger mixing instabilities exist in the newly formed neutron star, which accelerate the energy transport from the neutron star interior to the neutrinosphere. Thus the neutrino luminosities are boosted and the neutrino heating behind the supernova shock is enhanced. The analysis by \\cite{brudin96} and more recently by \\cite{brural04}, however, has demonstrated that neutrino diffusion leads to lepton number equilibration between perturbed fluid elements and their surroundings that is faster than assumed by \\cite{wilmay88,wilmay93}. Therefore neutron fingers are unlikely to occur in the supernova core. \\cite{brural04} instead discovered a new mode of doubly-diffusive instability, which they termed ``lepto-entropy fingers'' and which is also associated with neutrino-mediated thermal and lepton diffusion. The importance of this phenomenon during the early, critical phases of the explosion, however, was recently questioned by \\cite{desbur05} because of its slow growth compared to Ledoux convection. The latter, in turn, was predicted to play a role in supernovae on grounds of 1D models of the neutrino cooling phase of nascent neutron stars. A Ledoux-type of convection was indeed found to be present during the first second after neutron star formation in 2D hydrodynamic simulations by \\cite{kei97}, \\cite{keijan96}, \\cite{jankei98} and \\cite{jankif01}. The latter simulations, however, considered only the proto-neutron star without self-consistently following its feedback with the environment of the supernova core. Moreover, a grey, flux-limited equilibrium ``ray-by-ray'' diffusion code for the neutrino transport was used, with strong simplifications in the description of the opacities.  Only recently multi-dimensional simulations of stellar core collapse and post-bounce evolution with a spectral treatment of the neutrino transport have become possible (\\citealp{burram03,burram06:I}; \\citealp{livbur04}; \\citealp{walbur05}; \\citealp{swemyr05a,swemyr05b}; \\citealp{burliv06}). Although these current approaches are the first steps of removing the severe deficiencies of the previous generation of multi-dimensional models, all of them still contain approximations of various, and different, aspects in the treatment of 2D transport. \\cite{swemyr05a,swemyr05b}, for example, use a flux-limited diffusion description, an approximation also made by \\cite{walbur05} and \\cite{desbur05}, who in addition solve the transport for all neutrino energy groups independently. In contrast, \\cite{burram06:I} have developed a ``ray-by-ray plus'' approximation based on a variable Eddington factor solver for the coupled set of neutrino moments equations and Boltzmann equation, including a full coupling of the energy bins by neutrino reactions and by Doppler and gravitational redshift effects.  The approximations employed by the different groups are diverse and might hamper a detailed quantitative comparison of the results in the near future, and might constrain such efforts to a purely qualitative level. Eventually it will be necessary to test and possibly replace the current approximations by a more rigorous solution of the transport problem in the five- or six-dimensional phase space and in a relativistic framework, once the corresponding codes have become available and the necessary substantial increase of computer power has happened (\\citealp{car04}, \\citealp{caretal05}).  Here we present results obtained with the multi-dimensional neutrino-hydrodynamics code \\textsc{MuDBaTH}, which is the ``ray-by-ray plus'' implementation of the \\textsc{Prometheus/Vertex} code described in detail in Buras et al.\\ (2005; Paper~I). In continuation of our previous work (\\citealp{burram03,burram06:I}), where also a broader introduction into the status of the field and its open questions is provided, we present here 1D simulations for nine different progenitor stars with masses between 11.2$\\,M_\\odot$ and 25$\\,M_\\odot$, and compare them with 2D simulations for three of these stars. The core collapse and post-bounce evolution of these models was followed until nearly 300$\\,$ms after shock formation.  Using a state-of-the-art treatment of spectral neutrino transport for hydrodynamical supernova simulations, the main goals of our work are: \\begin{itemize} \\item We compare 1D and 2D models in order to obtain quantitative information about the influence of convection below the neutrinosphere on the neutrino emission, the evolution, and the structure of the nascent neutron star. We analyse the influence of proto-neutron star convection on the conditions in the neutrino-heating layer behind the shock, and assess quantitatively the impact of convective activity in the postshock layer on the possibility for reviving the stalled shock and for getting a delayed explosion. \\item We also study the differences of convection between a non-rotating and a rotating 15$\\,M_\\odot$ model, whose iron core spins rigidly before collapse with a period of about 12 seconds, leading to an ``extreme'' period of the settled neutron star of the order of 1$\\,$ms. In addition, we investigate the effects of low-mode (dipolar, $l = 1$, or quadrupolar, $l = 2$) hydrodynamic instabilities during the post-bounce evolution of an 11.2$\\,M_\\odot$ star, comparing simulations with a full 180$\\degr$ polar grid and simulations which are contrained to a $\\sim\\,$90$\\degr$ equatorial wedge (with periodic angular boundary conditions), thus preventing the development of such low modes in the pattern of the fluid flow. \\item Moreover, we perform tests for the influence of (i) the numerical resolution, in particular in the lateral direction of our 2D polar grid, (ii) of the chosen size of the angular wedge, and (iii) of the way in which we perturb our models to initiate the growth of convective instabilities, i.e., whether we follow a perturbed 2D model through core collapse and core bounce, or whether we map a 1D model to the 2D grid shortly after bounce, imposing random perturbations at that time. \\end{itemize}   The paper is structured in the following way. Main results of our 1D supernova simulations for the chosen set of progenitor models --- whose basic properties are compared in Appendix~\\ref{app:progs} --- will be discussed in Sect.~\\ref{sec:1d_prog}, supplemented with more details in Appendix~\\ref{app:1Dresults}. The 2D models will be presented in Sect.~\\ref{sec:p2_tdm}, with an analysis of the effects of convection in the forming neutron star in Sect.~\\ref{sec:pnsc}, a discussion of convection in the neutrino-heating layer in Sect.~\\ref{sec:hbc}, a description of our full 180$\\degr$ model in Sect.~\\ref{sec:full_star}, of the rotating model in Sect.~\\ref{sec:rot_star}, and of neutrino emission anisotropies in Sect.~\\ref{sec:2dneutrinos}. Most of our 2D simulations without rotation were started from 1D collapse models only shortly after bounce, at which time small random perturbations were imposed to seed convective instabilities. Since the adequacy of such an approach may be disputed, we also performed simulations where the collapse phase was followed in two dimensions. This allowed us to investigate the growth of inhomogeneities during infall and to assess the possible influence of that on the growth of convection after bounce (Appendix~\\ref{app:p_coll}). A summary and conclusions will follow in Sect.~\\ref{sec:concl}. Appendix~\\ref{app:PNSstr} contains a linear analysis of the structural changes of the proto-neutron star which can be expected as a consequence of convection below the neutrinosphere, and Appendix~\\ref{app:ns_mix} introduces a simple mixing scheme by which we achieved to reproduce in 1D simulations the main effects of proto-neutron star convection as observed in our 2D models. ",
        "conclusions": "\\label{sec:concl}  We have presented results of a series of core-collapse and post-bounce simulations for different progenitor stars between 11.2$\\,M_\\odot$ and 25$\\,M_\\odot$, comparing 2D (axially symmetric) with 1D (spherically symmetric) calculations. Doing so, our main goals were (i) investigating the differences between convection in progenitors with different masses, (ii) exploring the effects of convection below the neutrinosphere (``PNS convection'') on the proto-neutron star structure, its neutrino emission, and the neutrino heating-layer behind the shock, (iii) investigating the role of hydrodynamic instabilities that affect the stalled accretion shock, i.e.\\ convective overturn in the neutrino-heated ``hot bubble'' layer (``HB convection'') and global low-mode nonradial instability of the accretion shock (termed SASI by \\citealp{blomez03} and possibly caused by the action of an advective-acoustic cycle according to \\citealp{fog01,fog02}), (iv) studying the effects of rotation, and (v) testing the influence of numerical aspects like the grid resolution, size of the angular wedge, and magnitude of seed perturbations for convection. Since our 2D neutrino-hydrodynamics code is a direct descendant of our 1D \\textsc{Prometheus/Vertex} code, it is particularly well suited for performing such comparisons of 1D and 2D supernova models.  Convection inside the proto-neutron star starts 30--40$\\,$ms after bounce in all of our 2D models and encompasses a layer growing in mass until the end of our simulations (which were typically terminated about 250$\\,$ms after bounce). It leads to a more extended neutron star than in the 1D simulations with lower temperatures at the neutrinosphere. For this reason the mean energies of the neutrinos emitted from the neutrinosphere are reduced (up to 10\\% after 200$\\,$ms of PNS convection). Despite the larger radiating surface, the lower neutrinospheric temperatures also cause a slight reduction of neutrino luminosities during the first 150$\\,$ms after bounce. This holds in particular for $\\bar\\nu_{\\mathrm{e}}$, because the convective transport of lepton number maintains a higher electron degeneracy in the neutrinospheric region and accelerates the lepton number loss compared to 1D simulations. Only at $t\\ga 150\\,$ms after bounce, convectively enhanced energy transport in the nascent neutron star also leads to increased energy loss, and the luminosities of heavy-lepton neutrinos become significantly (15\\%--20\\%) higher than in the spherical models.  PNS convection of the kind found in our simulations leads to a slightly {\\em reduced} total energy deposition in the gain layer mainly because of the lower average energies of the radiated $\\nu_{\\mathrm{e}}$ and $\\bar\\nu_{\\mathrm{e}}$. Since the effects of convection below the neutrinosphere are hard to disentangle from those of hydrodynamic instabilities in the neutrino-heating layer behind the stalled shock, we developed a simple ``mixing algorithm''. It allowed us to reproduce all major effects of PNS convection in 1D simulations and thus to separate them from the consequences of multi-dimensional fluid flow in the postshock layer and to arrive at the above conclusion.  Convective overturn in the neutrino-heating layer remained rather weak in case of the 15$\\,M_\\odot$ and 20$\\,M_\\odot$ progenitors. The main reason for that is the rapid contraction of the accretion shock after its maximum expansion. This causes the gain layer to be very narrow and the infall velocities of the gas ahead and behind the shock to be very high. As a consequence, the advection timescale of the gas through the gain layer is very short compared to the typical neutrino-heating timescale. Buoyancy forces therefore hardly achieve bubble rise in the flow of gas accreted from the shock to the gain radius. As suggested by \\cite{jankei98} and \\cite{jankif01} and verified by \\cite{thoqua05}, the ratio of the advection timescale to the neutrino-heating timescale, $\\tau_{\\mathrm{adv}}/\\tau_{\\mathrm{heat}}$, turned out to be a useful diagnostic parameter to measure the proximity of a model to a neutrino-driven explosion. A necessary condition for an explosion is that the timescale ratio rises above unity for a time interval of at least the neutrino-heating timescale. In case of the 15$\\,M_\\odot$ and 20$\\,M_\\odot$ models, HB convection increases the heating rate and the timescale ratio to values only slightly larger than in the 1D simulations, but still roughly a factor of two below the critical limit. We therefore found explosions of these stars neither in spherical symmetry (in agreement with \\citealp{liemes02}, \\citealp{thobur03}, \\citealp{sumyam05}) nor in 2D.  Also rotation did not change this negative outcome. We studied one 15$\\,M_\\odot$ model with pre-collapse rigid iron core rotation of $\\sim\\,$12$\\,$s period, which leads to a neutron star with a spin period of about 1$\\,$ms, if the angular momentum of the core after collapse is conserved. Rotation of this size is probably on the extreme side of what can be expected for the cores of ``normal'' supernovae, which are supposed to give birth to neutron stars with an initial period of 10$\\,$ms or more (see the discussions in \\citealp{hegwoo05} and \\citealp{ottbur06}). Our simulations reveal a number of important differences of the rotating model compared to its non-rotating counterparts. The proto-neutron star develops an eccentricity of more than 0.6 until we stopped the simulation at nearly 300$\\,$ms after bounce. At this time the luminosities of the radiated neutrinos are significantly smaller (10--20\\%) and their mean energies up to 2$\\,$MeV lower than in the non-rotating 2D model, because the equatorially more extended neutrinosphere is significantly cooler and energy is stored in rotation instead of being released by neutrinos. Despite the clear oblateness of the proto-neutron star, its rotation-induced emission anisotropy is very small. Nevertheless, rotation has a favorable influence on the conditions and parameters which determine neutrino-driven explosions. Centrifugal forces stabilize the accretion shock at larger radii, increase the advection timescale of the postshock gas significantly, and thus allow for a layer of well developed, strong convective overturn activity behind the shock. Because more mass stays in the gain layer for a longer time, the total energy deposition rate behind the shock is higher at later post-bounce times ($t \\ga 130\\,$ms p.b.) than in the non-rotating models. In spite of these healthy effects, however, the timescale ratio $\\tau_{\\mathrm{adv}}/\\tau_{\\mathrm{heat}}$ remains still well below unity ($\\tau_{\\mathrm{adv}}/\\tau_{\\mathrm{heat}} \\la 0.6$).  Even without rotation postshock convection becomes violent in case of the 11.2$\\,M_{\\odot}$ star. The shock in this model is able to stay longer at large radii than in the more massive stars. This is due to the fact that the rather low-mass progenitor has a steeper density decline at the transition to the Si+O layer, which leads to a rapid decrease of the mass accretion rate of the shock at about 90$\\,$ms after bounce. This allows the shock to reexpand in adjusting to the situation of reduced ram pressure. The increased advection timescale gives convection the possibility to gain strength and thus to support the shock at a much larger radius than in the corresponding 1D model. Also the total neutrino heating rate behind the shock and the efficiency of net neutrino-energy transfer to the gas in the gain layer is higher by up to a factor of two. The timescale ratio $\\tau_{\\mathrm{adv}}/\\tau_{\\mathrm{heat}}$ approaches unity and remains close to --- but slightly below --- this threshold until the end of our simulations. The 11.2$\\,M_{\\odot}$ model computed with a 90$\\degr$ lateral wedge therefore lingers at the border to success.  Such a situation is extremely sensitive to relatively little changes. We saw this when we repeated the simulation with a full 180$\\degr$ grid instead of using the wedge around the equator. While PNS convection turned out not to depend on the wedge size, convective activity in the neutrino-heating layer can change significantly when the available degrees of freedom are not constrained by periodic boundary conditions of a 90$\\degr$ equatorial wedge and therefore low-mode deformation of dipolar ($l = 1$) and quadrupolar ($l = 2$) character is allowed for. Convection becomes sufficiently strong so that the accretion shock continues to expand. This ensures that the effective advection timescale does not decrease after it has reached its maximum. At $t \\ga 140\\,$ms after bounce, the timescale ratio $\\tau_{\\mathrm{adv}}/\\tau_{\\mathrm{heat}}$ then becomes larger than unity, thus further improving the conditions for efficient energy deposition by neutrinos in the postshock layer. After about 180$\\,$ms of post-bounce evolution the total energy in the gain layer becomes positive and continues rising because the mass in the gain layer and the energy per nucleon grow. The model has passed the critical threshold and is on its way to explosion. A closer inspection of the involved energies shows that this explosion is powered by neutrino heating.  This qualitative difference of the outcome of 2D simulations with 90$\\degr$ and 180$\\degr$ grids is another confirmation of the proximity of our 2D simulations, and in particular of the 11.2$\\,M_{\\odot}$ case, to a success of the convectively supported neutrino-driven mechanism. Together with the recent models for stars in the 8--10$\\,M_\\odot$ range with O-Ne-Mg cores, which explode even in spherical symmetry (\\citealp{kitjan06}), our current results seem to indicate that the neutrino-heating mechanism is viable at least for stars near the low-mass end of supernova progenitors.  The sensitivity to numerical variations, however, also stresses the need to remove some of the shortcomings and limitations of axially symmetric simulations. One must suspect that in 3D simulations morphological differences of the structures (plumes instead of azimuthal tori), different growth rates of instabilities, or additional degrees of freedom (e.g. triaxial asymmetries and vortex motion caused by Coriolis forces) might lead to sizable quantitative differences which could be crucial when collapsing stellar cores are close to the threshold for explosion. Also the existence of the polar axis of a spherical or cylindrical coordinate grid is a potential source of numerical uncertainties, because it is a coordinate singularity which is impenetrable for approaching fluid flow and thus defines a preferred grid direction.  Our results therefore suggest the need to strive for 3D simulations, preferentially without the disadvantages connected with the polar grid axis. The importance of low-mode convection or low-mode hydrodynamical instabilities as suggested by our results implies that such simulations will have to be done for the full star and cannot be contrained to a limited wedge."
    },
    "0512/astro-ph0512376_arXiv.txt": {
        "abstract": "We use the Sloan Digital Sky Survey (SDSS) spectroscopic sample to constrain the projected radial distribution of satellites around isolated $\\sim L_{\\ast}$ galaxies.  We employ mock galaxy catalogs derived from high-resolution cosmological simulations to investigate the effects of interloper contamination and show that interlopers significantly bias the estimated slope of the projected radial distribution of satellites. We also show that the distribution of interlopers around galaxies is expected to be non-uniform in velocity space because galaxies are clustered and reside in crowded environments. Successful methods of interloper contamination correction should therefore take into account environments of the host galaxies.  Two such  new methods are presented and the most reliable of them is used to correct for interloper contamination in analyses of the SDSS galaxy sample. The best fit power-law slope of the interloper-corrected surface density distribution of satellites, $\\Sigma(R)\\propto R^{\\alpha}$, in the volume-limited SDSS sample is $\\alpha \\simeq -1.7 \\pm 0.1$, independent of the galaxy and satellite luminosities. Comparison with $\\Lambda$CDM simulations shows that the radial distribution of the SDSS satellites is more concentrated than that of subhalos around galaxy-sized halos, especially at $R<100h^{-1}$~kpc. The predicted dark matter radial distribution is somewhat more concentrated than the profile of  the SDSS satellites, but the difference is not statistically significant for our sample. ",
        "introduction": "\\label{sec:intro}  In the Cold Dark Matter (CDM) paradigm, satellite galaxies are expected to be associated with the dark matter subhalos -- halos which lie within the virial radius of a larger halo -- ubiquitous in the cosmological CDM simulations. The abundance and radial distribution of satellite galaxies can therefore serve as a useful test of CDM galaxy formation models, constraining the relation between galaxies and subhalos. In addition, satellite dynamics can provide useful constraints on the total mass distribution in galactic halos \\citep[e.g.,][]{zaritsky_white94,zaritsky_etal97,prada_etal03,vandenbosch_etal04,conroy_etal05}. This, however, requires a good understanding of how the spatial distribution and kinematics of satellites and dark matter are related.  Many recent studies based on numerical simulations have shown that the radial distribution of subhalos in cluster-sized systems is less concentrated than that of dark matter in the inner $\\approx 20-50\\%$ of the virial radius of host halos, but approximately follows the dark matter distribution at larger radii \\citep{ghigna_etal98,colin_etal99,ghigna_etal00, springel_etal01,delucia_etal04,diemand_etal04,gao_etal04,nagai_kravtsov05}. Theoretical predictions for galaxy distributions in clusters have also been accompanied by rapidly improving observational measurements \\citep[e.g.,][]{lin_etal04,hansen_etal05,collister_lahav05,yang_etal05,coil_etal05}, which also find concentrations of galaxy radial profiles lower than the concentrations expected for the matter distribution of their parent halos.  The observed distribution of satellite galaxies in galactic halos has been studied less extensively.  The Local Group dwarf population is more radially concentrated than subhalos in dissipationless numerical simulations \\citep{kravtsov_etal04b,taylor_etal04,willman_etal04}, a bias that is likely related to the physics of the formation of the smallest dwarf galaxies \\citep{kravtsov_etal04b,diemand_etal05}. The known population of the Local Group satellites is, however, quite small compared to the expected population of CDM subhalos \\citep{klypin_etal99,moore_etal99}. Moreover, the strong radial bias exhibited by the faint Milky Way satellites is not expected to apply to the brighter satellites (such as, for example, the Magellanic Clouds).  More accurate, statistical constraints on the satellite distribution can be obtained by using galaxy redshift surveys. Several early studies attempted to constrain the small-scale galaxy correlation function by estimating the surface density of objects projected near galaxies, $\\Sigma(R) \\propto R^{\\alpha}$, finding slopes ranging from $\\alpha=-0.5$ to $-1.25$ \\citep{lake_tremaine80,phillipps_shanks87,vader_sandage91,lorrimer_etal94,smith_etal04,madore_etal04}.  Recently, the availability of large galaxy redshift surveys has allowed construction of large statistical samples of parent galaxies and satellites with well defined selection criteria.  The large sample sizes and redshift information make it possible to understand the biases and completeness of the sample.  In addition, isolation criteria for the primaries can be introduced in order to reduce the interloper contamination and simplify the interpretation of results. \\citet{vandenbosch_etal05} use mock galaxy redshift samples derived from large cosmological simulations to develop an iterative method of interloper rejection for the Two Degree Field Galaxy Redshift Survey (2dFGRS) and find that the data is generally consistent with the dark matter profile at large projected radii, but conclude that incompleteness of close pairs in the survey prevent strong constraints. In an independent analysis, \\citet{sales_lambas05} account for the close-pair bias in the data by estimating completeness with control samples of objects that are not physically bound to the primaries.  They estimate the power-law slope of the satellite distribution to be $\\alpha = -0.96 \\pm 0.03$ for projected radii between 20 and 500 $h^{-1}$ kpc with a significant dependence on morphological type of the parent galaxies ($\\alpha\\approx -1.1$ for the early type, and $\\approx -0.7$ for the late type galaxies). Note, however, that these values of $\\alpha$ are obtained without any correction for interlopers.  Given that the satellite distribution can be directly probed only in projection, with only limited information about positions of likely satellites in three dimensions, one has to worry about contamination by interlopers, the objects that are not true satellites but are simply close to the parent due to projection. Unfortunately, in practice it is often tricky to estimate and correct for the interloper contamination. This is especially difficult if the redshift information is absent as was the case in the earliest studies of the satellite distribution. However, even in studies in which redshift information is available, the interlopers are often neglected \\citep[e.g.,][]{sales_lambas05}. Nevertheless, as we show in this paper, the effect of interlopers must be corrected for in order to obtain an unbiased measurement of the satellite projected radial distribution.  The simplest assumption one can make is that the surface density of interlopers is uniform. The interloper contamination can then be estimated by sampling the environments around random points in the field. This method thus presumes that the volume around a random point on the sky and in redshift space contains a representative density of interlopers.  However, bright galaxies are strongly correlated in space and thus can be expected to be preferentially located in crowded environments. One may suspect, then, that the random points method can underestimate the interloper number density around real galaxies. Therefore, more sophisticated methods, which sample interlopers in the environments similar to those of the primary galaxies, need to be developed.  In this study, we develop two new methods to estimate the contribution of interlopers to the surface density of satellites, which take into account the clustering of parent galaxies.  We use cosmological simulations to test different methods of interloper subtraction and present detailed discussion of their strengths and weaknesses.  We show that interloper contamination can significantly bias measurements of the projected radial distribution of satellite galaxies. Proper interloper subtraction is, therefore, a must in studies of the radial distribution of satellites. We use the Sloan Digital Sky Survey (SDSS) spectroscopic sample to measure the projected radial distribution of satellites around nearby bright galaxies, corrected for interlopers. We compare the result to the predictions for the dark matter and subhalo distribution in the $\\Lambda$CDM cosmology.  The paper is organized as follows.  In \\S~\\ref{sec:testing} we discuss the interloper contamination and different methods of interloper subtraction, testing each of them using mock satellite catalogs derived from cosmological simulations.  We then describe our SDSS spectroscopic galaxy sample and the selection of primaries and satellites in \\S~\\ref{sec:sdss}. In \\S~\\ref{sec:data} we derive the interloper-corrected surface density profile of satellites in volume-limited and flux-limited SDSS samples, and their subsamples, and compare results to the $\\Lambda$CDM cosmological simulations. We also discuss comparisons to simulations results. Our main results and conclusions are summarized in section \\S\\ref{sec:conclusions}. Throughout this paper, we assume flat $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$ and $h=0.7$.   \\begin{table*}[t] \\begin{center} \\caption{Selection \\& Isolation Criteria for Test Samples\\label{tab:select_test}} \\begin{tabular}{lccc} \\tableline \\\\ \\multicolumn{1}{l}{Parameters} & \\multicolumn{1}{c}{Test Sample 1 (TS1)} & \\multicolumn{1}{c}{Test Sample 2 (TS2)} & \\multicolumn{1}{c}{Test Sample 3 (TS3)} \\\\ \\\\ \\tableline \\\\ & \\multicolumn{1}{l}{}\\\\ Constraints on primaries& $V_{\\rm max}$ = 100-150, ..., 300-350 km s$^{-1}$& ~~$M_{r}<$ -20~~ & ~~$M_{r}<$ -20~~\\\\  Satellite objects & DM particles & DM particles & satellite galaxies\\\\  Isolation criteria: \\\\ ~~~Size difference & $V_{\\rm max} > 0.5 V_{\\rm max}^{\\rm pri}$ & $V_{\\rm max} > 0.5 V_{\\rm max}^{\\rm pri}$ & $\\Delta M_{r} < 2$ \\\\ ~~~Minimum projected distance, $\\Delta R (h^{-1}$ Mpc)& 0.5 & 0.5 & 0.5\\\\ ~~~Minimum velocity separation, $\\Delta V$ (km ${\\rm s^{-1}}$)& 1000 & 1000 & 1000 \\\\  Satellite sample criteria: \\\\ ~~~Magnitude difference from the primary & ----- & ----- & $\\Delta M_{r} > 2$ \\\\ ~~~Maximum projected distance, $\\delta r (h^{-1}$ Mpc)& 0.6 & 0.6 & 0.6\\\\ ~~~Maximum velocity separation, $\\delta v$ (km ${\\rm s^{-1}}$)& 500 & 500 & 500, 1000\\\\  Number of isolated primaries&380, 289, 236, 143, 89& 728 & 728  \\\\  Number in satellite sample&7475, 11165, 16614,15354, 14899 & 50608 & 343, 401\\\\  Limiting magnitude $M_{r}$& ------ & ------ & -18\\\\ \\\\ \\tableline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec:conclusions}  Modern large galaxy redshift surveys allow one to study the distribution of satellites around galaxies and clusters with unprecedented statistical power, while controlling biases and completeness in a systematic way.  In addition, redshift information can be used to select galaxies from relatively non-crowded environments and to account for interloper contamination in a rigorous way.  Cosmological simulations are also sufficiently mature and allow systematic tests of the interloper subtraction algorithms.  Galaxy and satellite samples, for instance, can be constructed to mimic observational selection criteria. Examples of studies using such surveys are constraints on the DM halos of galaxies from satellite kinematics \\citep{mckay_etal02,prada_etal03,brainerd05b} and the anisotropy of the distribution of satellite galaxies \\citep{sales_lambas04,brainerd05}.  In this work, we use the SDSS spectroscopic survey to estimate the projected radial distribution of satellites around isolated primaries. We use areas of the survey which are at least 90\\% complete and check for the effects of incompleteness by comparing the surface density profile and conclude that our results are not affected by the (small) incompleteness of the sample.  We construct samples of primary and satellite galaxies with isolation criteria similar to those used by \\citet{prada_etal03}. We use high-resolution cosmological simulations of the concordance $\\Lambda$CDM cosmology to develop and carefully test new methods of correcting for interloper contamination. Our main results and conclusions can be summarized as follows.  \\begin{itemize}  \\item[1.] Using mock galaxy catalogs derived from high-resolution cosmological simulations, we show that interlopers can significantly bias the shape of the projected surface density profile of faint satellites around bright galaxies, making it shallower (biasing the power-law slope $\\alpha$ of the radial profile, $\\Sigma(R)\\propto R^{\\alpha}$, by $\\Delta\\alpha_{\\rm bias}\\gtrsim 0.5$). We also show that the most straightforward methods do not correct interloper contamination properly.  For example, the random points method, which assumes uniform distribution of interlopers in space, underestimates the fraction of interlopers in the satellite sample by oversampling voids compared to clustered areas where most galaxies in the sample reside.  \\item[2.] We develop two new methods to account for the interloper contamination: the clustered random points method and the nearby points method, variants of the random points method, designed to sample environments similar to those of the clustered galaxies in the observed samples.  Tests on the mock samples show that the methods perform consistently well, reducing the interloper bias in the best-fit power-law slopes of the satellite profiles to only $\\Delta\\alpha_{\\rm bias}\\approx 0.1$ for the nearby points method.  \\item[3.]  We apply these methods in our analyses of the volume- and flux-limited SDSS spectroscopic samples. The best fit power-law slope for the volume-limited SDSS satellite sample, after interloper contamination correction, is $-1.58 \\pm 0.11$ in the range of projected separations of $32.9<R<500h^{-1}$~kpc.  Note that we estimated a systematic bias in the derived slope of $\\Delta\\alpha_{\\rm bias}\\approx 0.1$ for nearby points interloper subtraction method, which implies that the true slope of the SDSS satellites may be $\\alpha=\\alpha_{\\rm raw}-\\Delta\\alpha_{\\rm bias} \\approx -1.7$.  We find similar values of the best fit slopes for the flux-limited samples, and for samples of primary galaxies in different absolute magnitude ranges. We thus do not find evidence for the dependence of the shape of satellite radial distribution on the luminosity of their host galaxy.  \\item[4.] Comparison of the observed radial distribution of the SDSS satellites to the distribution of subhalos and dark matter around galactic halos in dissipationless $\\Lambda$CDM simulations shows that the slope of the SDSS satellite radial profile is in between those measured for subhalos and for dark matter (closer to dark matter).  Subhalos thus appear to have more extended and shallower radial distributions than the observed satellites. The dark matter distribution is somewhat steeper than the observed satellite profiles, but the difference is not statistically significant. \\end{itemize}  Recently, \\citet{vandenbosch_etal05} and \\citet{sales_lambas05} studied the projected radial distribution of satellite galaxies around isolated galaxies using the Two Degree Field Galaxy Redshift Survey (2dFGRS). \\citet{vandenbosch_etal05} used mock galaxy redshift samples derived from large cosmological simulations to develop an iterative method of interloper rejection for the 2dFGRS and found that the data is generally consistent with the dark matter profile at large projected radii, but concluded that incompleteness of close pairs in the survey prevent strong constraints. \\citet{sales_lambas05} used isolation criteria for the primary galaxies, enforcing that any neighbor within a region of 700 $h^{-1}$ kpc and $\\Delta V$ = 1000 km s$^{-1}$ must be at least 1 magnitude fainter.  Satellites were assumed to be any object within projected separation of $500$~kpc and $\\Delta V = 500\\ {\\rm km\\, s^{-1}}$ that is at least 2 magnitudes fainter than the host galaxy. \\citet{sales_lambas05} estimated the effect of close-pair incompleteness by deriving a control sample of projected satellites with velocity difference of $2000 < \\Delta V < 10,000\\ \\,{\\rm km\\, s^{-1}}$. At $R \\gtrsim$ 20 $h^{-1}$ kpc, where the number density profile of the control sample is flat, the satellite samples were considered complete.  For the total sample, \\citet{sales_lambas05} quoted the best-fit power-law of slope $\\alpha = -0.96 \\pm 0.03$. This value, however, was derived without applying correction for interloper contamination.  Note that we derive significantly steeper profiles of satellites compared to those of \\citet{sales_lambas05}, which we attribute to the rigorous correction for interlopers and to possible fiber-collision bias in the 2dFGRS.  The results presented in this study provide interesting hints of the possible differences between observed satellite distributions and the expected distribution of subhalos in their parent halos. However, the statistical errors are still rather large. Significant improvements in the statistics are needed to address this question further. Larger statistics would also allow us to go beyond the average profiles and study the distribution of satellites as a function of satellite (e.g., luminosity and color) and host galaxy properties."
    },
    "0512/astro-ph0512006_arXiv.txt": {
        "abstract": "Recently, there have been suggestions that the apparent accelerated expansion of the universe is not caused by repulsive gravitation due to dark energy, but is rather a result of inhomogeneities in the distribution of matter. In this work, we investigate the behaviour of a dust dominated inhomogeneous Lema\\^itre-Tolman-Bondi universe model, and confront it with various astrophysical observations. We find that such a model can easily explain the observed luminosity distance-redshift relation of supernovae without the need for dark energy, when the inhomogeneity is in the form of an underdense bubble centered near the observer. With the additional assumption that the universe outside the bubble is approximately described by a homogeneous Einstein-de Sitter model, we find that the position of the first CMB peak can be made to match the WMAP observations. Whether or not it is possible to reproduce the entire CMB angular power spectrum in an inhomogeneous model without dark energy, is still an open question. ",
        "introduction": "The first indications that the universe is presently in a state of accelerated expansion were given by J.-E. Solheim as far back as in 1966 \\cite{solheim66}. Using the observed luminosity of several cluster galaxies he found that the model giving the best fit to the data was one with a non-vanishing cosmological constant and negative deceleration parameter. It is, however, only after the more recent observations of the luminosity of supernovae of type Ia (SNIa) that this claim has grown in popularity. The first SNIa observations supporting this claim were those of Riess et al. in 1998 \\cite{riess98} and Perlmutter et al. in 1999 \\cite{perlmutter99}. Since then, more recent observations of supernovae seem to strengthen this claim even further \\cite{tonry03, knop03, Riess:2004nr}. Other independent observations that appear to favour the picture of a universe in a phase of accelerating expansion, are the measurements of the anisotropies in cosmic microwave background (CMB) temperature \\cite{Spergel:2003cb} and the galaxy surveys \\cite{tegmark03}. With these observations in mind, the current period of accelerated expansion seems to be well-established. The physical mechanism that drives this accelerated expansion is, however, still an open question. It is usually ascribed to an exotic energy component dubbed \\emph{dark energy}, whose nature remains a mystery.   Recently, there have been several papers discussing the possibility that the apparent accelerated expansion of the universe is not caused by this mysterious dark energy, but rather by inhomogeneities in the distribution of matter. Most of these papers look at the backreaction effects arising from perturbing homogeneous models, and try to explain the accelerated expansion as corrections to the zeroth order evolution from the higher-order, inhomogeneous terms (see e.g. \\cite{barausse:2005,Kolb:2005me,Wiltshire:2005fw, Carter:2005mu, Moffat:2005jc,Buchert:2005kj,Kolb:2005da}). However, several papers criticizing some of this work have appeared \\cite{Flanagan:2005,Geshnizjani:2005,Hirata:2005,Siegel:2005,Ishibashi:2005}.  Another approach is to look at inhomogeneities of a larger scale in the form of underdense bubbles. The basic idea behind this line of explanation is that we live in an underdense region of the universe, and the evolution of this underdensity is what we perceive as an accelerated expansion. An analysis of early supernova data by Zehavi et al. gave the first indications that there might indeed exist such an underdense bubble centered near us \\cite{zehavi98}.  Specific models that give rise to such underdensities have been studied previously in the form of a local homogeneous void \\cite{Tomita:1999rw,Tomita:2001,Tomita:2000jj}. In these works both the underdensity and the region outside it are assumed to be perfectly homogeneous Friedmann-Robertson-Walker (FRW) models with a singular mass shell separating the two regions. The inhomogeneity manifests itself as a discontinuous jump at the location of the mass shell.  In this article, we wish to investigate more realistic models where there is a continuous transition between the inner underdensity and the outer regions. Therefore we consider an isotropic but inhomogeneous dust dominated universe model, where the inhomogeneity is spherically symmetric. The  model can then be described within the Lema\\^i{}tre-Tolman-Bondi (LTB) class of spherically symmetric universe models \\cite{Lemaitre:1933,Tolman:1934,Bondi:1947}. To make contact with the ordinary FRW models, we assume that the universe is homogeneous except for an isotropic inhomogeneity of limited spatial extension, where the transition between these two regions is continuous.  In a homogeneous universe, it is possible to infer the time evolution of the cosmic expansion from observations along the past light cone, since the expansion rate is a function of time only. In the inhomogeneous case, however, the expansion rate varies both with time and space. Therefore, if the expansion rates inferred from observations of supernovae are larger for low redshifts than higher redshifts, this must be attributed to cosmic acceleration in a homogeneous universe, whereas in our case it can simply be the consequence of a spatial variation, with the expansion rate being larger closer to us. As shown in \\cite{celerier:99}, this results in an  expression for the luminosity distance-redshift relation where the inhomogeneity mimics the role of the cosmological constant in homogeneous models.  However, the supernova observations are not the only data that support the claim of an accelerating expansion. As mentioned above, CMB observations also seem to lend support to this claim. Therefore, in order for our model to be considered realistic, it should also be able to explain the observed CMB temperature power spectrum. We will not attempt to reproduce the whole CMB temperature spectrum for our inhomogeneous model in this paper. For simplicity, we will limit ourselves to the location of the first acoustic peak. As we will show in section~\\ref{sec:CMB}, it is possible to obtain a very good match to both the supernova data and the location of the first acoustic peak simultaneously. In fact, the match to the supernova data is better than for the $\\Lambda$CDM model.   The observed isotropy of the CMB radiation implies that we must be located close to the center of the inhomogeneity. According to this picture, we are positioned at a rather special place in the universe. On the other hand, this model has the attractive feature that there is no need for dark energy. Also the model is sufficiently simple so that it can be solved exactly. It is therefore a good toy model for testing the ideas of inhomogeneities as a solution to the mystery of the dark energy.   The structure of this paper is as follows: First, we will present our model in section~\\ref{sec:LTB}, parameterized by two functions $\\alpha(r)$ and $\\beta(r)$ related to the distribution of matter and spatial curvature, respectively. Still in section~\\ref{sec:LTB}, we present the formalism needed in order to obtain the luminosity distance-redshift relation for spherically symmetric, inhomogeneous models. In section~\\ref{sec:CMB} we discuss the physics behind the first peak of the CMB spectrum and define a shift parameter that quantifies the deviation of the location of this peak relative to that of the concordance $\\Lambda$CDM model. In section~\\ref{sec:results} we present the results from the confrontation of our model with the physical tests presented in the preceding sections, and discuss briefly the possibility of using the recently detected baryon oscillations in the matter power spectrum to constrain the model even further. Finally, in section~\\ref{sec:conclusion} we summarize our work. ",
        "conclusions": "\\label{sec:conclusion}  The main goal of this paper has been to present a simple model with the ability to explain the apparent accelerated expansion of the universe without the need to introduce dark energy. Inspired by the recent discussions about the possibility of explaining the apparent acceleration by inhomogeneities in the matter distribution, we have studied a model where the observer is assumed to be situated near the center of an underdense bubble in a flat, matter dominated universe. If this model is realistic, we live in a perturbed Einstein-de Sitter universe within 130 million light years from the center of an underdensity that extends about 5 billion light years outwards. Under the assumption of spherical symmetry this universe is described by the Lema\\^itre-Tolman-Bondi space-time.  The two main observations we sought to explain were the luminosity distance-redshift relation inferred from SNIa observations and the CMB temperature power spectrum. These two sets of observations are made at opposite ends of the redshift spectrum, respectively low redshifts for the supernovae and high redshifts for the CMB. The fact that our model is inhomogeneous allows us therefore to choose the geometry and matter distribution such that the physical conditions are favourable for explaining the SNIa at low redshift while they at the same time are favourable for explaining the CMB at high redshifts. We find that a very good fit to the supernova data is obtained if we allow the transverse Hubble parameter to decrease with the distance from the observer. On the other hand, we get a good fit to the location of the first peak of the CMB power spectrum if we assume the universe to be flat with a value of $0.51$ for the Hubble parameter outside the inhomogeneity. Interpolating between these two limiting behaviours we get a good fit to both the supernova data and the location of the first peak.  Our model yields a better fit to the Riess data set of supernovae than the concordance $\\Lambda$CDM model. However, for the CMB fit we tested only for the location of the first peak. Although the model yields a good fit to this, it does not necessarily mean that it matches the whole CMB spectrum. Indeed, since the physics responsible for the acoustic peaks is determined by the pre-recombination era, we would expect the peaks to look more or less the same as for a flat, homogeneous model with $h=0.51$. This suggests that our model might fail to explain the third peak. Furthermore, the model does not appear to be able to explain the observed length scale of the baryon oscillations in the SDSS matter power spectrum either, although one may question whether the data quoted by the SDSS team can be used directly to test our model.   The most powerful way to rule out inhomogeneous universe models would be to do a full analysis of the evolution of perturbations in these models. In that way, one could confront the model with both the full CMB angular power spectrum and the matter power spectrum. Only after such an analysis is carried out can one say whether our model is ruled out or if it is a viable alternative to dark energy."
    },
    "0512/astro-ph0512230_arXiv.txt": {
        "abstract": " ",
        "introduction": "FU~Orionis objects are a small but remarkable class of stars that display several-magnitude outbursts in visible light, followed by decade-long fading phases.  They also exhibit fluxes in excess of photospheric levels at infrared wavelengths, broad and doubled spectral absorption lines, and wavelenght-dependent rotational velocities. Two main theories compete to explain their nature. \\citet{herbig2003} (and references therein) favor a model in which the flare occurs in an unstable star rotating near breakup, and where the spectral properties are explained by a rapidly rotating G$-$supergiant photosphere overlaid with a rising cooler shell.  In contrast, a more widely accepted view first proposed by \\citet{hartmann1985}, is that FU~Orionis stars are pre-main sequence objects and that all the observables have their origin in Keplerian circumstellar disks and pre-natal infalling envelopes that surround them \\citep[see also][for recent arguments in relation to the disk vs. rapidly rotating star debate]{hartmann2004}.  In this picture, the outbursts have their origin in sudden increases in the accretion rate through the disk, and all young stars may experience (probably recurrent) FU~Orionis phases during their evolution \\citep[see also the review by][]{hartmann1996}.  Although the existence of circumstellar disks around young stars is well established observationally, very little is directly known about the structure and properties of the inner-most disk regions ($\\lesssim 1$~AU). For solar-mass and intermediate mass young stars, disk models have been reasonably successful at reproducing the infrared flux excesses characteristic of these systems, as well as the spectroscopic signatures of the accretion process itself. However, these efforts are fundamentally limited by the fact that the observations are spatially unresolved, which introduces important degeneracies with respect to the spatial distribution of the emitting material.  The degeneracies with respect to the spatial distribution of material inherent in spatially unresolved observations can be lifted by observations using a long baseline interferometer, and important progress has already been made using this technique.  For Herbig~Ae/Be and T~Tauri objects (young stars of intermediate and solar-type mass respectively), all observations to date have revealed surprisingly large spatial scales for the near-infrared (NIR) emission \\citep{rmg1999,tuthill1999,rla2000,rmg2001,rla2002,eisner2004,rla2005,eisner2005,rla2005b,jdm2005}; prompting in part a significant revision of disk models, whereby previously ignored detailed physics of the disk inner edge are now incorporated with some success \\citep{natta2001, dullemond2001, jdm2002, muzerolle2003, vinkovic2003}. For a review of observational breakthroughs in this area recently made possible by interferometric techniques at infrared wavelengths, the reader is also referred to \\citet{rmg2006}.  For FU~Orionis objects, far fewer direct observational constraints exist at those scales. On the other hand, following the initial outburst the disk emission in FU~Orionis objects outshines that of their central stars by many orders of magnitude, making them in principle ideal laboratories for testing theories describing the structure and physical properties of accretion disks. Indeed, the interpretation is freed from several complicating factors that affect the more evolved Herbig~Ae/Be and T~Tauri systems: instead of having multiple source components (star plus disk) and disk heating mechanisms (stellar irradiation plus active accretion), only the disk component and active accretion matters for interpreting observations of FU~Orionis objects.  Given that these same theories are being applied to the disks inferred to exist around all young stars, and that the physical conditions in the inner parts of those disks set the initial conditions for planet formation, it is important to validate the disk paradigm using FU~Orionis objects as ground truth tests.  The very first young stellar object to be resolved by long baseline optical interferometry was in fact FU~Orionis itself, and in their original analysis and follow up work \\citet{malbet1998,malbet2005} concluded that the visibility amplitudes measured agreed well with those predicted by thermal emission in an accretion disk with the canonical temperature profile $T \\sim r^{-3/4}$ with parameters tuned to fit the spectral energy distribution (SED).  In this paper we present new interferometer data on three other well known FU~Orionis objects, and find instead that the NIR emission is over-resolved relative to predictions of models that consist only of a standard disk emitting thermally.  In $\\S$~\\ref{obs_sec} we present the observations, instrumental arrangement and data analysis; in $\\S$~\\ref{sizes_sec} we derive spatial scales for the NIR emission using simple geometrical models of general applicability; in $\\S$~\\ref{models1_sec} we construct accretion disk models with which we attempt to simultaneously fit the broadband SEDs and visibility amplitude data; in $\\S$~\\ref{models2_sec} we consider additional source components that may be required in order to explain the measured visibilities; and in $\\S$~\\ref{conclusions} we discuss our results in the broader context and suggest future lines of investigation. ",
        "conclusions": "\\label{conclusions}  We have presented new observations which spatially resolve the NIR emission of three FU~Orionis objects on mas scales, a first for two of the objects (V1515~Cyg and Z~CMa-SE). The spatial scales for NIR emission are surprisingly large, compared to predictions of simple accretion disk models.  This indicates that the accretion models require modification, or that additional model components are required.  As is well known, the infrared SED slopes sensitively probe the radial temperature profile in the disk, and we have shown that while exponents $q=-0.6$ or flatter would be required to match the measured visibilities, these solutions are not able to maintain good agreement with the SED data.  We note however that in this approach, tests of the validity of small variations to the disk temperature profile are limited by the accuracy of the SED photometry adopted, and therefore follow-up studies should ideally be performed using contemporaneous photometry of uniform quality.  It is interesting to contrast our results with the case of the prototype for the class, FU~Orionis, which has only modest long wavelength fluxes in excess of thermal emission in a flat accretion disk, and for which the NIR interferometer data does agree well with this model according to \\citet{malbet1998,malbet2005}. We note that this conclusion rests on the ability to incline the FU~Ori disk by a relatively large amount ($i \\sim 50-60 \\degr$), which in this case does not imply unreasonably low stellar masses given the measured rotational velocities \\citep{kenyon1988}.  Given the existing indications for dense envelopes around V1057~Cyg and V1515~Cyg, we hypothesize as a natural explanation for the low measured visibilities that scattering by material located on 10s of AU scales adds incoherent flux in the interferometer FOV, diluting the visibilities from the relatively compact inner disk. We note that this mechanism may also explain the somewhat larger size (Gaussian FWHM~=~$1.36 \\pm 0.07$~mas) found for V1057~Cyg by \\citet{wilkin2003}, since the Palomar Testbed Interferometer has a $1 \\arcsec$ FOV, $\\times 20^2$ larger than that of the KI. A full calculation of FOV effects in specific scattering models is beyond the scope of this paper, we simply note that considering the PTI measurement alone, the minimum incoherent flux fraction becomes $f_i = 21 \\pm 5$~\\%.  An important consequence of this interpretation is that it compromises the notion to use FU~Orionis objects as {\\em straightforward} test cases of accretion disk theories, at least for most of the better known candidate targets. Indeed, if the disk is not isolated, multi-baseline and multi-wavelength interferometer observations must be used to better discriminate between competing models.  In particular, the hypothesis that scattering is at the origin of the low visibilities measured is testable using measurements in different NIR bands: at shorter wavelengths were scattering is more efficient (e.g. J or H bands) the characteristic sizes should be larger; while thermal disks predict the opposite trend of larger sizes at longer wavelengths (and with a detailed dependence that probes the flatness of the temperature profile). These multi-wavelength observations are already possible at existing interferometers (e.g. KI and the AMBER instrument on the VLTI, \\citet{petrov2000}) or upcoming instruments (e.g. the MIRC combiner at CHARA, \\citet{monnier2004}).  For objects with evidence for dense envelopes, self-consistent radiative transfer codes that include scattering should also be used in order to include this effect in both the SEDs and visibilities for realistic geometries and dust properties.  The exciting possibility that all three objects are instead resolved due to the presence of stellar companions, sometimes invoked to explain the outburst themselves, is also testable with follow-up interferometer observations with improved spatial frequency coverage. Given the KI visibility calibration accuracy (which translates into a 100:1 maximum detectable binary contrast), the telltale and un-ambiguous sinusoidal signatures would be easily detected even for the most pessimistic values of the binary contrast derived above."
    },
    "0512/astro-ph0512094_arXiv.txt": {
        "abstract": "We report on the serendipitous discovery of $\\sim$ 500 low-mass candidate PMS stars in the vicinity of the stellar association LH 52 in the Large Magellanic Cloud. We present evidence that the red faint sequence of these stars seen in the CMD of LH 52 from HST/WFPC2 observations belongs only to the association and follows almost perfectly isochrone models for PMS stars of masses down to $\\sim$ 0.3 M{\\solar}. We find that this feature has a Galactic counterpart and that the mass spectrum of the candidate PMS stars in LH 52 seems to correspond to a Salpeter IMF with a slope $\\Gamma \\simeq -1.26$ in the mass range 0.8 - 1.4 M{\\solar}. ",
        "introduction": "The search for pre-main sequence (PMS) stars in young stellar systems of the Galaxy and the Large Magellanic Cloud (LMC) has verified their existence from both ground-based and space observations. Hillenbrand et al. (1993)  observed the galactic open cluster NGC 6611 ``in the act of forming intermediate-mass stars'' of 3 - 8 M{\\solar}. They found that high-mass stars of age $\\sim 2$ Myr coexist with a PMS population as young as 0.25 Myr, suggesting that the massive stars did not alter the star formation process. Brandl et al. (1999) using VLT/ISAAC observations revealed a PMS population in NGC 3603, which is the most massive {\\sc Hii} region in the Galaxy. Recent WFPC2 imaging with the {\\em Hubble Space Telescope} (HST) showed that the CMD position of stars fainter than $M_{\\rm V} \\simeq 0$ mag indicates PMS stars with masses down to 0.1 M{\\solar} with an age between 0.3 - 1 Myr (Stolte et al. 2004).  In the LMC, Brandl et al. (1996) detected with near-infrared imaging of the central cluster of 30 Doradus (R 136) 108 ``extremely red sources'', which are most likely PMS stars of masses down to $\\sim$ 3 M{\\solar}. PMS stars in R 136 were originally revealed by Hunter et al. (1995) with HST/WFPC2 observations. A red population with masses down to $\\sim$ 1.35 M{\\solar}, well tracked by PMS isochrones of ages 1 - 10 Myr, was also discovered in R 136 by Sirianni et al. (2000) from HST/WFPC2, although no variable extinction was taken into account (Brandl \\& Andersen 2005). HST/NICMOS observations showed that in the region of 30 Doradus there are stars of lower mass down to 0.4 M{\\solar} (Zinnecker 1998) with no evidence for a lower mass cutoff (Brandner et al. 2001). Romaniello et al. (2005) identified a population of objects with H{\\alp} and/or Balmer continuum excess, which they interpret as low mass ($\\sim$ 1 - 2 M{\\solar}) PMS stars.     As far as non-star-burst young stellar systems are concerned, (e.g. stellar associations), studies in our galaxy have revealed their PMS population, but such stars {\\em have not yet been detected, although they should exist also in stellar associations of the Magellanic Clouds (MCs)}. According to Massey et al. (1995) the regions of galactic OB associations show evidence of PMS stars with masses between 5 - 10 M{\\solar} and ages $<$ 1 Myr. Preibisch \\& Zinnecker (1999)  discovered 95 low-mass PMS stars with an age of $\\sim$ 5 Myr in the Upper Scorpius OB association and they suggest a coevality in the formation of low- and high-mass stars. Recent studies of the Orion OB1 association revealed its PMS population in the mass range 0.2 - 2 M{\\solar} (Sherry et al. 2004; Brice\\~{n}o et al. 2005). The only comparable study in the MCs is the one around SN 1987A by Panagia et al. (2000), who identified several hundreds PMS stars through their H{\\alp} emission, and they found that most of them have masses in the range 1 - 2 M{\\solar} and ages between 1 - 2 and 20 Myr.  Associations in the MCs {\\em are close enough for low-mass stars to be observed} with high-resolution ground-based and space imaging, and they {\\em do not suffer from crowding or high extinction}, like the star-burst of 30 Doradus or compact young clusters. Therefore, {\\em if PMS stars are to be found in extra-galactic stellar systems, stellar associations of the MCs should be considered as the best places to search}. In this letter we report the serendipitous discovery of low-mass candidate PMS stars found with HST/WFPC2 in the Color-Magnitude Diagram (CMD) of the LMC association LH 52 earlier published by Gouliermis, Brandner \\& Henning (2005a, from here on Paper I). ",
        "conclusions": ""
    },
    "0512/astro-ph0512577_arXiv.txt": {
        "abstract": "Gamma-Ray Bursts (GRBs) are widely thought to originate from collimated jets of material moving at relativistic velocities.  Emission from such a jet should be visible even when viewed from outside the angle of collimation.  I summarize recent work on the special relativistic transformation of the burst quantities $\\eiso$ and $\\ep$ as a function of viewing angle, where $\\eiso$ is the isotropic-equivalent energy of the burst and $\\ep$ is the peak of the burst spectrum in the power $\\nu F_{\\nu}$.  The formulae resulting from this work serve as input for a Monte Carlo population synthesis method, with which I investigate the importance of off-jet relativistic kinematics as an explanation for a class of GRBs termed ``X-ray Flashes'' (XRFs).  I do this in the context of several top-hat shaped variable opening-angle jet models.  I find that such models predict a large population of off-jet bursts that are observable and that lie away from the $\\epei$ relation.  The predicted burst populations are not seen in current datasets.  I investigate the effect of the bulk $\\gamma$ value upon the properties of this population of off-jet bursts, as well as the effect of including an $\\ojet$-$\\egamma$ correlation to jointly fit the $\\epei$ and $\\epeg$ relations, where $\\ojet$ is the opening solid angle of the GRB jet.  I find that the XRFs seen by \\hetetwo and \\bsax cannot be easily explained as classical GRBs viewed off-jet.  I also find that an inverse correlation between $\\gamma$ and $\\ojet$ has the effect of greatly reducing the visibility of off-jet events.  Therefore, unless $\\gamma > 300$ for all bursts or unless there is a strong inverse correlation between $\\gamma$ and $\\ojet$, top-hat variable opening-angle jet models produce a significant population of bursts away from the $\\epei$ and $\\epeg$ relations, in contradiction of current observations. ",
        "introduction": "\\label{cha:intro}  The importance of collimated jets in GRBs was highlighted by the extremely large isotropic-equivalent energies ($\\eiso$) of very luminous events like GRB 971214 \\citep{kulkarni1998} and GRB 990123 \\citep{kulkarni1999} and by the observation of breaks in afterglow light-curves \\citep{rhoads1997,sph1999,harrison1999}.  \\cite{frail2001} and \\cite{bloom2003} corrected the isotropic-equivalent energies by the beaming fraction obtained from the jet break time in afterglow light-curves and found that the values of the energy released in $\\gamma$-rays ($\\egamma$) were tightly clustered around $10^{51}$ ergs. Recently \\cite{ghirlanda2004} have shown that a tight correlation exists between $\\egamma$ and the peak of the $\\nu F_{\\nu}$ spectrum in the rest-frame, $\\ep$.  Recent results by \\cite{sakamoto2005a} obtained from \\hetetwo \\citep{ricker2003} observations have shown that XRFs \\citep{heise2000,kippen2002}, X-ray-rich GRBs and GRBs lie along a continuum of properties and that XRFs with known redshift extend the $\\epei$ relation predicted by \\cite{lloyd-ronning2000} and found by \\cite{amati2002} to over five orders of magnitude in $\\eiso$ \\citep{lamb2005}.  Relativistic kinematics implies that even a ``top-hat''-shaped jet will be visible when viewed outside its angle of collimation; i.e., off-jet \\citep{ioka2001}.  \\cite{yamazaki2002,yamazaki2003} used this fact to construct a model where XRFs are simply classical GRBs viewed at an angle $\\thetav > \\thetan$, where $\\thetan$ is the half-opening angle of the jet and $\\thetav$ is the angle between the axis of the jet and the line-of-sight.  The authors showed that such a model could reproduce many of the observed characteristics of XRFs.  \\cite{yamazaki2004a} showed that in such a model, the distribution of both on- and off-jet observed bursts was roughly consistent with the $\\epei$ relation.  In this paper, I use the population synthesis method developed by \\cite{ldg2005} and incorporate the relativistic emission profiles calculated by \\cite{graziani2005}, to predict the global properties of bursts localized by \\hetetwo and \\bsaxnosp.  I consider the possibility that the XRFs observed by \\hetetwo and \\bsax are primarily regular GRBs observed off-jet \\citep{yamazaki2004a} and show that it is difficult to account for the observed properties of XRFs in this model.  However, since the effect of special relativity on off-jet emission must exist, I seek to understand its relative importance in the context of current models of GRB jets.  I revisit the top-hat variable opening-angle (THVOA) jet model put forward in \\cite{ldg2005}, now including the effects of relativistic kinematics on off-jet emission.  I present results for several models which explore various regions of the parameter space in $\\gamma$, $\\egamma$ and $\\thetan$.  For this paper, I only consider the effect of relativistic kinematics on off-jet emission\\footnote{In this paper I use the terms``off-jet emission'' or ``off-jet relativistic kinematics'', rather than ``off-axis beaming'', to emphasize that such emission is a direct consequence of special relativity and that it is primarily important beyond the edge of the jet.} from uniform or ``top-hat'' jets; we will consider the effects on Fisher-shaped \\citep{dlg2005a} and Gaussian-shaped \\citep{zhang2004} jets in a future publication \\citep{dlg2005b}.  I describe my population synthesis method in \\S 2 and present the results for various models in \\S 3.  I discuss the results in \\S 4 and draw some conclusions in \\S 5.  Preliminary results were reported in \\cite{donaghy2005a}. ",
        "conclusions": "\\label{cha:conclude}  Bursts with known redshifts have been found to obey the $\\epei$ relation, and a large population of bursts away from this relation is not readily apparent in current datasets.  In particular, the limited sample of XRFs with known redshift information is consistent with an extension of the $\\epei$ relation to over 5 orders of magnitude in $\\eiso$.  \\cite{liang2004} found that the $\\epei$ relation holds internally within a large sample of bright BATSE bursts without known redshift, perhaps indicating that the relation is a signature of the physics of the emission mechanism.  However, recently some authors have argued that the $\\epei$ relation may be an artifact of some unknown selection effect arising from the process of determining the burst redshift, and that $25$\\% \\citep{nakar2004a} to $88$\\% \\citep{band2005} of the BATSE bursts may be inconsistent with the $\\epei$ relation. However, these results are controversial \\citep{ghirlanda2005a,bosnjak2005,lamb2005} and depend sensitively on the quality of the spectral fit that generates the $\\eop$ parameter.  I therefore regard the question of the percentage of BATSE bursts that are inconsistent with the $\\epei$ relation to be an open one.  For low values of $\\gamma$, top-hat variable opening-angle jet models predict a sizable population of detectable, off-jet bursts that lie away from the $\\epei$ relation and that are not seen in current data sets. It may be that such bursts are in fact present in the BATSE catalog and form a population of bursts that do not obey the $\\epei$ relation.  On the other hand, if such a population is found not to exist, it implies that the bulk $\\gamma$ of the jet is large ($\\sim 300-1000$).  For models that include the $\\epeg$ relation, the off-jet burst population lies closer to the $\\epei$ relation than for other models, but a similarly discrepant population of off-jet bursts is found to lie above the $\\epeg$ relation, leading to similar conclusions.  Regardless of the size of the population of off-jet bursts, it seems unlikely that such a population can make up the bulk of the XRFs observed by \\hetetwonosp. The larger sample of \\hetetwo bursts without redshift information contains XRFs which lie toward smaller $\\eop$ values than is predicted by the off-jet emission model and hence are not easily explained as classical GRBs viewed off-jet.  Models in which XRFs are the product of larger jet opening-angles seem to better match the observed distributions of XRF properties.  This seems to match the evidence arising from X-ray afterglows of XRFs.  \\cite{granot2005} calculated the afterglow light curves predicted by various models of burst emission seen off-jet.  They find that a general feature of off-jet afterglows is an initial rising light-curve that peaks at about the jet break time and then declines rapidly, similar to an on-jet event.  Afterglows with initially rising components have not been observed.  In particular, XRFs with well-observed X-ray afterglows, for example XRF 020427 \\citep{amati2004} and XRF 050215b \\citep{sakamoto2005b}, have afterglow light curves that join smoothly onto the end of the prompt emission and that show no evidence of a jet break for many days after the burst, implying large jet opening-angles.  It is straightforward to arrange for top-hat variable opening-angle jet models to match the empirical $\\epeg$ relation, and such models also provide a natural explanation for XRFs.  Figures \\ref{fig_ggl1} and \\ref{fig_ggl2} illustrate the consequences of adopting the correlation between $\\ojet$ and $\\egt$ that ensures that on-jet events obey the $\\epeg$ relation.  Most importantly, incorporating the $\\epeg$ relation in the THVOA model removes two of the main drawbacks of the THVOA model presented in \\cite{ldg2005}.  The requirement of very small ($\\sim 2^{\\circ}$) jet opening angles to explain the largest $\\eiso$ values was criticized \\citep{stern2003} as being difficult to achieve in a hydrodynamic jet.  The high end of the $\\eiso$ distribution is here explained by jets with moderate opening angles but larger $\\egamma$ values.  The need to re-scale the central value of $\\egt$ downward to $\\sim 10^{49}$ ergs to incorporate the XRFs in a unified model was criticized as being difficult to reconcile with afterglow models.  The $\\epeg$ relation naturally produces a range of $\\egamma$ values that extends down into the XRF regime.  Matching the $\\epeg$ relation also mitigates the problem of a large population of bursts seen away from the $\\epei$ relation in the low $\\gamma$ case.  In these models the off-jet events hew more closely to the $\\epei$ relation, but a similar problem arises in that these off-jet events are seen away from the $\\epeg$ relation instead.  A possible way out of this dilemma might be to impose a relationship whereby narrower jets have larger bulk $\\gamma$ values and broader jets have smaller $\\gamma$ values.  Using a simple model where $\\gamma \\propto 1/\\ojet$ results in a substantial reduction in the percentage of detected bursts seen off-jet.  \\cite{yamazaki2004b} have proposed a multiple subjet model for unifying short and long GRBs, X-ray-rich GRBs and XRFs.  The model employs emission from multiple subjets (seen off-subjet) to explain X-ray-rich GRBs and XRFs.  The authors performed Monte Carlo simulations for a universal multiple subjet model and find that the results are consistent with the $\\epei$ relation, albeit with considerable scatter (see Figure 4 in \\cite{toma2005}).  There are two reasons why the multiple subjet model better satisfies the $\\epei$ relation.  First, the authors choose $\\gamma = 300$, which satisfies the constraint on $\\gamma$ that I find above.  The behavior of bursts viewed outside the envelope of subjets should approximate the top-hat models considered in this work.  If the authors had adopted a lower value of $\\gamma$, the model would have produced a large number of bursts that lie away from the $\\epei$ relation; i.e., they would have encountered the same problems as those of model THVOA1.  Second, since $\\gamma = 300$, the spectrum for each line of sight is dominated by that of the closest subjet, and since there are many subjets, each line of sight lies very close to at least one subjet, mitigating the effects of relativistic kinematics produced by viewing a subjet well off the jet.  Finally, \\cite{toma2005} find that for values of $\\gamma$ lower than $300$, the ratio between GRBs, X-ray-rich GRBs and XRFs becomes highly skewed toward hard GRBs, in contradiction with the HETE-2 results. Thus, all of the results in \\cite{toma2005} support the requirement I find in this paper that large gamma values are needed in order to match the observed data for XRFs, X-ray-rich GRBs, and GRBs.  Off-jet relativistic kinematics will be important in non-uniform jets as well as in top-hat jets.  Models employing Gaussian \\citep{zhang2004} or Fisher-shaped \\citep{dlg2005a} jets rely on the exponential fall off of the emissivity with viewing angle to match the wide spread of observed burst quantities.  By including off-jet emission in these models, the exponential fall off will be dominated at some angle by the power-law fall off due to relativistic kinematics, thereby broadening the emissivity distribution and reducing the range of generated $\\eiso$ values.  \\cite{graziani2005} showed that different underlying burst profiles may have radically different observational distributions.  We hope to use population synthesis Monte Carlo simulations to further explore these models in future work.  In conclusion, off-jet emission from collimated GRB outflows should exist simply as a consequence of relativistic kinematics.  Monte Carlo population synthesis simulations of top-hat shaped variable opening-angle jet models predict a large population of off-jet bursts that are observable and that lie away from the $\\epei$ and $\\epeg$ relations. Such off-jet events are not apparent in current datasets. These discrepancies can be removed if $\\gamma > 300$ for all bursts or if there is a strong inverse correlation between $\\gamma$ and $\\ojet$. The simulations show that XRFs seen by \\hetetwo and \\bsax cannot be easily explained as classical GRBs viewed off-jet, and are more naturally explained as jets with large opening-angles."
    },
    "0512/astro-ph0512431_arXiv.txt": {
        "abstract": "We have studied the unbarred Sb galaxy with a nuclear star-forming ring, NGC 7742, by means of 2D spectroscopy, long-slit spectroscopy, and imaging, and have compared the results with the properties of another galaxy of this type, NGC 7217, which is studied by us earlier. Both galaxies have many peculiar features in common: each has two global exponential stellar disks with different scalelengths, each possesses a circumnuclear inclined gaseous disk with a radius of 300 pc, and each has a global counterrotating subsystem, gaseous one in NGC 7742 and stellar one in NGC 7217. We suggest that past minor merger is the probable cause of all these peculiarities, including appearance of the nuclear star-forming rings without global bars; the rings might be produced as resonance features by tidally induced oval distortions of the global stellar disks. ",
        "introduction": "Nuclear rings, looking prominent features due to their intense star formation, are found mostly in barred galaxies so they are commonly treated as linked to inner Lindblad resonances where all radial gas inflows are slowed down and where gas is accumulated \\citep{butcr93,shlosrings}. However there are several cases of spectacular nuclear rings in unbarred galaxies, such as those in NGC 278, NGC 7217, NGC 7702, NGC 7742. Most of these galaxies are seen face-on, so the conclusion about the bar absence is quite safe in them. Suggestions about the nature of nuclear rings in unbarred galaxies includes: resonance effects produced by weakly triaxial potential \\citep{jungpal,bu95}; resonance effects produced by a past bar which is now dissolved \\citep{ath}; viscous gas accretion produced by rotation velocity shear in the global disk and its accumulation at a stagnation point at the turnover radius of the rotation curve \\citep{wegmc2}; finally, minor merger \\citep{n278}. Among these hypotheses, the last provides more opportunities to explain various combinations of observational facts. Indeed, \\citet{athcoll} have shown that vertical central impact of a small satellite whose mass is about 10\\%\\ of the host mass should produce a nuclear stellar ring which is morphologically indistinguishable from a resonance ring. On the other hand, if the merged satellite orbit was close to the main galaxy disk plane, their gravitational interaction might produce an oval disk distortion which could in its turn create a resonance nuclear ring. However, certain combinations of predictions are provided by each theoretical model, and by collecting more various observational data, both morphological and kinematical, for every galaxy in question, we would be able at last to restrict possible mechanisms of nuclear ring generation in any particular case.  In this paper we will consider NGC 7742 and NGC 7217; both galaxies have prominent nuclear star-formation rings with a radius of some $10\\arcsec$. We will attempt to find any general features which may be connected to the nuclear ring origin. As for the latter galaxy, now we are undertaking our third approach to its study. Earlier we have found a circumnuclear gas polar ring and two exponential stellar disks with different scalelengths in it \\citep{we97c,we2000}. Also, NGC 7217 is known to possess two counterrotating stellar subsystems \\citep{mk94,we2000}. Recently the SAURON team has also found a gas-stars counterrotation in the center of NGC 7742 \\citep{sau2}, so this fact promises interesting speculations. Both galaxies are moderate-luminosity unbarred spirals of Sb-type. ",
        "conclusions": "By using a variety of 2D kinematical data as well as deep images of NGC 7742, we analyse stellar and gaseous kinematics in this unbarred Sb galaxy possessing the nuclear star-forming ring; we compare it to NGC 7217, another unbarred spiral galaxy with rings which has been studied by us earlier. We have found some common features in NGC 7217 and NGC 7742.  \\begin{enumerate} \\item{ Both galaxies demonstrate global structure consisting of two exponential stellar disks with different scalelengths; the outer disks look quite normal whereas the inner disks are compact, with $r_0\\approx 1$ kpc, and have unusual high surface brightness.  We would like to propose the following qualitative scenario to form such a `multi-tiers' stellar disk. A few Gyrs ago there may be a sudden global gas redistribution in the disk, due perhaps to external tidal perturbation or minor merger. Before that event stars should form in the disk with a large, normal scalelength, and after that when all the gas had been dropped closer to the center the star formation should continue in the disk with a smaller scalelength and higher surface density.}  \\item{ Both galaxies, NGC 7742 and NGC 7217, have circumnuclear gaseous disks with the radius of some 300 pc, highly inclined to the global disk planes; the outer gas disks are, on the contrary, close to the main galactic symmetry planes. Both galaxies possess also some counterrotationg subsystems. NGC 7742 has all its gas in counterrotation with respect to all its stars, with exception of some newly born stellar population in the ring, while in NGC 7217 the gas outside $R=300$ pc corotates the bulk of stars, but there are some 30\\%\\ of all stars in the inner disk that counterrotates \\citep{mk94}.  Three-dimensional dynamical simulations of the self-consistent evolution of a stellar-gaseous galactic disk unstable with respect to bar-like perturbations presented by \\citet{sec1} proposed a scenario for the origin of circumnuclear inclined gas rings. If initially the gas of the global disk counterrotates the stars, then drifting to the center in a triaxial potential of a transient bar, this gas must leave the disk plane and accumulate on orbits strongly inclined to this plane -- only these inclined orbits remain stable for the initially counterrotating gas near the inner Lindblad resonances. We may suggest that the gas which is now observed as the circumnuclear strongly inclined disks in NGC 7217 and NGC 7742 has come from the outer parts of the galaxies, and when it was there, it counterrotated the stars.}  \\end{enumerate}  As for the problem of the origin of initially counterrotating gas, it may be solved together with the problem of the nuclear star-forming rings origin. If we suggest past minor merger of a dwarf gas-rich galaxy from a retrograde orbit, this event had to supply some amount of counterrotating gas and at the same time it might cause an oval distortion of the stellar disk of the host galaxy that in its turn had to produce rapid radial gas re-distribution and the nuclear star-forming ring appearance -- all the peculiar features observed in NGC 7217 and NGC 7742. NGC 7742 demonstrates strong vertical gas oscillations in its counterrotating gaseous disk implying rather recent gas accretion, NGC 7217 might possess the counterrotating gas in the past, but now it is fully reprocessed into counterrotating stars. \\citet{n278} have detected strongly peculiar kinematics of the neutral and ionized hydrogen beyond the optical stellar disk in the unbarred galaxy with the rings, NGC 278, though the galaxy is morphologically regular and quite isolated; they conclude that the galaxy has recently experienced a minor merger. In absence of the detailed neutral-hydrogen observations well outside the optical borders of the galaxy, one would treat NGC 278 as a twin for NGC 7217 and NGC 7742. To our opinion, NGC 278 may represent an early stage of the evolution having followed a minor merger, with respect to two galaxies considered in our work, and its nearest future is perhaps NGC 7742. {\\it The presence of numerous minor merger signatures in the three unbarred galaxies with nuclear star-forming rings makes the hypothesis of tidally induced oval distortion of the global stellar disks the most attactive scenario for the ring origin in unbarred galaxies.}"
    },
    "0512/astro-ph0512161_arXiv.txt": {
        "abstract": "We present spectroscopic observations of red giant branch (RGB) stars in the Andromeda spiral galaxy (M31), acquired with the DEIMOS instrument on the Keck~II 10-m telescope.  The three fields targeted in this study are in the M31 spheroid, outer disk, and giant southern stream.  In this paper, we focus on the kinematics and chemical composition of RGB stars in the stream field located at a projected distance of $R=20$~kpc from M31's center.  A mix of stellar populations is found in this field.  M31 RGB stars are isolated from Milky Way dwarf star contaminants using a variety of spectral and photometric diagnostics.  The radial velocity distribution of RGB stars displays a clear bimodality --- a primary peak centered at $\\bar{v_1}=-513$~km~s$^{-1}$ and a secondary one at $\\bar{v_2}=-417$~km~s$^{-1}$ --- along with an underlying broad component that is presumably representative of the smooth spheroid of M31.  Both peaks are found to be dynamically cold with intrinsic velocity dispersions of $\\sigma(v)\\approx16$~km~s$^{-1}$. The mean metallicity and metallicity dispersion of stars in the two peaks is also found to be similar: $\\rm\\langle[Fe/H]\\rangle \\sim -0.45$ and $\\rm\\sigma([Fe/H])=0.2$.  The observed velocity of the primary peak is consistent with that predicted by dynamical models for the stream, but there is no obvious explanation for the secondary peak.  The nature of the secondary cold population is unclear: it may represent: (1) tidal debris from a satellite merger event that is superimposed on, but unrelated to, the giant southern stream; (2) a wrapped around component of the giant southern stream; (3) a warp or overdensity in M31's disk at $R_{\\rm disk}>50$~kpc (this component is well above the outward extrapolation of the smooth exponential disk brightness profile). ",
        "introduction": "\\label{intro}  Non-baryonic matter is now known to comprise a significant fraction (23\\%) of the total matter/energy content of the Universe (Spergel et~al.\\ 2003).  It is generally believed that this matter is most likely in the form of weakly interacting cold dark matter (CDM).  Cosmological simulations suggest that the large-scale distribution of matter in the Universe is consistent with CDM predictions (i.e., observations of galaxy clustering).  Simulations also suggest that the evolution of this dark matter imprints a signature on much smaller scales.  In fact, it is expected that CDM will cluster gravitationally on subgalactic scales.  Within the framework of hierarchical formation scenarios (Searle \\& Zinn 1978), it is the merging and accretion of these subhalos that is thought to build up massive galaxies such as the Milky Way and our neighbor the Andromeda spiral galaxy (M31).  The cores of the CDM subhalos that fall into massive potentials to form large galaxies are expected to remain intact, even after experiencing strong tides through several orbital timescales around the host galaxy (Hayashi et~al.\\ 2003).  A fraction of the subhalos will have also likely experienced star formation either prior to, or during assimilation with the host.  A parent galaxy the size of the Milky Way or M31 may be expected to have cannibalised, or at the least be host to, 100--500 dwarf galaxy mass systems over its lifetime (e.g., Klypin et~al.\\ 1999; Moore et~al.\\ 1999; Bode, Ostriker, \\& Turok\\ 2001).  Numerical simulations suggest that the accretion of dwarf galaxies by massive hosts should leave fossil signatures in the form of tidal streams around massive galaxies (Johnston 1998; Bullock, Kravtsov, \\& Weinberg 2001).  The Sagittarius dwarf galaxy (Ibata, Gilmore, \\& Irwin 1994; Majewski et al. 2003) is an excellent example of such a merger in our own Galaxy.  Other merger events are also seen in the Milky Way, such as the Magellanic stream (Mathewson, Cleary, \\& Murray 1974) and the Monoceros stream (Yanny et~al.\\ 2003).  The recent discovery of the giant stellar stream to the south of M31 (Ibata et~al.\\ 2001) confirms that tidal streams are in fact a remnant of the accretion process, and that this process is still occurring today in galaxies.  Studying the dynamics, composition, and structure of these streams will help produce a clear picture of the recent accretion history of these host galaxies and further help constrain what fraction of halos may originate in these accretion events.  In some ways such studies are more easily carried out in M31 than the Milky Way since our location within the latter cause local streams to have a larger angular extent and lower stellar density on the sky.  The giant stellar stream of M31 is of particular interest as it is easily seen in optical star-count maps spanning over a very large radial extent ($>$125 kpc), and therefore can set important constraints on the potential of M31 (e.g., Geehan et~al.\\ 2005) ",
        "conclusions": "\\label{conclusion}  We present Keck/DEIMOS observations of red giant stars in three~fields in the Andromeda spiral galaxy: a field on the giant southern stream, a minor-axis spheroid field, and a major-axis disk field.  In this paper, we discuss the kinematics and chemical composition of stars in the stream field at a projected distance of $R\\sim21$~kpc from M31's center.  The disk and spheroid fields will be addressed in future papers.  We isolate RGB stars in M31 by removing foreground Galactic dwarf star contaminants through probability distribution functions calculated from a training set of known RGB and dwarf stars.  These functions combine four~individual parameters --- (1)~radial velocity, (2)~Na\\,{\\sc i} equivalent width (surface gravity sensitive), (3)~position in the CMD, and (4)~comparison between photometric and spectroscopic [Fe/H] estimates --- to compute a likelihood that any given object is an M31 RGB star or a Milky Way dwarf star.  The M31 RGB sample shows the cleanest detection to date of the inner part of the giant southern stream: a kinematically-cold population ($\\sigma(v_1)=16$~km~s$^{-1}$) of 50~stars moving with a mean heliocentric velocity of $\\bar{v_1}=-513$~km~s$^{-1}$.  These stars are found to be metal-rich, with a mean metallicity $\\rm\\langle[Fe/H]\\rangle=-0.47$, and a small metallicity dispersion.  We also find clear evidence for a second population of stars that is also moving with a relatively large negative radial velocity with respect to M31's systemic velocity: $\\bar{v_2}=-417$~km~s$^{-1}$ that is comparably cold: $\\sigma(v_2)=16$~km~s$^{-1}$).  This second cold population has never been reported before.  A close look at previous spectroscopic observations reveals evidence that this second group of stars is not associated with the smooth disk of M31, but rather may represent a new satellite debris trail, a wrapped around component of the giant southern stream, or a tidally-disrupted population of disk stars.  Future observations of the inner regions of the stream will help constrain models for its orbit and put better limits on which, if any, of M31's satellites may be the progenitor."
    },
    "0512/astro-ph0512211_arXiv.txt": {
        "abstract": "We consider a model in which \\SgrA, the $3.7\\times10^6\\,\\Ms$ supermassive black hole candidate at the Galactic Center, is a compact object with a thermally emitting surface.  For very compact surfaces within the photon orbit, the thermal assumption is likely to be a good approximation because of the large number of rays which are strongly gravitationally lensed back onto the surface.  Given the very low quiescent luminosity of {\\SgrA} in the near infrared, the existence of a hard surface, even in the limit in which the radius approaches the horizon, places a severe constraint upon the steady mass accretion rate onto the source: $\\dot{M}\\lesssim10^{-12}\\,\\Ms\\yr^{-1}$.  This limit is well below the minimum accretion rate needed to power the observed submillimeter luminosity of \\SgrA: $\\dot{M} > 10^{-10}\\,\\Ms\\yr^{-1}$.  We thus argue that \\SgrA~does not have a surface, i.e., it must have an event horizon.  The argument could be made more restrictive by an order of magnitude with $\\muas$ resolution imaging, \\eg, with submillimeter very-long baseline interferometry. ",
        "introduction": "Infrared observations of individual stars in the Galactic Center imply the existence of a dark object of mass $M \\approx 3.7\\times10^6\\,\\Ms$, constrained to lie within $45\\,{\\rm AU}$ (or $10^3 GM/c^2$) of the radio source \\SgrA~\\citep{Scho_et_al:03,Ghez_et_al:05,Eise_et_al:05}. Observations at $3.5\\,\\mm$ and $7\\,\\mm$ further constrain the extent of the radio emission from \\SgrA~to less than $1-2\\,{\\rm AU}$ (or $10-20GM/c^2$) \\citep{Shen_etal:05,Bowe_etal:04}.  The favored interpretation of these observations is that \\SgrA~is a supermassive black hole.  Indeed, the current constraints rule out many alternative explanations, including clusters of stellar mass compact objects \\citep{Maoz:98} and fermion balls \\citep{Scho_etal:02}.  Nonetheless, it remains to be conclusively demonstrated that the dark mass at the Galactic Center is a true black hole with an event horizon.  If {\\SgrA} is not a black hole, then it must have a surface at some radius $R$.  Although general relativistic considerations coupled with reasonable assumptions on the equation of state of matter require $R \\geq 9GM/4c^2$ \\citep[see, \\eg,][]{Shap-Teuk:86}, alternatives to general relativity exist which allow smaller radii, despite the fact that the exterior spacetime may be arbitrarily close to that predicted by general relativity (\\eg, scalar-tensor theories, \\citealt{Fuji-Maed:03}, gravastars, \\citealt{Mazu-Mott:01}, boson stars \\citealt{Torr-Capo-Lamb:00}). Thus, in principle, $R$ could have any value greater than $2GM/c^2$, the horizon radius (we restrict our analysis to non-spinning objects).  In this {\\it Letter}, we show that current observations do not favor {\\SgrA} having such a surface.  We assume that any putative surface of {\\SgrA} is in steady state in the presence of accreting gas, and that it emits the accreted energy thermally.  The latter assumption is reasonable since, even for models in which the radius of the surface approaches $2GM/c^2$, the thermalization timescale is short in comparison to the lifetime of {\\SgrA} (or of an observer)\\footnote{The thermalization timescale is expected to be on the order of the cooling time of the infalling matter, which is not very different from the free-fall time-scale $\\sim 100$ s at the surface.  Since the timescale diverges only logarithmically as measured at infinity even for extremely compact configurations (\\eg, the gravastar), the thermalization time is increased by only a few orders of magnitude.}.  Indeed, for surfaces contained well within the photon orbit, strong gravitational lensing significantly decreases the number of photon trajectories that escape to infinity.  Most outgoing rays return to be absorbed by other parts of the surface, so that the object will resemble the classical example of a blackbody: a thermal cavity with a pinhole.  If \\SgrA~accretes at the Bondi rate from the hot gas surrounding it, the accretion rate is expected to be $\\dot M \\sim 10^{-6} ~M_\\odot\\,{\\rm yr^{-1}}$ \\citep{Baga_etal:03}.  A more likely scenario is that the source accretes via a radiatively inefficient accretion flow \\citep[RIAF;][and references therein]{Nara-Yi-Maha:95,Yuan-Quat-Nara:03}, with a mass accretion rate in the range $\\dot M \\sim 10^{-8.5}-10^{-6} ~M_\\odot\\,{\\rm yr^{-1}}$.  In a RIAF model, essentially all the potential energy released by the accreting gas would be radiated from the surface of {\\SgrA} (assuming it has a surface), with a predicted luminosity at infinity of \\begin{equation} L_{\\rm surf} \\approx \\eta \\dot M c^2 ,\\label{Lsurf} \\end{equation} where the efficiency factor $\\eta$ is the fraction of the rest mass energy of the infalling gas that is released as radiation.  In the Newtonian case, $\\eta$ is simply $GM/c^2R$.  In general relativity, $\\eta$ is given in terms of the gravitational redshift $z$ at the surface: \\begin{equation} \\eta = z/(1+z), \\quad 1+z = (1-2GM/c^2R)^{-1/2}. \\label{eta} \\end{equation} Although it is highly unlikely that \\SgrA~has a radiatively {\\it efficient} accretion disk, even such a model requires a fairly large $\\dot M$.  For example, the observed bolometric luminosity of {\\SgrA} of $10^{36} ~{\\rm erg\\,s^{-1}}$ implies a minimum accretion rate of $\\dot M \\sim 2\\times 10^{-10} ~M_\\odot\\,{\\rm yr^{-1}}$ for a radiative efficiency of 10\\%.  To within a factor of a few (depending on the nature of the boundary layer at the inner edge of the disk), the luminosity from the surface of {\\SgrA} is again predicted to be $\\sim \\eta\\dot Mc^2$ with $\\eta$ not very different from (\\ref{eta}).  All the above estimates are for gas accretion.  When stellar capture events are considered, the average accretion rate can be as high as $10^{-5}~M_\\odot\\,\\yr^{-1}$ to $10^{-3}~M_\\odot\\,\\yr^{-1}$, though this would be expected to be in the form of transient accretion events \\citep{Mago-Trem:99}.  Note that the $\\dot M$ estimates given here are from the point of view of a distant observer, i.e., they represent the rate of accretion of rest mass per unit time as measured at infinity.  In \\S~2, we derive upper limits on $\\dot M$ from the observed near-infrared (NIR) fluxes of {\\SgrA} and compare these with the various estimates of $\\dot M$ given above.  On this basis we argue that \\SgrA~is unlikely to have a surface and therefore that it must be a black hole.  In \\S~3, we present theoretical images of the RIAF model discussed in \\citet{Yuan-Quat-Nara:03} and \\citet{Brod-Loeb:05b}, and show that imaging experiments alone cannot distinguish between a black hole and a compact object with a surface. However, by combining imaging with the argument presented in \\S~2, we show that one could make the case for an event horizon stronger.  We conclude in \\S~4 with a discussion.  In what follows, unless otherwise noted, we use geometrized units ($G=c=1$). ",
        "conclusions": "({\\em i}) the surface is in steady state with respect to the accreting material, ({\\em ii}) the surface radiates thermally, and ({\\em iii}) general relativity is an appropriate description of gravity external to the surface.  As mentioned briefly in \\S~1, the assumption of steady state is likely to be a good one, even for surfaces very near the horizon (including those which are separated from the horizon by a Planck length, the minimum scale for which a horizon will not develop).  However, it should be noted that for a black hole the unradiated binding energy of the accreting matter contributes to an increase of the black hole's mass.  Thus a black hole is an explicit example of an accreting compact object which is not in steady state.  Of more concern is the assumption that the surface emits thermally. For models in which large-scale correlations play a significant role (\\eg, the gravastar) it is unclear what happens to accreting material. For instance, it is conceivable that one may obtain coherent emission with wavelengths comparable to the correlation length of the surface, which in principle could introduce large deviations from the Planck spectrum.  While we cannot rule out such a model, we would like to emphasize that, in general, for the gravastar (or similar) model to remain a viable alternative to the black hole model of {\\SgrA} necessarily requires an extremely exotic emission mechanism that deviates enormously from blackbody emission.  Finally, some assumption regarding the description of gravity external to the surface is necessary to compute the flux due to a compact surface near a strongly gravitating object.  In the absence of a well tested alternative, general relativity is the natural choice and this is what we have selected for our calculations.  While we have explicitly considered a non-rotating black hole, we expect rotation (the most obvious extension that one would wish to consider) to primarily broaden the thermal spectrum, changing our results by no more than factors of order unity.  Multi-wavelength high-resolution imaging of flares in the Galactic Center has been proposed as a method by which the nature of the spacetime surrounding the Galactic Center's black hole may be probed \\citep{Brod-Loeb:05,Brod-Loeb:05c}.  To summarize, in the absence of unknown exotic phenomena, the current NIR flux measurements already conclusively imply the existence of an event horizon in the black hole candidate {\\SgrA} at the Galactic Center."
    },
    "0512/astro-ph0512027_arXiv.txt": {
        "abstract": "We use a high-resolution grid-based hydrodynamics method to simulate the multi-phase interstellar medium in a Milky Way-sized quiescent disk galaxy.  The models are global and three-dimensional, and include a treatment of star formation and feedback.  We examine the formation of gravitational instabilities  and show that a form of the Toomre instability criterion can successfully predict where star formation will occur.  Two common prescriptions for star formation are investigated.  The first is based on cosmological simulations and has a relatively low threshold for star formation, but also enforces a comparatively low efficiency.  The second only permits star formation above a number density of $10^3$ cm$^{-3}$ but adopts a high efficiency.  We show that both methods can reproduce the observed slope of the relationship between star formation and gas surface density (although at too high a rate for our adopted parameters).  A run which includes feedback from type II supernovae is successful at driving gas out of the plane, most of which falls back onto the disk.  This feedback also substantially reduces the star formation rate.  Finally, we examine the density and pressure distribution of the ISM, and show that there is a rough pressure equilibrium in the disk, but with a wide range of pressures at a given location (and even wider for the case including feedback). ",
        "introduction": "Star formation in galactic disk systems is the product of a large number of physical processes and so is potentially quite complicated.  In outline, gravity tries to form the dense molecular clouds out of which stars form, while rotational shear, thermal pressure, turbulence, magnetic  fields, and cosmic ray pressure resist the collapse.  Since the gas in disks does not turn into stars on a free-fall time, one or more of these resistive mechanisms must be effective.  The classic condition for disk instability, the Toomre $Q$ parameter \\citep{toomre64} encodes the impact of shear, which suppresses large-scale perturbations, and the effective sound speed, or pressure (which suppresses small-scale fluctuation, and may be any of the physical processes described earlier).  When $Q$ is above some critical value, the large and small-scale suppression ranges overlap and no (linear) perturbations are gravitationally unstable.  While this picture is theoretically pleasing and has substantial observational support \\citep{kennicutt89,boissier03,heyer04}, it is still not perfectly clear which of the physical processes listed above actually supply the local effective pressure.  Probably the leading candidate is turbulence because it is seen in all disks and has many of the correct properties, but it should be noted that the magnetic fields and cosmic-ray pressure have similar energy densities \\citep{boulares90}.  Turbulence in galactic disks can be generated from a number of sources, including gravity \\citep{wada02}, stellar winds, supernovae \\citep[e.g.,][]{MacLow2004}, the magneto-rotational instability \\citep{sellwood99} and even radiative heating \\citep{kritsuk02}.  Whatever the source, the combined effect of the effective pressure, gravity and shear must reproduce both the observed threshold density for star formation and also the observed relation between gas surface density and star formation \\citep{schmidt59, kennicutt89, kennicutt98, martin01}.  A related question is the structure and distribution of gas densities, temperatures and pressures within the interstellar medium (ISM).  Observational and theoretical work have suggested a picture of a multi-phase medium with substantial turbulent motions \\citep{mo77, mccray79, larson81, stanimirovic99, 2000MNRAS.315..479D, elmegreen01}.  However, the distribution of gas (both in terms of volume and density) in the various phases is not well understood, although substantial observational progress has been made \\citep{jenkins78, shelton94, ferguson96, chu99, shelton01}.  It is clearly important to try to probe this topic theoretically in order to tease apart the connection between star formation, feedback and the ISM, not simply for a better understanding of our galaxy and local galaxies, but also to model star formation and galaxy formation at high redshift.  Numerical work often focuses on one of two aspects; either a detailed analysis of the ISM and a smaller simulation area, or a study of the global disk instabilities and star formation at the cost of a simplified ISM model. Models that have tackled both topics have either been in  two-dimensions \\citep{wn2001} or restricted to a box size a few hundred parsecs across \\citep{2001ApJ...559L..41W}.  Work performed in two dimensions by Rosen \\& Bregman (1995) allowed the ISM to evolve self consistently, but treated the stars as a collisional rather than collisionless fluid.  They modeled the galaxy side-on so that the simulation region included a dimension out of the disk. They found that the gas formed a three-phase medium with cold and warm filaments surrounding bubbles of hot gas. The bubbles of hot gas extend to up to a kiloparsec across, with filamentary structure similar to that observed in our own Galaxy.  Self consistent treatment of the ISM and star formation has been performed in two dimensions (with both dimensions in the plane) by \\citet{wn2001} and in three dimensions over a small box size \\citep{2001ApJ...559L..41W}. They see three phases, but also gas that exists in unstable regions between these phases.  They therefore argue that a simple two or three phase model of the ISM is not sufficient to represent it properly, and that turbulence results in the smearing out of the phases so that gas exists outside pressure equilibrium. Wada also finds that the hot gas is a product of the supernovae explosions and both the hot and warm gas exist off the surface of the disk, leaving the cold gas on the disk plane.  In three dimensions but considering a small section of disk, \\citet{2000MNRAS.315..479D} also utilized a separate stellar disk to explore the collective effects of type I and II supernovae on the structure of the ISM. His simulations were performed in three dimensions for a section of a galactic arm, located $8.4$\\,kpc from the galactic centre using a fixed gravitational field. The supernovae locations were determined randomly but with constraints imposed to give a realistic distribution. His results show cold gas is present in a thin, irregular layer on the galactic plane, intercrossed with tunnels of hot gas from supernovae explosions. Around the cold gas is a thick disk of neutral warm gas up to $500$\\,pc followed by ionized warm gas and then hot gas at heights above $1.5$\\,kpc. Places where several supernovae merge form reservoirs of hot gas that have enough energy to break free of the gravitational pull of the stellar disk and expand upwards in large bubbles.  Korpi \\citep{1999ApJ...514L..99K} modeled a section of the ISM self consistently, including  the effects of type I and II supernovae and that of magnetic fields, but left out star formation and self gravity. Their ISM formed a two phase structure of warm and hot gas with a bimodal temperature-density distribution. The warm gas was found in scale heights less than $500$\\,pc and the hot gas above that. They also found a cold component, due to compression by the warm gas, which was found at heights of less than $100$\\,pc. The supernovae in their simulations clustered to produce large non-spherical shells.  Several numerical studies have simplified their treatment of the ISM to study the global properties of the disk and star formation.  Robertson et al. \\citep{2004ApJ...606...32R} and Semelin \\& Combes (2002) both assume a two-phase ISM consisting of cold clouds embedded in a warm gas in pressure equilibrium. The ratio of gas in these two phases is controlled statistically by allowing gas to switch phases during supernovae explosions, conduction and cooling processes.  In their paper, \\citet{2004ApJ...606...32R} compared simulations which used first no star formation, then star formation but no feedback and finally feedback with a two phase ISM. In the non-feedback cases, the gas cooled extremely efficiently, resulting in a near isothermal ISM. They found that in these cases, disk fragmentation was catastrophic and the stars (when present) ended up in two big clumps. The addition of star formation stabilised the disk for slightly longer than in the no star formation case. In both cases, the fragmentation was due to the Toomre instability \\citep{toomre64}. The addition of feedback and a multiphase ISM resulted in increased pressure support and a smoother distribution of gas and stars.  \\citet{Li2005a,Li2005b} use an isothermal gas for the ISM and examine star formation for two different temperatures. They find in both cases the gradient of the gas surface density plotted against the surface density of the star formation rate is around 1.5, in good agreement with the Schmidt law and the observations of \\citet{kennicutt89}. They also observe a threshold for star formation, where no stars are formed past two radial scale lengths. This is also the point where the Toomre Q parameter drops below one and the disk becomes stable to Toomre instabilities.  \\citet{2003ApJ...590L...1K} performed hydrodynamic simulations in a cosmological context. In this work, the gas is converted into stars on a characteristic gas consumption time scale, rather than on the dynamical time. Kravtsov finds that this can still reproduce the Schmidt law with a gradient of 1.4. The addition of feedback in these simulations results in considerably more hot gas at low densities although the PDF of the gas density remains unchanged.  This paper is the first step in a longer range plan to understand the fundamental processes of star formation and feedback in a galaxy disk. These are the first three-dimensional simulations of a global disk without the need to simplify the structure of the ISM, using a grid-based code which is better abled to resolve the multi-phase medium (except for the small-disk simulations of \\citet{2001ApJ...559L..41W}, which are also grid-based). We use this model to investigate local star formation throughout the evolution of the disk, from the early fragmentation of the gas into stars through to the global and local properties of the star formation rate and the effect on the evolution of the interstellar medium. We compare the results for two different models of star formation and with and without the inclusion of stellar feedback from type II supernovae. Ultimately, we hope the better understanding of star formation gained from looking at our isolated disk will act as a guide to cosmological simulations of galaxy formation. For these simulations we use a high-resolution adaptive mesh refinement (AMR) code, which includes a more sophisticated treatment of star formation and feedback as well as a full treatment of self-gravity of the gas rather than the fixed potential that is often used. We concentrate on hydrodynamical effects, ignoring (for the moment) magnetic fields and cosmic ray pressure.   In section~\\ref{sec:comp_method}, we describe our computational approach, in section~\\ref{sec:structure} we discuss the structural properties of the disk simulations, including the formation of instabilities and the vertical distribution.  In section~\\ref{sec:sf_properties}, we discuss how  star formation is related to the surface density and compare this to observations, while section~\\ref{sec:ism_properties} contains an analysis of the multi-phase structure of the resulting ISM.  \\newpage ",
        "conclusions": "We have performed high-resolution adaptive mesh refinement simulations of an isolated galactic disk evolved for more than 1 Gyr.  We include many of the physical processes which must be important for the long-term evolution of the gas in spiral galaxies including cooling, shocks, self-gravity, star  formation and supernova feedback in a global three-dimensional model.   Our adaptive-mesh methodology allows us to resolve scales from 100 kpc down to 25 pc, the size of typical giant molecular clouds.  The physical model for the galactic disk is clearly oversimplified in a number of respects: it does not include magnetic fields, cosmic rays, chemistry and the cooling/heating model is incomplete.  Still, this represents a substantial improvement over previous work in a number of ways (see the introduction for a discussion of previous simulation work) and represents some of the most realistic global disk simulations ever performed.  We performed a number of simulations while varying the input physics.   This included runs with cooling down to two minimum temperatures ($10^4$ K and 300 K), but no star formation.  A series of runs were performed with cooling and two different prescriptions for star formation, the first a cosmological-simulation inspired star formation algorithm  which allowed star formation at relatively low densities but with a low efficiency, and the second a more physically-minded algorithm which adopted a high density threshold (comparable to that found in giant molecular clouds) before stars could form.  These two forms we have denoted C-type and D-type, respectively.  We also performed some runs with spatial resolution two times better and mass resolution eight times better in order to investigate numerical convergence.  Finally, feedback from supernovae from massive stars was introduced.  Our results are summarized below: \\begin{itemize}  \\item Gravitational instabilities grow as long as the Toomre $Q$ parameter is less than a critical value (0.6).  Outside of this region, no stars form (although stars can be scattered into this region).  This appears to be a well-resolved result, and does not depend on the star formation algorithm.  \\item If no star formation occurs, the clumps merge and form more massive, denser clouds.  If star formation is permitted, stars form preferentially in the densest part of the clumps (particular for the D-type star formation algorithm).  Without some form of feedback, the clouds  are long-lived and convert a high fraction of their mass into stars.  \\item Both star formation algorithms reproduce the slope of the observed relation between star formation and gas surface density.  This appears to be because clump formation is controlled by the dynamical time.  The C-type (cosmological) method can be tuned (with a sufficiently low efficiency parameter  $\\epsilon < 0.005$) to reproduce the observed normalization of the relation as well.  The D-type method (with a high density threshold) produces too many stars and will require some additional form of feedback to match the normalization.  Energy input from type II supernovae does indeed decrease the star formation rate (although more feedback, such as photoionization, seems to be required to match observations).  \\item A multiphase ISM is naturally reproduced with most of the mass ($>$ 80\\%) in cold, dense clouds and peaks in the volume distribution at temperatures of approximately $10^4$ K and $10^{6.5}$ K.  \\item Feedback from type II supernovae drives material out of the plane of the disk (which then falls back).  However, it does not increase the mean pressure in the plane of disk or generate large amounts of hot gas, or substantially increase the vertical scale-height of the gas.  \\end{itemize}    We thank Adrianne Slyz, Julien Devriendt, Yuexing Li, Mordecai Mac Low and Frazer Pearce for useful discussions.  EJT and GLB acknowledge support from PPARC and GLB the Leverhulme Trust. Some simulations used in this paper were performed at the National Center for Supercomputing Applications."
    },
    "0512/gr-qc0512074_arXiv.txt": {
        "abstract": "We investigate the evolution of  the non-linear long wavelength fluctuations during preheating after inflation. By using the separate universe approach, the temporal evolution of the power spectrum of the scalar fields and the curvature variable is obtained numerically. We found that the amplitude of the large scale fluctuations is suppressed after non-linear evolution during preheating. ",
        "introduction": "Preheating after inflation is a crucial stage in the early universe. Fluctuations of matter fields and gravitational fields are amplified due to  the parametric resonance caused by the coherent oscillation of the inflaton field\\cite{TraschenJ:PRD42:1990,KofmanL:PRL73:1994,KodamaH:PTP96:1996, NambuY:1997,TaruyaA:1998}. One of the important question on preheating is how the long wavelength metric fluctuation is amplified by the parametric resonance. The linear analysis shows that non-adiabatic modes of long wave fluctuations are amplified and grow during preheating until the non-linear effect by the second order fluctuations becomes dominant. When the non-linear effect cannot be neglected,  the backreaction of the fluctuations on the evolution of background quantities  becomes dominant and the amplification of long wave modes stops \\cite{BassetB:PRD62:2000,ZibinJP:PRD63:2001}. To obtain complete understanding of the evolution of inhomogeneities during  preheating, we have to investigate non-linear evolution of fluctuations .  The separate universe approach\\cite{WandsD:2000} is an appropriate method to treat  the non-linear dynamics of long wavelength fluctuations during preheating. This approach neglects fluctuations of which wavelength is smaller than the Hubble horizon scale. For each spatial point, the basic equation of the separate universe reduces to that of the Friedmann equation for a homogeneous and isotropic flat universe. However, all dynamical variables include non-linear inhomogeneities of which wavelength is larger than the Hubble horizon scale. As the separate universe approach includes all order long wavelength non-linear gravitational fluctuations, we can apply this method to investigate the backreaction effect on  evolution of a Friedmann-Robertson-Walker(FRW) universe\\cite{NambuY:PRD71:2005}.  Tanaka and Basset\\cite{TanakaT:2003} investigated preheating using the separate universe approach. They found that initial small fluctuations are amplified and evolve to random spatial distribution. In the separate universe, dynamical variables at each spatial point evolve independently. Hence, in the early stage of evolution, fluctuations keep their initial spatial distribution  and only the amplitude of fluctuations  grows by the parametric resonance. When the amplitude of the massless field grows to be comparable to that of the inflaton field, the non-linear interaction between scalar fields becomes to be dominant and the system enters a chaotic regime. Then,  the field variables at each spatial point behave as independent random variables. At this stage, the power spectrum of fluctuations is  same as that of random white noise.  In this paper, we concentrate on the evolution of the power spectrum of fluctuations during preheating especially on curvature fluctuation that gives an impact on formation of large scale structure. By performing the numerical simulation of preheating based on the separate universe approach, we aim to understand the feature of evolution of long wavelength non-linear fluctuation during preheating.  The plan of paper is as follows. In Sec.~2, we review the separate universe approach. In Sec.~3, we analytically estimate the evolution of the power spectrum of scalar fields and a curvature variable. We present our numerical results in Sec.~4 and Sec.~5 is devoted to summary and conclusion. We use units in which $c=\\hbar=8\\pi G=1$ throughout the paper. ",
        "conclusions": "In this paper, we investigated the evolution  of long wavelength fluctuations during preheating after inflation. By using the separate universe approach, we obtained  the evolution of the power spectrum of long wavelength fluctuations numerically. During the linear stage of the evolution, the fluctuation of the massless field grows exponentially in time by the parametric amplification, but the power spectrum keeps its initial shape. When  the fluctuation of the massless field becomes non-linear, the amplification of fluctuations stops by the back reaction effect of massless field on the inflaton field.  After this time, the system enters the chaotic non-linear stage and  the shape of the spectrum changes. For large scale mode, the amplitude of the fluctuation is suppressed because the statistical nature of the large scale mode becomes that of white noise after the non-linear evolution. The evolution of the curvature variable that has an important role in the cosmological scenario is qualitatively same as the evolution of the scalar field fluctuations. Hence we do not expect the significant effect of the parametric amplification  during preheating on the evolution of the large scale fluctuations of metric variables.  On the other hand, for small scale fluctuations, the power grows for larger $k$ and this leads to the possibility of the formation of the non-linear structures such as primordial black holes\\cite{SuyamaT:PRD71:2005}."
    },
    "0512/astro-ph0512033_arXiv.txt": {
        "abstract": "We show that neutrino-driven pulsar kicks can increase the energy of the supernova shock. The observed large velocities of pulsars are believed to originate in the supernova explosion, either from asymmetries in the ejecta or from an anisotropic emission of neutrinos (or other light particles) from the cooling neutron star. In this paper we assume the velocities are caused by anisotropic neutrino emission and study the effects of these neutrino-driven kicks on the supernova explosion.  We find that if the collapsed star is marginally unable to produce an explosion, the neutrino-driven mechanisms can drive the convection to make a successful explosion.  The resultant explosion is asymmetric, with the strongest ejecta motion roughly in the direction of the neutron star kick.  This is in sharp contrast with the ejecta-driven mechanisms, which predict the motion of the ejecta in the opposite direction.  This difference can be used to distinguish between the two mechanisms based on the observations of the supernova remnants. ",
        "introduction": "Current observations of pulsar proper motions suggest that a large fraction of neutron stars are moving with velocities in excess of 400\\,km\\,s$^{-1}$ \\citep{Cor98,Fry98,Lai01,Arz02}.  The large energy and momentum released during the formation of the neutron star (and the ensuing supernova explosion), coupled with the growing evidence that many core-collapse supernovae exhibit asymmetric explosions, has led to a general consensus in astronomy that neutron stars receive these large ``kicks'' at birth.  The mechanisms driving these kicks can be separated into two classes: ejecta-driven kicks and the kicks driven by emission of neutrinos or other weakly interacting particles.  Ejecta-driven kicks can occur if a sufficient degree of anisotropy develops in the hydrodynamics of the explosion.  Since only 1\\% of the collapse energy accompanies the ejecta, large asymmetries are required to produce large supernova kicks.  A number of ejecta asymmetries have been proposed: asymmetric collapse \\citep{Bur96}, low mode convection \\citep{Her92,Bur03} and the related low-mode convection in an accretion shock instability \\citep{Blo03}.  Asymmetries in the progenitor star can not produce kicks in excess of 200\\,km\\,s$^{-1}$ \\citep{Fry04}, far short of the observed 1000\\,km\\,s$^{-1}$. Asymmetries produced by low mode convection has proven more successful \\citep{Bur03} in 2-dimensional studies.  Such mechanisms require, by momentum conservation, that the kick be along the explosion asymmetry, but moving in the opposite direction of the ejecta.  Asymmetric neutrino emission has been proposed as an alternate kick mechanism.  This mechanism takes advantage of the fact that most of the energy and momentum released in the collapse of a massive star is in the form of neutrinos, and asymmetries of a percent are sufficient to produce the observed kicks.  The proposed mechanisms range from collective effects, for example, turbulence near the neutrinosphere \\citep{Soc05}, to elementary processes involving neutrinos, including neutrino oscillations \\citep{Kus96,Kus97,Bar02,Ful03,Bar04,Kus04}. All these mechanisms require strong magnetic fields. Although the surface magnetic fields of ordinary radio pulsars are estimated to be of the order of $10^{12}-10^{13}$G, the magnetic field inside a neutron star may be much higher, probably as high as $10^{16}$G \\citep{magnetic,dt,magnetars}.  Naively, one might think that even the standard urca reactions responsible for production of neutrinos, $ p+e^-\\rightleftharpoons \\nu_e+n$ and $\\bar\\nu_e+p\\rightleftharpoons n+e^+$, have a sufficient asymmetry to give the neutron star a kick.  Indeed, in the rates of the urca processes depend on the relative orientations of the electron spins and the neutrino momentum.  Hence, there is a 10-20\\% anisotropy in the distribution of neutrinos in every one of these processes~\\citep{chugai,drt}. However, this asymmetry in production does not lead to any asymmetry in the emission of neutrinos, because the anisotropy is washed out by the re-scattering of neutrinos on their way out of the star \\citep{Vil95,Kus98,Arr99}.  If some other particles, with interactions weaker than those of neutrinos were produced anisotropically, their emission would remain anisotropic.  For example, if sterile neutrinos exist and have a small mixing with active neutrinos, they should be produced in the urca processes at the rate suppressed by the square of the mixing angle \\citep{Ful03}.  It is intriguing that the parameters of the sterile neutrinos required for the pulsar kicks \\citep{Kus97,Ful03} are consistent with the mass and mixing that make the sterile neutrino a good dark matter candidate \\citep{Ful03}.   There is a strong evidence that most of the gravitating matter in the universe is not made of ordinary atoms.  This evidence is based on a consensus of observations of galaxy rotation curves, cosmic microwave background radiation, gravitational lensing, and X-ray emission from galaxy clusters. None of the known particles can be the dark matter, and a number of candidates have been proposed.  Perhaps, the simplest extention of the Standard Model that makes it consistent with cosmology is the addition of a sterile neutrino with a 2-15 keV mass.  Unlike the active fermions, which must be added in the whole generations to satisfy the anomaly constraints, or the supersymmetric particles, which require a major modification of the particle content, the sterile neutrino does not entail and additional counterparts because it is gauge singlet. Sterile neutrinos can be produced from neutrino oscillations in the early universe in just the right amount to be the dark matter \\citep{dw,Aba01,Aba01a,Dol02,Map05}. If their mass exceeds 2~keV, they are sufficiently cold to explain the large-scale structure.  The discovery of neutrino oscillations points to the existence of some gauge singlets, at least those that make the right-handed counterparts of the active (left-handed) neutrinos.  However, the number of \\textit{sterile} neutrinos is still unknown. Unless some neutrino experiments are wrong, the present data on neutrino oscillations cannot be explained without sterile neutrinos. Neutrino oscillations experiments measure the differences between the squares of neutrino masses, and the results are: one mass squared difference is of the order of $10^{-5}$(eV$^2$), the other one is $10^{-3}$(eV$^2$), and the third is about $1\\,$(eV$^2$).  Obviously, one needs more than three masses to get the three different mass splittings which do not add up to zero.  Since we know that there are only three active neutrinos, the fourth neutrino must be sterile.  However, if the light sterile neutrinos exist, there is no compelling reason why their number should be limited to one.  Some theoretical arguments favor at least three sterile neutrinos \\citep{Asaka:2005an}.  If there are three sterile neutrinos, they can help explain the matter-antimatter asymmetry of the universe \\citep{Asaka:2005pn}.  Oscillations to sterile neutrinos add an intriguing additional consequence to the search for a neutron star kick mechanism; the opportunity to use supernovae as laboratories to study particle physics.  Other weakly interacting particles, for example, majorons, may cause the asymmetry as well \\citep{Far05}.  Supernova asymmetries can be used to discover or constrain a class of weakly interacting particles with masses below 100 MeV.  It is useful, therefore, to separate the details of a particular kick mechanism from its effects on the supernova, and to perform a model-independent analysis of how the non-ejecta kicks impact the rest of the supernova.  This is the main goal of the present paper.  The neutrino-driven explosion mechanism has evolved considerably since its introduction by \\cite{Col66}.  Although it is becoming increasingly accepted that convection above (and possibly within) the proto-neutron star can help make neutrino heating efficient enough to drive an explosion, the current state-of-the-art produces a range of results \\citep{Bur95,Jan96,Mez98,Fry99,Fry00,FW02, Bur03, FW04,Wal05}.  Over the past few years, a number of papers have studied ways to make the convection more vigorous, from asymmetries in the collapse \\citep{Bur96,Fry04} to instabilities in the accretion shock and a possible vortical-acoustic instability \\citep{Blo03,Sch04}.  The neutrino-driven kicks have the effect of breaking the spherical symmetry of the overall explosion, which may help stir the material and strengthen the convection.  In this paper, we study the effects of the neutrino-oscillation kick mechanism on the core-collapse engine.  We test its ability to help drive an explosion and study the observational implications of an explosion driven by the asymmetric emission from neutrinos. \\S 2 describes our computational set-up and the results of the simulations.  We find that, under some conditions, neutrino-driven kicks can affect the explosion.  In \\S 3, we study this effect and how it aids the explosion mechanism.  We conclude with a discussion of the observational implications from these effects and how these observations constrain what we know about neutrino oscillations. ",
        "conclusions": "To summarize, we have shown that pulsar kick mechanisms based on anisotropic emission of neutrinos or other weakly interacting particles from the cooling neutron star can increase the energy gained by the shock, hence improving the prospects for a successful explosion.  A distinguishing feature of this class of mechanisms is asymmetric explosion enhanced in the direction of the motion of the neutron star."
    },
    "0512/astro-ph0512205_arXiv.txt": {
        "abstract": "We report the first spatially resolved radio continuum measurements of the Mira~AB symbiotic binary system, based on observations obtained with the Very Large Array (VLA). This is the first time that a symbiotic binary has been resolved unambiguously at centimeter wavelengths. We describe the results of VLA monitoring of both stars over a ten month period, together with constraints on their individual spectral energy distributions, variability, and radio emission mechanisms. The emission from Mira~A is consistent with originating from a radio photosphere, while the emission from Mira~B appears best explained as free-free emission from an ionized circumstellar region $\\sim$(1-10)$\\times10^{13}$~cm in radius. ",
        "introduction": "Symbiotics are interacting binary systems in which a cool, evolved giant transfers material onto a hotter, compact companion through a stellar wind (e.g., Whitelock 1987). This symbiosis affects the late-stage evolution of both stars and their surrounding medium and may play an important role in shaping the formation of planetary nebulae and in the triggering of Type~Ia supernovae (e.g., Munari \\& Renzini 1992; Corradi et al. 2000). Mira~AB is the nearest example of a weakly symbiotic binary. It comprises a pulsating asymptotic giant branch (AGB) star (Mira~A=$o$~Ceti, the prototype of Mira variables) and a low-mass, accreting companion (Mira~B, possibly a white dwarf). The pair has a projected separation of $\\sim$\\as{0}{5} ($\\sim$65~AU; Karovska et al. 1997)\\footnote{All physical quanitities quoted in the this paper have been scaled to the {\\it Hipparcos} distance of 128~pc.}, making this one of the very few wind-accreting binaries in which the components can be spatially and spectrally resolved with current telescopes. This system therefore provides a rare opportunity to study the individual components of an interacting binary.  Multiwavelength observations of Mira AB have revealed a complex interacting system with tremendous temporal changes in the components and the circumbinary environment (e.g., Karovska et al. 1997,2001,2005; Wood \\& Karovska 2004 and references therein). Much of this activity is driven by mass-loss from Mira~A (${\\dot M}\\approx2.8\\times10^{-7}~M_{\\odot}$ yr$^{-1}$) through a cool, low-velocity wind ($V_{\\infty}\\approx 5$~\\kms; Bowers \\& Knapp 1988). Material from the wind is accreted onto Mira~B, forming a hot accretion disk, as evidenced by the presence of numerous rotationally broadened, ultraviolet (UV) emission lines (Reimers \\& Cassatella 1985).  {\\it Chandra} observations carried out in 2003 December detected an unprecedented X-ray outburst from Mira~A (Karovska et al. 2005). This was followed by an increase in UV emission and H$\\alpha$ flaring lasting for about one month. This outburst is  very unusual, since the X-rays seem to be originating from the AGB star rather than from the accretion disk of Mira~B. The outburst may be associated with a magnetic flare followed by a mass ejection event. To constrain the origin of the outburst and monitor its time evolution, we subsequently began multiwavelength monitoring of the Mira~AB system, including radio continuum measurements with the Very Large Array (VLA)\\footnote{The Very Large Array of the National Radio Astronomy Observatory is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.}. ",
        "conclusions": "Figure~1 shows several examples of our recently obtained radio continuum images of Mira~AB. Two continuum sources are clearly detected in each image. After accounting for proper motion, the two sources correspond with the positions of Mira A and B, respectively, as recently determined from {\\it HST} UV observations (Figure~2). This is the first time the components of a symbiotic binary have been resolved unambiguously at centimeter wavelengths.  For each observation, we measured the flux densities of the two stellar components by fitting elliptical Gaussians to the data in the image plane. Our results are summarized in Table~1. Figure~3 plots the derived flux densities of Mira~A and B as a function of time.  With one exception, we found Mira~A to be brighter than Mira~B at all times and at all frequencies. However, our 8.5~GHz measurements on JD~2,453,394 seem to indicate a simultaneous brightening of Mira~B and dimming of Mira~A compared with our previous measurements,  while nine days later, the flux densities of the two are nearly identical. We believe these apparent synchronous variations to be an artifact of undersampling the $u$-$v$ plane, and were able to reproduce an analogous ``redistribution'' in flux using pairs of similarly separated artificial point sources introduced into the visibility data. Note however that total flux density (A+B) was conserved. The observations in question were obtained using the BnA configuration of the VLA, which produced a highly elongated beam pattern, with the point spread functions for Mira~A \\& B partially overlapping at 8.5~GHz.  Our observations spanned $\\sim78$\\% of the 332 day pulsation period of Mira~A, during which the optical brightness of the star varies by $\\sim$6-7 magnitudes (e.g., Reid \\& Goldston 2002). Our measurements imply that any changes in the radio brightness of either star linked to these pulsations have a substantially smaller amplitude. At 8.5~GHz, where we have the longest observational baseline, the data are consistent with brightness changes of $\\lsim$30\\% from the mean for both Mira~A \\& B during the 189-day interval between 2004 October 19 and 2005 April 26 ($\\lsim$15\\% for Mira~A excluding the BnA array data; see above).  We have used a weighted mean of the measurements at each frequency where the  components of Mira~AB were resolved  to estimate the spectral indices of the two stars. Figure~4 shows the resulting radio frequency spectra. Both Mira~A and B have positive spectral indices, $\\alpha$ (where  flux density $S_{\\nu}\\propto \\nu^{\\alpha}$), indicating that the emission mechanism for both sources is likely to be predominantly thermal. Nonlinear least-squares fits to the data show the radio spectra of Mira~A and B can be described as $S_{\\nu,A}=(0.009\\pm0.004)\\nu^{1.50\\pm 0.12}_{\\rm GHz}$~mJy and $S_{\\nu,B}=(0.010\\pm0.008)\\nu^{1.18\\pm0.28}_{\\rm GHz}$~mJy, respectively.  The initial goal of our radio observations of Mira~AB was to search for radio signatures of the recent X-ray outburst detected by Karovska et al. (2005). From the continuum observations obtained to date, we have not detected any unambiguous aftereffects of this event. Some of our radio images of Mira~A (e.g., Figure~1) show elongations along a position angle of $\\sim120^{\\circ}$, analogous to those seen in X-rays and in the UV (Karovska et al. 2005). These features could be consistent with an outflow; however, given the phase noise and limited $u$-$v$ sampling of our data (observations $\\lsim$1 hour), they may also be spurious. If the X-ray outburst of Mira~A was connected with a mass ejection event and/or shock formation, we might also expect a brightening in one or both components of Mira~AB. We detected some signs of statistically signficant flux density fluctuations in Mira~A and in the total flux from the binary at 8.5, 22.5, and 43.3~GHz; however, these are generally limited to a single frequency on a given date, and cannot be unambiguously linked with the outburst. Future monitoring should reveal whether similar-scale fluctuations  are the norm for the system. We note that  material ejected from Mira~A at speeds of $\\sim$200~\\kms\\ (see Karovska et al. 2005) would require $\\gsim$1.5 years to reach the accretion disk of Mira~B. Therefore radio signatures of such an event may still occur during late 2005 or 2006.  To gauge the longer-term variability of the Mira~AB system, we have also analyzed unpublished archival 8.5~GHz data taken in 1996 December using the VLA in its A configuration. Both Mira~A and Mira~B were detected with flux densities $S_{A}=0.28\\pm 0.05$~mJy and $S_{B}=0.15\\pm 0.04$~mJy, respectively. These values are consistent with our recent measurements, and imply changes in the mean brightness of $\\lsim$30\\% for Mira~A and $\\lsim$20\\% for Mira~B during the past 8.5 years.  Previous VLA 8.4~GHz measurements of Mira~A have been published by Reid \\& Menten (1997), based on data obtained during 1990. However, Mira A \\& B were spatially unresolved in these observations, thus the authors likely measured the {\\it combined} emission from the two binary components. Indeed, their published flux densities agree  with the values we derive from the sum of Mira~A+B. Reid \\& Menten (1997) interpreted the radio emission they observed from Mira and other long period variable stars as arising from radio photospheres located at $R\\approx2R_{\\star}$. Their models predict that at centimeter wavelengths ($\\sim$8-22~GHz) this emission should exhibit a spectral index $\\alpha\\approx$1.86. This is slightly steeper than the  value we derive for Mira~A (Figure~4). However, after accounting for calibration uncertainties, it is consistent with our data in the range 8.5-22.5~GHz. At higher frequencies, a gradual turnover may be expected owing to opacity changes, consistent with our 43.3~GHz measurements. The Reid \\& Menten model also predicts an absolute 8.5~GHz flux density for Mira~A in agreement with our mean observed value. Our new unblended radio continuum measurements of Mira~A therefore remain consistent with the emission  arising primarily from a radio photosphere. The limits on brightness fluctuations of $\\lsim$30\\% over the course of several months furthermore place a limit on shocks or other disturbances in the radio photosphere of Mira~A (e.g., resulting from pulsations or the X-ray outburst event)  to speeds of less than a few \\kms\\ (see Reid \\& Menten 1997).  In the case of Mira~B, thermal emission from the stellar photosphere and/or accreting surface will be undetectable at centimeter wavelengths owing to the small emitting area. However, Lyman continuum photons from its hot accretion disk or boundary layer may ionize a portion of Mira~A's wind, providing a source of free-free emission. In general, free-free radiation from a fully ionized stellar wind results in  radio emission with a spectral index $\\alpha=0.6$ (Wright \\& Barlow 1975). Our measured spectral index for Mira~B is a factor of two steeper than this, and is also steeper than values previously measured for other (unresolved) symbiotics over this frequency range (Seaquist \\& Taylor 1990). However, in the limit where only a very small fraction of the wind is ionized, the model of Seaquist et al. (1984) and Taylor \\& Seaquist (1984) (the so-called ``STB'' model) predicts $\\alpha\\rightarrow$1.3-2.0 across the optically thick portion of the spectrum, consistent with our measurements. Assuming a canonical electron temperature of $T_{e}=10^{4}$~K (Osterbrock 1989), the  radius of an optically thick ionized sphere required to produce the observed $\\nu$=8.5~GHz emission from Mira~B can be estimated as $R=[S_{\\nu}c^{2}d^{2}/2kT_{e}\\nu^{2}\\pi]^{0.5}\\approx 1.6\\times10^{13}$~cm ($\\sim$1~AU) where $c$ is the speed of light and $k$ is Boltzmann's constant. This suggests an ionized volume much smaller than the binary separation.  An alternate constraint on the size of an ionized region surrounding Mira~B may be obtained by incorporating results from previous studies. Based on UV spectroscopy, Reimers \\& Cassatella (1985) estimated a radius and thickness for the Mira~B accretion disk of $r\\sim1.7\\times10^{11}$~cm and $t\\sim6.8\\times10^{9}$~cm, respectively, and an electron temperature $T_{e}=11,000$~K. A blackbody with this emitting area and temperature will produce $N_{\\rm UV}\\approx5\\times10^{41}$ hydrogen-ionizing photons per second. Assuming a separation between Mira~A \\& B of $a$=100~AU (slightly larger than the projected separation), we can also estimate the particle density from the wind of Mira~A at the location of Mira~B as $n_{e}={\\dot M}/4\\pi\\mu m_{\\rm H}V_{\\infty}a^{2}$ (Wright \\& Barlow 1975, Eq.~2), where ${\\dot M}$ and $V_{\\infty}$ are Mira~A's wind parameters, $\\mu$ is the mean molecular weight ($\\sim$1), and $m_{H}$ is the mass of a hydrogen atom. This predicts $n_{e}\\approx7.2\\times10^{5}$~cm$^{-3}$, neglecting density enhancements caused by the gravitational field of Mira~B. In the idealized case of a pure hydrogen medium of uniform density, the size of the resulting ionized region as predicted by the Str\\\"omgren relation (e.g., Osterbrock 1989) is $R_{s}\\approx(3N_{\\rm UV}/4\\pi\\alpha_{r}n^{2}_{e})^{\\frac{1}{3}}\\approx9.8\\times10^{13}$~cm. Here $\\alpha_{r}$ is the recombination coefficient to  all levels but the ground state of hydrogen ($\\approx2.6\\times10^{-13}$~cm$^{3}$~s$^{-1}$; Osterbrock 1989). Although this order-of-magnitude estimate neglects such factors as a non-uniform density within the ionized volume and the presence of molecules and heavier elements in the wind, it  supports the STB model as an explanation for the radio emission from Mira~B and reaffirms that any ionized nebula around the star should be quite small compared with the binary separation and the overall extent of Mira~A's circumstellar envelope ($>10^{17}$~cm; Bowers \\& Knapp 1988). It is also consistent with the nebula being unresolved by  our recent VLA observations (predicted angular diameter $\\theta\\lsim$\\as{0}{1}) . Future VLA measurements should provide additional constraints on the size and shape of this ionized region, and ultimately place new, independent constraints on the Mira~B accretion disk properties and the true binary separation.  We have considered other possible mechanisms as the source of the radio emission from Mira~B, but none appear likely to contribute appreciably to the observed flux. For example, if Mira~B is a magnetically active dwarf rather than a white dwarf (Jura \\& Helfand 1984; Kastner \\& Soker 2004), its quiescent radio luminosity is expected be $\\sim10^{12}$-$10^{14}$ erg s$^{-1}$ Hz$^{-1}$ (Benz \\& G\\\"udel 1994)---undetectable at Mira~B's distance. Mira~B is also known to power a variable wind  (Wood et al. 2002), but even if fully ionized, it is too feeble to produce detectable radio emission. Based on the parameters derived by Wood et al. (2002): ${\\dot M}\\approx$(0.14-2.8)$\\times10^{-11}~M_{\\odot}$ yr$^{-1}$ and $V_{\\infty}$=250-400~\\kms, the predicted radio flux density from Mira~B's wind (see Wright \\& Barlow 1975, Eq.~8) would be $\\sim$4 orders of magnitude smaller than the observed emission. Finally, the steep spectral index  of the radio emission from Mira~B ($\\alpha>$0.6) and low energy flux of its wind argue against any significant contribution from non-thermal emission produced by collisions between the two winds of the binary (see Dougherty \\& Williams 2000; Kenny \\& Taylor 2005).  In summary, we have presented new centimeter wavelength images of the symbiotic binary system,  Mira~AB. These data allow for the first time measurements of the radio properties of the individual components of a symbiotic binary. Over the frequency range 8-43~GHz, the radio properties of the evolved giant, Mira~A, are consistent with the radio emission originating predominantly from a radio photosphere. The emission from Mira~B is consistent with arising from a circumstellar \\HII\\ region $\\sim$(1-10)$\\times10^{13}$~cm in radius. We find the radio variability of both stars to be $\\lsim$30\\% over a ten month period during 2004-2005. Flux densities we derived from archival data taken in 1996 are consistent with our recent measurements, and imply changes in the mean 8.5~GHz flux densities of $\\lsim$30\\% for Mira~A  and $\\lsim$20\\% for Mira~B during the past 8.5 years."
    },
    "0512/astro-ph0512175_arXiv.txt": {
        "abstract": "In the context of recent observational results that show massive ellipticals were in place at high redshifts, we reassess the status of monolithic collapse in a $\\Lambda$CDM universe. Using a sample of over 2000 galaxies from the Sloan Digital Sky Survey, by comparing the dynamical mass and stellar mass (estimated from colours) we find that ellipticals have `cores' which are baryon-dominated within their half-light radius. These galaxies correspond to 3-sigma peaks in the spherical collapse model if the total mass in the halo is assumed to be 20 times the dynamical mass within the half-light radius. This value yields stellar mass to total mass ratios of 8\\%, compared to a cosmological baryon fraction of 18\\% derived from WMAP3 alone. We further develop a method for reconstructing the concentration halo parameter $c$ of the progenitors of these galaxies by utilizing adiabatic contraction. Although the analysis is done within the framework of monolithic collapse, the resulting distribution of $c$ is log-normal with a peak value of $c\\sim3-10$ and a distribution width similar to the results of N-body simulations. We also derive scaling relations between stellar and dynamical mass and the velocity dispersion, and find that these are sufficient to recover the tilt of the fundamental plane. ",
        "introduction": "The large new galaxy surveys such as 2dF \\citep{2dFmain} and the Sloan Digital Sky Survey (SDSS) \\citep{York} have transformed the way we can study galaxy properties statistically.  Here we focus on a sample of elliptical galaxies derived from those selected from the SDSS by \\citet{Bernardi1} (hereafter B03).  Our motivation is to revisit fundamental issues of galaxy formation from the perspective afforded by a modern data set.  A theory of galaxy formation in a universe dominated by Cold Dark Matter was first set out in a seminal paper by Blumenthal, Faber, Primack \\& Rees (1986a) (hereinafter BFPR) which considers the monolithic collapse of isolated dark matter halos. According to this picture, the redshift at which a halo collapses is determined only by its mass, for a given cosmology and choice of amplitude of fluctuation. The baryons follow the dark matter distribution until the radius of the collapsing halo reaches the virial radius and is halted.  Studies of the evolution of the fluctuations in density, however, suggest that galaxy formation is dominated by merging of small dark matter halos \\citep{WhiteRees}, and this picture is supported by simulations \\citep{Springel}. There have been many attempts to discriminate between these two pictures of the formation of galaxies, referred to respectively as monolithic or spherical collapse and hierarchical merging. Simulation results notwithstanding, the small scatter of the observed colour magnitude relation and its evolution with redshift provides evidence that massive ellipticals were already in place at a redshift of $z\\approx 1-2$ with little subsequent merging (e.g. \\citet{sed98}), whereas less massive ellipticals present features characteristic of recent star formation (e.g. \\citet{FS00}).  This `inverted hierarchy' or `downsizing' effect \\citep{Cow96} illustrates the complexity of galaxy formation compared with a simple picture of the assembly of dark matter halos.  In this paper we use the predictions of the spherical collapse scenario as a benchmark, which can then be challenged with further comparisons with observations and detailed simulations.  A new aspect of our analysis is the estimation of the stellar mass of each of the B03 ellipticals, allowing us to distinguish between baryonic and total mass.  We also update the calculations of BFPR for `concordance cosmology' (matter density $\\Omega_m=0.3$,dark energy density $\\Omega_\\Lambda=0.7$,Hubble constant $h=0.7$ and amplitude of fluctuations $\\sigma_8=0.9$) \\footnote{We note that the recent WMAP release prefers $\\Omega_{M}\\approx 0.24$ and $\\sigma_8 \\approx 0.74$. However, $\\sigma_8=0.9$ is still preferred by weak lensing.} and place the observed ellipticals on a revised cooling diagram. Despite the attention BFPR has received, this relatively simple generalization to $\\Lambda$CDM has not previously appeared in the literature. Our analysis of the B03 galaxies also appears to be the first attempt to incorporate data for individual galaxies (as opposed to schematic data) in this parameter space. From the first three years of WMAP observations of the cosmic microwave background alone (Spergel et al. 2006), $\\Omega_m=0.24^{+0.03}_{-0.04}$ and $\\Omega_b=0.042^{+0.003}_{-0.005}$. For simplicity, we use a `cosmological' baryon fraction of 1/6 unless otherwise stated.  On galactic scales, \\citet{Klypin} model the Milky Way within the virial radius ($\\approx 250$ kpc) and find that a substantial `feedback' mechanism which removes baryons from the galaxy itself must have operated. Do similar processes operate in massive ellipticals? Are their cores dominated by baryonic matter?  \\citet{Romanowsky} and \\citet{Dekeletal} have reached conflicting conclusions to this important question, based on the analysis of planetary nebul\\ae~in elliptical galaxies out to five times the effective radius $R_e$. By using the central velocity dispersion and stellar mass out to $R_e$ we seek to answer the same question by focussing on the central region only. We find that the cores of the galaxies are baryon dominated, at least to a distance of $R_e$ from the centre.   By making the straightforward assumption that the total mass of the galaxy is proportional to the mass of the galaxy `cores' we have studied, we reproduce the BFPR results for a modern `concordance' cosmology, and investigate the regime in which their benchmark model of spherical collapse is consistent with the data. We also recover the profile of the undisturbed dark matter halo from present-day observables via the procedure of adiabatic contraction, which enables us to compare our derived profile concentration with both simulations and observations. Defining $\\alpha=M_{\\rm tot}$/$M_{\\rm dyn}\\left(<R_e\\right)$ - the total mass of the collapsing halo divided by the dynamical mass within $R_e$ -- we find a strong constraint that $\\alpha \\ge 10$. $\\alpha \\approx 20$ provides a good fit to the data.   The outline of the paper is as follows: Section 2 presents the B03 sample and a volume-limited subset of 2040 galaxies selected from it.  Section 3 describes the derivation of stellar masses of these ellipticals from the SDSS colours.  Section 4 derives dynamical masses and Section 5 covers the relations between luminosity, velocity dispersion, stellar mass and dynamical mass, extending the classic \\citet{FJ76} relation.  In Section 6 we comment on the acceleration of stars in the ellipticals and the connection to MOND. Section 7 extends the benchmark spherical collapse model for the concordance $\\Lambda$-CDM cosmology and use it to contrast the observed ellipticals with the level of rms fluctuations on the cooling diagram.  Finally, in Section 8 we present a new reconstruction method which recovers the initial halo of dark matter from the present-day observables using the adiabatic contraction.  We summarize the results and outline future work in Section 9. ",
        "conclusions": "The sample of early-type galaxies selected by Bernardi et al. (2003a) from the Sloan Digitial Sky Survey has been used to investigate the properties of baryonic and dark matter within these systems. For the work presented in this paper, the sample was further restricted to be volume limited, containing only massive elliptical systems, leaving a total of more than 2000 systems. Using photometry from Sloan Data Release 4 and the synthetic stellar populations of Bruzual \\& Charlot (2003) we calculated the mass in stars of each system.   We recover the well-known Faber-Jackson relation between luminosity and observed velocity dispersion and also find a tight correlation between stellar mass and velocity dispersion, which we have called the `baryonic Faber-Jackson relation'. We find a slope of  $M_{\\mathrm{star}}\\left(<R_e\\right) \\propto \\sigma^{2.33 \\pm 0.18}$ (in appropriate units), less steep than previous work \\citep{Thomasetal} based on a study of an order of magnitude fewer galaxies.   In addition to determining the stellar mass, it is possible to determine a dynamical mass via the velocity dispersion and the radius of the system. Comparing the two estimates of mass reveals that dark matter is underabundant in the centre of such galaxies when compared to the cosmological baryon fraction. However, our results are in good agreement with observers who have claimed that the inner regions of massive elipticals are baryon dominated.   In order to examine whether this large data set is sufficient to discriminate between galaxy formation via monolithic collapse or the merging of small dark matter haloes, the data were compared to a simple model of spherical collapse. It was necessary to assume that the total mass of the systems under study is proportional to the mass within the baryon-rich core (within $R_e$), but given such an assumption we find that for $\\alpha \\approx 20$ (where $\\alpha=M_{\\mathrm{tot}}$/$M_{\\mathrm{dyn}}\\left(<R_e\\right)$), the data were found to be consistent with simple monolithic collapse of 3-sigma peturbations (assuming $\\Lambda$CDM cosmology).  Interestingly, this choice of $\\alpha$ also produces a value for the fraction of the mass which lies in the form of cooled baryons which is consistent with other work. A choice of 20 seems large, but corresponds to the stellar mass contributing half of the cosmological baryon fraction. Given the other reservoirs of baryons in these systems, this is a reasonable value.  Having motivated our choice of $\\alpha$, we used the adiabatic contraction method first due to \\citet{BFFP} to calculate the response of the dark matter to the presence of the observed baryons. We assume that the dark matter has an initial NFW profile, and produce the distribution of initial concentration shown in figure \\ref{fig:c}. The mean concentration is lower than that found in dissipation-less N-body simulations.   In summary, we have shown that the B03 sample of early-type galaxies is a powerful resource for studying the distribution of matter within these objects. However, despite this excellent data set we cannot rule out spherical collapse as a mode of formation for massive ellipticals. We find that the central regions of the galaxies, which are baryon dominated, must represent only a small fraction of the total mass of the system. In future work, we will extend our analysis to a new large sample  of early-type systems, and undertake a detailed comparision of simulation results in order to understand the differences between observations and semi-analytic models."
    },
    "0512/astro-ph0512496_arXiv.txt": {
        "abstract": "We present observations of the 2002 September 30 white-light flare, in which the optical continuum emission near the \\Ha\\ line is enhanced by $\\sim$10\\%. The continuum emission exhibits a close temporal and spatial coincidence with the hard X-ray (HXR) footpoint observed by {\\it RHESSI}. We find a systematic motion of the flare footpoint seen in the continuum emission; the motion history follows roughly that of the HXR source. This gives strong evidence the this white-light flare is powered by heating of nonthermal electrons. We note that the HXR spectrum in 10--50 keV is quite soft with $\\gamma \\approx 7$ and there is no HXR emission above 50 keV. The magnetic configuration of the flaring region implies magnetic reconnection taking place at a relatively low altitude during the flare. Despite a very soft spectrum of the electron beam, its energy content is still sufficient to produce the heating in the lower atmosphere where the continuum emission originates. This white-light flare highlights the importance of radiative backwarming to transport the energy below when direct heating by beam electrons is obviously impossible. ",
        "introduction": "Compared to ordinary solar flares, white-light flares (WLFs) manifest themselves with an enhanced emission in the optical continuum of a few or tens of percent, or even more in some extreme cases. The continuum emission originates in the lower chromosphere and below. Therefore, research on WLFs can reveal critical clues to energy transport and heating mechanisms in the solar lower atmosphere \\citep{nei89, din99}.  Most of the WLFs observed so far belong to Type I WLFs that are spectrally characterized with a Balmer and Paschen jump and strong broadened hydrogen Balmer lines \\citep{mac86}. Observations usually demonstrate a close temporal correspondence between the continuum and the hard X-ray (HXR) emissions in these WLFs \\citep[e.g.,][]{hud92, fan95, nei93, mat03, met03, hud05}. This implies that the continuum emission is closely related to nonthermal heating of energetic electrons, which are assumed to be accelerated in the corona and then stream downward to the chromosphere. However, this process alone, in most cases, cannot efficiently produce the continuum emission since very few electrons ($E$ $\\geq$ 200 keV) can penetrate to the lower chromosphere and below \\citep{din03b}. Therefore, radiative backwarming is further invoked to transfer the enhanced chromospheric radiation to deeper layers to produce the heating there and finally an enhanced continuum emission \\citep{hud72, abo87, mac89, met90a, met90b,gan94,din03b}. A recent radiative hydrodynamic model of solar flares, which includes the electron beam heating and radiative transport in a consistent way, reproduces the continuum enhancement comparable to the observed values \\citep{all05}. On the other hand, observations show that in some few very energetic events, electrons can be accelerated to very high energies, e.g., 300--800 keV \\citep{xu04,xu05}. In such cases, a large amount of nonthermal electrons can penetrate deep into the atmosphere, and direct collisional heating may play an important role in producing the continuum emission.  It has been widely accepted that nonthermal electrons, accelerated by magnetic reconnection in the corona, stream downward along magnetic fields and produce the HXR emission in the chromosphere via bremsstrahlung radiation \\citep{bro71}. According to this scenario, the apparent motions of HXR sources reflect the successive reconnection process between neighboring field lines in the corona \\citep[e.g.,][]{qiu02}. Observations of footpoint motions can thus provide a test or constraint on theoretical models of solar flares. There are, up to now, quite some observations revealing the motions of HXR sources indicative of successive reconnection under various magnetic topologies \\citep[e.g.,][]{sak00, kru03}. However, few examples have shown the motions of white-light continuum kernels that are related with the HXR sources.  In this paper, we present observations of a white-light flare on 2002 September 30. We examine the temporal and spatial relationship between the continuum emission and the HXR emission during the flare and discuss its energetics. In particular, We discover a fairly good correlation between the motion history of the white-light kernel and the corresponding HXR source. We explain the observed continuum emission in terms of nonthermal heating by electron beams followed by the radiative backwarming effect. ",
        "conclusions": "After a synthesising analysis of the WLF on 2002 September 30, we find that there is a fairly close relationship in both time and space between the continuum and the HXR emission during the flare. We discover a footpoint motion seen in the optical continuum; the motion history follows roughly that of the HXR source. This gives strong evidence that this WLF is powered by energetic electrons followed by the radiative backwarming effect. The magnetic configuration of the flaring region implies magnetic reconnection taking place at a relatively low layer during the flare. The HXR spectrum is thus very soft, possibly, due to a low efficiency of electron acceleration at low altitudes. However, the energy content of the electron beam is still enough to produce the heating in the lower atmosphere where the continuum emission originates. Radiative backwarming is of course an important mechanism to transport the energy below when direct heating is obviously impossible."
    },
    "0512/astro-ph0512343_arXiv.txt": {
        "abstract": "The XBo\\\"otes Survey is a 5-ks \\chandra\\ survey of the Bo\\\"otes Field of the NOAO Deep Wide-Field Survey (NDWFS). This survey is unique in that it is the largest (9.3 deg$^2$), contiguous region imaged in X-ray with complementary deep optical and near-IR observations. We present a catalog of the optical counterparts to the 3,213 X-ray point sources detected in the XBo\\\"otes survey. Using a Bayesian identification scheme, we successfully identified optical counterparts for 98\\% of the X-ray point sources. The optical colors suggest that the optically detected galaxies are a combination of $z<$1 massive early-type galaxies and bluer star-forming galaxies whose optical AGN emission is faint or obscured, whereas the majority of the optically detected point sources are likely quasars over a large redshift range. Our large area, X-ray bright, optically deep survey enables us to select a large sub-sample of sources (773) with high X-ray to optical flux ratios (f$_x$/f$_o>$10). These objects are likely high redshift and/or dust obscured AGN. These sources have generally harder X-ray spectra than sources with 0.1$<$f$_x$/f$_o<$10. Of the 73 X-ray sources with no optical counterpart in the NDWFS catalog, 47 are truly optically blank down to $R\\sim$25.5 (the average 50\\% completeness limit of the NDWFS $R$-band catalogs). These sources are also likely to be high redshift and/or dust obscured AGN. ",
        "introduction": "Active Galactic Nuclei (AGN) are complex objects which radiate across the entire electromagnetic spectrum. To gain a better insight into their nature and how they evolve, we require large samples of AGN in multiple wave-bands. Previous studies at optical and soft X-ray wavelengths have failed to obtain a complete census of the AGN population because dust and gas obscures many AGN from view. In fact in the local Universe, optically obscured AGN may outnumber optically unobscured AGN by a factor of four \\citep{mai95}. In all but the shallowest surveys, the number density of AGN identified in the hard X-ray and mid-IR bands is far greater than is found in optical surveys of comparable depth (\\citealt{bau04}; \\citealt{ris04}; \\citealt{ste05}). Thus it is important to determine the extent to which there exists a hidden population of AGN whose optical properties do not identify them as AGN. Because of the ability of hard X-rays to penetrate dust and all but the most extreme column densities of gas, X-ray surveys with sensitivities above $\\sim$4 keV can provide relatively unbiased samples of AGN over a range of redshifts.  We summarise some of the existing extragalactic X-ray surveys in Table~\\ref{tab:xsurveys}. The \\chandra\\ \\citep{wei02} and {\\it XMM-Newton} \\citep{jan01} observatories have led to great advances in X-ray astronomy, producing surveys which are 10 - 100 times deeper than those by previous X-ray telescopes. \\chandra 's imaging resolution is superior to all previous and current X-ray telescopes (\\citealt{van97}; \\citealt{gar03}); this is crucial for determining the correct optical counterpart in deep optical imaging data in which the surface density of sources is large. While the deepest surveys have resolved $\\sim$80\\% of the hard (2-7 keV) X-ray background into discrete sources (\\citealt{bra01}; \\citealt{ros02}; \\citealt{mor03}; \\citealt{wor04}) and an even larger fraction of the soft (0.5-2 keV) X-ray background, they typically cover only small ($\\sim$0.1 deg$^2$) areas of the sky. Surveys covering a larger volume are needed to overcome cosmic variance and to better determine the properties of the most luminous and rarest sources such as powerful quasars, whose number densities are low. Large contiguous areas and extensive multi-wavelength coverage are also necessary for detailed studies of AGN clustering and environment.  \\begin{deluxetable*}{lllll} \\tabletypesize{\\scriptsize} \\setlength{\\tabcolsep}{0.02in} \\tablecolumns{4} \\tablewidth{0pc} \\tablecaption{\\label{tab:xsurveys} Properties of recent extragalactic deep and medium-deep X-ray surveys ordered by increasing survey area. The X-ray flux is quoted for the 0.5-2 keV band and has been converted to the flux assuming a power-law with $\\Gamma$=1.7 where needed. The optical magnitude limit is taken from the referenced paper and may be defined in different ways. {\\bf a.} The LALA field is in the north-east corner of the Bo\\\"otes field. {\\bf b.} The number in parenthesis indictates the total expected area of the survey when completed.} \\tablehead{ \\colhead{Survey} & \\colhead{Area} & \\colhead{X-ray Flux limit} & \\colhead{$R$ limit} & \\colhead{Reference} \\\\ \\colhead{} & \\colhead{deg$^2$} & \\colhead{${\\rm ergs~s^{-1}~cm^{-2}}$} & \\colhead{magnitudes} & \\colhead{} } \\startdata Chandra Deep Field North & 0.1 & 2.4$\\times 10^{-17}$ & $\\sim$26.4 &\\citet{bar03}, \\citet{ale03}\\\\ Chandra Deep Field South & 0.1 & 5.7$\\times 10^{-17}$ & $\\sim$26.5 &\\citet{gia02}\\\\ Extended Chandra Deep Field South & 0.3 & 8$\\times 10^{-17}$ &$\\sim$26.5 & \\citet{leh05}\\\\ Chandra Extended Growth Strip (DEEP)& 0.1 & 1.0$\\times 10^{-16}$ & 23.4 & \\citet{nan05}\\\\ Chandra LALA$^{\\bf a}$ & 0.1 & 1.4$\\times 10^{-16}$ & 25.7 & \\citet{wan04}\\\\ SPICES-II & 0.1 & 1.8$\\times 10^{-16}$ &25.4 & \\citet{ste02}\\\\ XMM Lockman hole & 0.25 & 3.0$\\times 10^{-16}$ & $\\sim$27.1 &\\citet{has01},\\citet{wil04}\\\\ CLASXS & 0.36 & 5.0$\\times 10^{-16}$ & 27.0 & \\citet{yan04}\\\\ Chandra survey of 13h XMM/ROSAT  &0.25 & 5.0$\\times 10^{-16}$& 27.0 &\\citet{mch03}\\\\ Chandra Multi-wavelength Project & 0.8 (14.0$^{\\bf b}$) & variable & 24.0 & \\citet{kim04}, \\citet{gre04}\\\\ ROSAT Ultra Deep & 0.36 & 1.1$\\times 10^{-15}$ & 25.5 & \\citet{leh01}\\\\ {\\bf XBo\\\"otes} & {\\bf 9.3} & ${\\bf 4.0\\times 10^{-15}}$ & {\\bf 25.5 } & Murray et al.~(2005), Kenter et al.~(2005); this work\\\\ XMM-Newton/2dF & 1.5 & 4.0$\\times 10^{-15}$& 19.5 & \\citet{geo03}\\\\ HELLAS2XMM & 0.9 (4.0$^{\\bf b}$)&7.5$\\times 10^{-15}$& 24.5 & \\citet{bal02}\\\\ \\hline \\enddata \\end{deluxetable*}   Here we identify the candidate optical counterparts to the X-ray point sources in a new, medium depth (5-ks/pointing), wide-field X-ray survey performed with \\chandra\\  ACIS-I (XBo\\\"otes; Murray et al. 2005) of the Bo\\\"otes field of the NOAO Deep Wide-Field Survey (NDWFS; \\citealt{jan99}; Jannuzi et al. in prep.; Dey et al. in prep.). The XBo\\\"otes survey is unique because of its contiguous imaging coverage over a large (9.3 deg$^2$) area with {\\it Chandra}. While not as deep as many past surveys made with {\\it Chandra} and {\\it XMM}, the large co-moving volume surveyed allows us to obtain large samples of sources comprising the bright end of the X-ray luminosity function and to determine the nature of rare populations whose number densities are too small to obtain meaningful statistics in small area surveys. The large contiguous area is also critical for studies of AGN clustering: one of the key goals of the XBo\\\"otes survey. We can also perform stacking analyzes of the X-ray data to determine the mean X-ray properties of different populations selected in different wavebands (e.g., \\citealt{bra05}; Watson et al.~in prep.). The Bo\\\"otes field is unique in the extent of its multi-wavelength coverage over such a large area. The multi-wavelength data includes X-ray (\\chandra), UV ({\\it GALEX}; \\citealt{hoo04}), optical (NDWFS), near-IR (FLAMEX; \\citealt{els05}), mid-IR ({\\it Spitzer}/IRAC --  \\citealt{eis04}; {\\it Spitzer}/MIPS --  \\citealt{soi04}), and radio (VLA/FIRST; \\citealt{bec95} and WSRT; \\citealt{dev02}). Optical spectroscopic follow-up observations have also been undertaken for all X-ray sources with $I<$21.5 as part of the AGES survey (Kochanek et al. in prep.). %  In Murray et al.~(2005; Paper I), we described the general characteristics of the X-ray survey and determined the angular clustering of the sources. In Kenter et al.~(2005; Paper II) we presented the X-ray catalog. The main aim of this paper is to provide an accurate catalog of the optical counterparts of the X-ray point sources in the Bo\\\"otes region which were presented in Kenter et al.~(2005). We begin with a summary of the X-ray, optical, and near-IR data used in the XBo\\\"otes and NDWFS surveys (Section~2). We discuss our method of associating the X-ray source with optical counterparts in Section~3. The robustness of the resulting matched catalog as well as the properties of the catalog are described in Section~4. In Section~5, we describe the X-ray and optical properties of the matched catalog. The main points of the paper are summarized in Section~6. ",
        "conclusions": "We have presented the catalog of optical counterparts of the 3,213 X-ray point sources in the Bo\\\"otes field of the NOAO Deep Wide-Field Survey. Our Bayesian technique finds optical counterparts for 98\\% of the X-ray sources. This provides the basis for further investigation of the properties of the X-ray sources. In particular, a large program of optical spectroscopy with the Hectospec instrument on the MMT is underway, targeting all ($\\approx$1900) X-ray sources with $R<$21.5 or $I<$21.5 optical counterparts (AGES; Kochanek et al. in prep.).  Approximately half of the XBo\\\"otes sources are identified with optical point sources;  the other half have extended optical counterparts dominated by extended galaxy light.  The bright ($R<23.0$) optical point sources fall in a region of optical color-color space consistent with them being quasars at $z\\le$3. The fainter point sources appear to be redder in their optical colors, suggesting that they may be more obscured QSOs.  The optical colors of the extended sources suggest that they are a combination of $z<$1 early-type galaxies and bluer star-forming galaxies.  The large area, X-ray shallow and optically deep nature of the XBo\\\"otes survey allows us to identify a large sample of 773 bright X-ray sources with high X-ray to optical flux ratios ($f_x/f_o>10$). We interpret these large X-ray to optical flux ratios as resulting from extinction of the optical light by dust. These X-ray sources are generally harder in their X-ray spectra than the bulk of the X-ray source population ($0.1<f_x/f_o<10$), suggesting that AGN that are obscured by dust in the optical are also more likely to be obscured in the soft X-ray by gas. Such a large sample of bright X-ray sources with extreme properties holds much promise for follow-up multiwavelength spectroscopy and detailed studies of obscured AGN. An extreme subset of the $f_x/f_o>10$ population are 47 X-ray sources which have no optical identification down to $R\\sim 25.5$, the average 50\\% completeness limit of the NDWFS $R$-band images."
    },
    "0512/astro-ph0512425_arXiv.txt": {
        "abstract": "The phenomena customly called Dark Matter or Modified Newtonian Dynamics (MOND) have been argued by Bekenstein (2004) to be the consequences of a covariant scalar field, controlled by a free function (related to the MOND interpolating function $\\mut(g/a_0)$) in its Lagrangian density.  In the context of this relativistic MOND theory (TeVeS), we examine critically the interpolating function in the transition zone between weak and strong gravity. Bekenstein's toy model produces too gradually varying $\\mut$ and fits rotation curves less well than the standard MOND interpolating function $\\mut(x)=x/\\sqrt{1+x^2}$.  However, the latter varies too sharply and implies an implausible external field effect (EFE).  These constraints on opposite sides have not yet excluded TeVeS, but made the zone of acceptable interpolating functions narrower. An acceptable ``toy\" Lagrangian density function with simple analytical properties is singled out for future studies of TeVeS in galaxies. We also suggest how to extend the model to solar system dynamics and cosmology, and compare with strong lensing data (see also astro-ph/0509590). ",
        "introduction": "On galaxy scales, dark matter generally dominates over baryons (stars plus gas) at large radii.  At intermediate radii in a galaxy where dark matter and baryons overlap with comparable amounts, the two mass profiles are not uncorrelated (McGaugh 2005). The correlation between the Newtonian gravity of the baryons ${\\bf g}_N$ and the overall gravity ${\\bf g}$ (baryons plus dark matter) can be loosely parameterized by the Milgrom's (1983) empirical relation \\begin{equation}\\label{mueq} ~{\\mut}(g/a_0) \\, {\\bf g} = {\\bf g}_{{\\rm N}}, \\end{equation} where the interpolating function $\\mut(x)$ is a function which runs smoothly from $\\mut(x)=x$ at $x\\ll 1$ to $\\mut(x)=1$ at $x\\gg 1$ with a dividing gravity scale $a_0 \\sim 10^{-8} {\\rm cm}\\,{\\rm s}^{-2} \\sim cH_0/6$ at the transition.  This simple correlation was taken as the basis for the MOND theory (or more precisely the aquadratic Lagrangian theory) by Bekenstein \\& Milgrom (1984, hereafter BM84), where one modifies the Newtonian gravity of a baryonic galaxy to eliminate the need for dark matter.   Recently interests on the subject of MOND have been further stimulated since Bekenstein (2004, hereafter B04) provided a Lorentz-covariant theory (dubbed TeVeS), which passes standard tests to check General Relativity, and allows for rigourous modeling of Hubble expansion and gravitational lensing (e.g. Zhao et al. 2005). In TeVeS the MONDian behaviour originates from a scalar field, the dynamics of which is controlled by a Lagrangian density involving a free function that yields the expected dynamics in the low-acceleration limit (although BM84 theory is not precisely recovered). This freedom of the Lagrangian density, that echoes the freedom in the choice of the interpolating function $\\mut$ in MOND, means that every choice of the free function defines a distinct theory. As this class of theories do not at present derive from any basic principle and are purely phenomenological, the only constraints on the free function must come from phenomenological grounds. A refinement of the function studied by B04 as a toy model is surely needed. In this letter, we differentiate popular choices of the MOND $\\mut$ function by fitting a benchmark rotation curve, and argue that many of those functions are likely unphysical in the TeVeS context. We then propose a new free function for TeVeS in the domain relevant for galaxies, with a possible extension to solar system dynamics and cosmology.  \\subsection{Warming up to TeVeS} It is a tensor-vector-scalar Lorentz-covariant field theory, where the tensor is the Einstein metric $g_{\\alpha \\beta}$ out of which is built the usual Einstein-Hilbert action, $U_\\alpha$ is a dynamical normalized vector field, and $\\phi$ a dynamical scalar field. The action is the sum of the Einstein-Hilbert action for the tensor $g_{\\alpha \\beta}$, the matter action, the action of the vector field $U_\\alpha$, and the action of the scalar field $\\phi$. Einstein-like equations are obtained for each of these fields by varying the action w.r.t. each of them. The action of the scalar field $\\phi$ involves a dimensionless parameter $k$ (of the order of a few percents), a length scale parameter $l$ ($\\sim {\\sqrt{3k} c^2 \\over 4\\pi a_0}$), and a free dimensionless function linking $k l^2 |\\grad \\phi|^2 \\propto y$ with the auxiliary nondynamical scalar field $\\mu$.  More relevant to us, the physical metric in TeVeS near a quasi-static galaxy or the solar system is identical to that of General Relativity, with a potential \\beq\\label{twopart} \\Phi = \\Xi \\Phi_N + \\phi \\eeq where $\\Xi \\simeq 1$. This means that the scalar field $\\phi$ plays the role of the dark matter gravitational potential. It is related to the Newtonian potential $\\Phi_N$ (generated by the baryonic density $\\rho$) through the equation (similar to the field equation for the full $\\Phi$ in BM84) \\beq\\label{poisson} \\grad . [ \\mu_s \\grad \\phi] = \\grad^2 \\Phi_N = 4 \\pi G \\rho, \\eeq where $\\mu_s$ is a function of the scalar field strength $g_s=|\\grad \\phi|$.  It is related the $\\mu$ function of B04 and the interpolating function $\\mut$ of MOND by \\beq\\label{mus} \\mu_s \\equiv {\\mu \\over k'} = \\frac{\\mut}{1-\\mut}, ~k' \\equiv {k \\over 4 \\pi}. \\eeq In the intermediate to deep-MOND regime, the toy free function in the scalar action of B04 gives rise to the following interpolating function: \\begin{equation}\\label{btoy} \\mut(x) = {\\sqrt{1 + 4x} -1\\over\\sqrt{1 + 4x} +1}. \\end{equation} ",
        "conclusions": "Part of the amazing successes of the non-relativistic version of MOND in explaining galaxy dynamics is due to its ``standard\" interpolating function $\\mut = x/\\sqrt{1+x^2}$. When exploring a range of other empirical functions in the context of TeVeS, we see that the fit to rotation curves becomes poorer for more gradual $\\mut$ function, whilst an external field effect with imaginary dilation happens for more rapid changes of the $\\mut$ function. These two independent constraints from opposite sides suggest a fairly narrow range of TeVeS free functions. Among these we propose a simple expression which works for both very weak and very strong gravity (Eq.~\\ref{propmus}), with a possible extension to cosmology. Unlike the one in B04, the new function has the nice feature that it links quantities of TeVeS and of MOND easily, hence facilitates future examinations using galaxy dynamics and solar system data. The explicit simple monotonic interpolating function $\\mu_s(g_s/a_0)$ could be easily fed into a numerical solver for Eq.~(\\ref{poisson}) (e.g., as developed by Ciotti et al. 2005), and could allow for the modelling of realistic galaxy geometries (the curl-field of BM84, neglected here following conventional wisdom, could be put back with a realistic amount).  Combined with a galaxy with sensitive kinematical data, this may confirm or falsify our ``toy\" function, hence further establish or squeeze the parameter space of the TeVeS theory.  As two final remarks, (i) we note that multiple-imaged gravitational systems present a challenge to all MOND/TeVeS interpolating functions (cf. Fig.~\\ref{lens}). Zhao et al. (2005) found that elliptical galaxies of comparable luminosity and redshift show a large scatter in their Einstein ring sizes.  Among the previously proposed $\\mu$ functions the B04 toy function is the most effective in lensing, but none of the functions seems to fit all lenses in the point lens case (shaded zone); the fit is poorer with realistic lens mass profile (cf. Fig. 17 of Zhao et al. 2005). (ii) We also note that the dark matter potential is fundamentally different from the scalar field; although the two are sometimes degenerate in fitting rotation curves, there is no equivalent of EFE in dark matter, hence the dark matter potential enjoys more freedom."
    },
    "0512/astro-ph0512613_arXiv.txt": {
        "abstract": "We use radio, near-IR, optical, and X-ray observations to examine dynamic processes in the central region of the rich galaxy cluster Abell 2125. In addition to the central triple of E and  cD galaxies, including members of both major dynamical subsystems identified from a redshift survey, this region features a galaxy showing strong evidence for ongoing gas stripping during a high-velocity passage through the gas in the cluster core. The disk galaxy C153 exhibits a plume stretching toward the cluster center seen in soft X-rays by {\\it Chandra}, parts of which are also seen in [O II] emission and near-UV continuum light. {\\it HST} imaging shows a distorted disk, with star-forming knots asymmetrically distributed and remnant spiral structure possibly defined by dust lanes. The stars and ionized gas in its disk are kinematically decoupled, demonstrating that pressure stripping (possibly turbulent as well as ram) must be important, and that tidal disruption is not the only mechanism at work. Comparison of the gas properties seen in the X-ray and optical data on the plume highlight significant and poorly-known features of the history of stripped gas in the intracluster medium which could be clarified through further observations of this system. The nucleus of C153 also hosts an AGN, shown by the weak and distorted extended radio emission and a radio compact core. The unusual strength of the stripping signatures in this instance is likely related to the very high relative velocity of the galaxy with respect to the intracluster medium, during a cluster/cluster merger, and its passage very near the core of the cluster. An additional sign of recent dynamical events is the diffuse starlight asymmetrically placed about the central triple in a cD envelope. Transient and extreme dynamical events as seen in Abell 2125 may be important drivers of galaxy evolution in the cores of rich clusters. ",
        "introduction": "Progressively richer data in the optical and X-ray have revealed that clusters of galaxies, in addition to being important sites for galaxy evolution, are themselves evolving systems. This is seen in the evidence for cluster mergers, as shown by galaxy position/velocity arrays and substructure in the hot intracluster medium (ICM). Particularly intriguing has been evidence that such events on a cluster-wide scale might affect the individual galaxies, as manifested via star formation, occurrence of nuclear activity, and/or modifications due to the effects of tidal and ram-pressure forces. We have conducted an intensive study of Abell 2125, a cluster of richness class 4 at $z=0.247$, originally motivated by fractions of blue galaxies and radio detections which are exceptional for its redshift and richness. As discussed by \\citet{m03} and \\citet{o03b} (Paper II), A2125 appears to be a major merger in progress seen close to the line-of-sight. Spatially the two largest galaxy concentrations overlap in projection. Modeling suggests a position angle to the line-of-sight of about 30$^\\circ$ and that the two systems are within 0.2 Myr of core passage. The projected scale of the full A2125 concentration is about 5 Mpc, making the total extent about 10 Mpc, consistent with a major merger. Most of the radio-detected galaxies are distributed  uniformly in projection throughout this region. We argued in Paper II that the radio emission from most of these systems is due to star-formation activity and that these galaxies live in intermediate density environments, more conducive to mergers, interactions, and thus star-formation. The on-going major cluster merger may also be enhancing the galaxy-galaxy interactions in group-like environments we are seeing.  Additional aspects of the cluster merger appear to be active in the richest parts of A2125. The core region has four, fairly luminous radio sources ($> 10^{23}$ W Hz$^{-1}$). All the evidence which we will present here suggests that, unlike most rich clusters as we currently observe them, a complex interaction is taking place which is the result of the merger of two dense sub-clumps in the A2125 system. Most striking, one of these radio sources arises from a smaller galaxy in the process of losing much of its interstellar medium to the ram pressure of a rapid passage through the densest part of the ICM. We  discuss here the optical, radio, and X-ray evidence for this interpretation. ",
        "conclusions": "We have used optical, X-ray, and radio properties of galaxies near the core of the populous and apparently merging cluster Abell 2125 to probe the effects of cluster-scale events on individual galaxies. The radio source C153 shows spectacular features which suggest both an AGN and a starburst have been triggered during the interaction. A tail of emission, seen in X-rays, [OII] line emission and UV-light are also seen. A ``toy'' model suggests the event is consistent with triggering by tidal effects, and stripping of gas during the last 10$^8$ years, roughly the time it would take to cross the cluster core at an unusually high relative velocity.  Individual events of this intensity may be rare, requiring a gas-rich galaxy to cross deep within the cluster core at very high velocity. However, they do furnish an opportunity to test our ideas about how stripping proceeds, with less confusion from purely gravitational effects than is generally seen in more protracted instances. In particular, simultaneous detection of optical and X-ray gas in the plume from C153 in A2125 may show us the thermal history of gas as it moves from the interstellar medium to the hot intracluster medium. This event may also be showing us another way in which galaxy evolution can occur."
    },
    "0512/astro-ph0512339_arXiv.txt": {
        "abstract": "We present phase-resolved low resolution $JHK$ and higher resolution $K$-band spectroscopy of the polar VV Pup. All observations were obtained when VV Pup was in a low accretion state having a K magnitude near 15. The low resolution observations reveal cyclotron emission in the $J$ band during some phases, consistent with an origin near the active 30.5 MG pole on the white dwarf. The secondary in VV Pup appears to be a normal M7V star and we find that the $H$ and $K$ band fluxes are entirely due to this star at all orbital phases during the low accretion state. We use our higher resolution Keck spectroscopy to produce the first $K$-band radial velocity curve for VV Pup. Our orbital solution yields $K_2$=414$\\pm27$ km sec$^{-1}$ and leads to mass estimates of M$_1$=0.73$\\pm$0.05 M$_{\\odot}$ and M$_2$=0.10$\\pm$0.02 M$_{\\odot}$. We find that the mass accretion rates during the normal low states of the polars VV Pup, EF Eri, and EQ Cet are near 10$^{-13}$ M$_{\\odot}$ yr$^{-1}$. The fact that \\.M is not zero in low state polars indicates active secondary stars in these binary systems, including the sub-stellar donor star present in EF Eri. ",
        "introduction": "VV Puppis is considered to be a typical magnetic cataclysmic variable or ``polar''. It contains a highly magnetic white dwarf and a late-type secondary star. VV Pup has an orbital period of 100 minutes, and magnetic field strengths of 30.5 MG and 56 MG at the two white dwarf poles. As with all polars, VV Pup shows large brightness variations of a few magnitudes between what are termed high and low states. These are times of normal and very low mass transfer respectively for material accreted from the secondary star to the primary white dwarf. VV Pup has been studied in detail in the optical since the mid-1960's (Walker 1965; Liebert et al., 1978; Schneider \\& Young, 1980; Cowley et al., 1982; Wickramasinghe et al., 1989) and at high energy (X-ray, FUV, and UV: see Patterson et al., 1984; Vennes et al., 1995, Imamura et al., 2000; and Hoard et al., 2002). Nearly all studies have occurred during a high state when emission from the accretion stream, the accretion column near the white dwarf magnetic pole, and the accretion region at the pole itself dominate its flux output. One recent X-ray study of VV Pup (Pandel and Cordova 2005) was performed during a low state. At this time, it was observed that the mass transfer occurred in an irregular fashion and the X-ray flux from the accretion pole region varied by more than an order of magnitude during orbital phases when the active pole was in view. These authors concluded that irregular accretion is common in low-state polars, and that such flares are due to ``accretion rate fluctuations\" representing strongly varying mass transfer from the companion star. The flares are likely to result either from capture of coronal mass ejections from the secondary star by the white dwarf's strong magnetic field, or from solar flares near the L1 point. This indicates stellar activity in the secondary star in polars.  VV Pup has been in an extended low state much of the past three years. Based on our stochastic monitoring efforts, the current low state appears to have started near December 2002 and continued througout most of 2005 (until October or November) with a brief (2-3 weeks?) volley back to an apparent normal high state in December 2003. During these past few years, we have undertaken extensive photometric and spectroscopic studies of VV Pup in the optical and IR in order to confidently identify the secondary star and assess its properties and to study the non-accreting white dwarf surface Zeeman field and low state cyclotron emission. Results of our optical low state spectroscopy and photometry as well as our magnetic modeling effort are presented elsewhere (Mason, et al., 2005, 2006). Here we concentrate on IR spectroscopy and the secondary star. We present the first low state, phase-resolved $K$-band spectroscopy of VV Pup wherein we determine an orbital solution for the binary. We find the true K$_2$, the orbital phase of the secondary star at inferior conjunction, and the spectral type of the secondary star. We are also able to provide robust mass estimates for M$_1$ and M$_2$. ",
        "conclusions": "We present the first K-band, low state radial velocity solution for VV Pup. These observations allow, for the first time, a true orbital phasing for the binary to be determined. We note that during the low state, all the $H$ and $K$ band flux is from the secondary star and we can use the secondary star's spectral appearance and sodium absorption features in its photosphere to derive mass estimates for both components of the binary. The secondary star in this system seems to have a mass (M$_2$=0.10$\\pm$0.02 M$_{\\odot}$) and luminosity which are consistent with a normal, late-type M7V star. We determine a white dwarf mass of M$_1$=0.73$\\pm$0.05 M$_{\\odot}$. Normal secondary stars seem to be the rule for polars but the exception for the remainder of the cataclysmic variables (Harrison et al., 2005). We also provide an improved orbital ephemeris. Table 5 summarizes our determined parameters for VV Pup.  During this low state, the cyclotron harmonics in VV Pup were very weak, much weaker than in the LARP EQ Cet (Campbell \\& Harrison 2004), or EF Eri during its extended (8+ year) low state (Harrison et al., 2004). It seems that during low states, the accretion rates in polars can be just as low as those observed for LARPs ($\\sim$10$^{-13}$ M$_{\\odot}$ yr$^{-1}$), values in agreement with single, highly active M stars.. Our findings suggest that low state polar white dwarfs can remain at a low activity level by accreting material from the secondary star stellar wind and/or flare or star-spot activity. The tidally locked, spun-up secondaries in CVs should be enhanced in activity and capable of easily providing the needed mass loss levels. For VV Pup, low state observations indicate that a stellar wind and/or chromospheric activity is still on-going in a M7V star. For the secondary object in EF Eri, similar processes need to be present even though it is a cool, sub-stellar brown dwarf-like star."
    },
    "0512/astro-ph0512563_arXiv.txt": {
        "abstract": "{ We present a new \\element[][12]{CO}(J=1--0)--line survey of the Andromeda galaxy, M\\,31, with the highest resolution to date ($23\\arcsec$, or 85~pc along the major axis), observed {\\em On-the-Fly} with the IRAM 30-m telescope. We mapped an area of about $2\\degr\\times 0\\fdg 5$ which was tightly sampled on a grid of $9\\arcsec$ with a velocity resolution of $2.6\\kms$. The r.m.s. noise in the velocity-integrated map is around $0.35\\Kkms$ on the $T_\\mathrm{mb}$-scale.\\\\ Emission from the \\element[][12]{CO}(1--0) line is detected from galactocentric radius $R=3$\\,kpc to $R=16$\\,kpc, but peaks in intensity at $R\\sim 10$\\,kpc. Some clouds are visible beyond $R=16$\\,kpc, the farthest of them at $R=19.4$\\,kpc.  \\\\ The molecular gas traced by the (1--0) line is concentrated in narrow arm-like filaments, which often coincide with the dark dust lanes visible at optical wavelengths. The \\ion{H}{i} arms are broader and smoother than the molecular arms. Between $R=4$\\,kpc and $R=12$\\,kpc the brightest CO filaments and the darkest dust lanes define a two-armed spiral pattern that is well described by two logarithmic spirals with a constant pitch angle of 7\\degr--8\\degr. Except for some bridge-like structures between the arms, the interarm regions and the central bulge are free of emission at our sensitivity. The arm--interarm brightness ratio averaged over a length of 15~kpc along the western arms reaches about 20 compared to 4 for \\ion{H}{i} at an angular resolution of $45\\arcsec$.\\\\ In several selected regions we also observed the \\element[][12]{CO}(2--1)--line on a finer grid. Towards the bright CO emission in our survey we find normal ratios of the (2--1)--to--(1--0) line intensities which are consistent with optically thick lines and thermal excitation of CO.   \\\\ We compare the (velocity-integrated) intensity distribution of CO with those of \\ion{H}{i}, FIR at $175\\,\\mu$m and radio continuum, and interpret the CO data in terms of molecular gas column densities. For a constant conversion factor $X_\\mathrm{CO}$, the molecular fraction of the neutral gas is enhanced in the spiral arms and decreases radially from 0.6 on the inner arms to 0.3 on the arms at $R\\simeq 10$~kpc. We also compare the distributions of \\ion{H}{i}, H$_2$ and total gas with that of the cold (16\\,K) dust traced at $\\lambda175\\,\\mu$m. The ratios $N(\\ion{H}{i})/I_{175}$ and $(N(\\ion{H}{i})+2N(\\mathrm{H}_2))/I_{175}$ increase by a factor of $\\sim20$ between the centre and $R\\simeq 14\\kpc$, whereas the ratio $2N(\\mathrm{H}_2)/I_{175}$ only increases by a factor of 4. For a constant value of $X_\\mathrm{CO}$, this means that either the atomic and total gas--to--dust ratios increase by a factor of $\\sim20$ or that the dust becomes colder towards larger radii. A strong variation of $X_\\mathrm{CO}$ with radius seems unlikely. The observed gradients affect the cross-correlations between gas and dust. In the radial range $R=8$--14~kpc total gas and cold dust are well correlated; molecular gas is better correlated with cold dust than atomic gas. At smaller radii no significant correlations between gas and dust are found.\\\\ The mass of the molecular gas in M\\,31 within a radius of 18~kpc is $M (\\mbox{H}_2) = 3.6\\times 10^8\\, \\mbox{M}_{\\sun}$ at the adopted distance of 780~kpc. This is 12\\% of the total neutral gas mass within this radius and 7\\% of the total neutral gas mass in M\\,31. ",
        "introduction": " ",
        "conclusions": "The new \\element[][12]{CO}(J=1--0)--line survey of the Andromeda Galaxy presented here covers an area of about $2\\degr\\times 0\\fdg5$, which is fully sampled with a velocity resolution of $2.6\\kms$ and an angular resolution of $23\\arcsec$, the highest angular resolution to date of a map of this extent. At the adopted distance of 780\\,kpc the linear resolution is $\\rm 85\\,pc \\times 400\\,pc$ in the plane of M\\,31. The {\\em On-the-Fly} method of observing made it possible to measure nearly 1.7 million spectra (before gridding) in about 500 hours of effective observing time on M\\,31. The r.m.s. noise in the CO(1--0) line per 1~MHz channel is 25~mK in the northern half of the map and 33~mK in the southern half.  The velocity-integrated distribution, $I_{1-0}$, and the velocity field are shown in Fig.~\\ref{fig:co-map}. The molecular gas is concentrated in narrow, filamentary arms between 4 and 12~kpc from the centre with a maximum near 10~kpc. The inner arm at $R\\simeq 5$~kpc is remarkably bright. Only few clouds, often forming bridge-like structures, are detected in between the arms above $3\\times$ r.m.s. noise of typically $0.35\\Kkms$. The region within 2~kpc from the centre is almost free of molecular clouds brighter than the sensitivity of this survey.  The thin CO arms define a two-armed spiral pattern that can be well described by two logarithmic spirals with constant pitch angle of $7\\degr - 8\\degr$. At a resolution of $45\\arcsec$ the arm--interarm contrast reaches a maximum of 20 in $I_{1-0}$ compared to 4 in \\ion{H}{i} for the western bright arms. The \\ion{H}{i} arms are much wider than the molecular arms, and diffuse \\ion{H}{i} exists everywhere in between the arms and at large radii. Few molecular clouds are visible outside a radius of 16\\,kpc.  The velocity field of the molecular gas is very similar to that of the disk component in \\ion{H}{i}. At some positions perturbed velocity profiles occur that are possibly caused by nearby \\ion{H}{ii} regions.  Several selected regions were also observed in the $^{12}$CO(2--1) line. At a resolution of $23\\arcsec$ the line ratios are nearly constant with mean values of $I_{2-1}/I_{1-0} = 0.5-0.7$ in the arms. These line ratios are similar to those observed in other galaxies and show no indication of subthermal excitation.  Averaged radial profiles of the velocity-integrated CO and \\ion{H}{i} distributions show that for a constant conversion factor of $X_\\mathrm{CO} = 1.9\\times 10^{20}\\, \\mbox{mol\\, cm}^{-2}\\, (\\Kkms )^{-1}$ the molecular fraction of the neutral gas is enhanced on the spiral arms and decreases radially from about 0.6 on the inner arms to about 0.3 on the arms at $R\\simeq 10$~kpc (see Fig.~\\ref{fig:11}). Along the arms the molecular fraction also varies considerably. Comparisons with averaged radial profiles of the $\\lambda175\\,\\mu$m emission, which traces the cold (16\\,K) dust, revealed a strong, continuous increase of the apparent atomic gas--to--dust ratio of nearly a factor 20 between the centre and $R\\simeq 14$~kpc, whereas the apparent  molecular gas--to--dust ratio increases by about a factor of 4. The apparent  total gas--to--dust ratio also increases by about a factor of 20.  The strong apparent gradients in the molecular fraction and the gas--to--dust ratios influence the cross-correlations between CO, \\ion{H}{i} and FIR($175\\,\\mu$m) intensities.  In the radial range $R= 35\\arcmin - 60\\arcmin$ (about 8--14\\,kpc) the best correlation exists between total neutral gas and $175\\,\\mu$m, followed by that between CO and $175\\,\\mu$m and between \\ion{H}{i} and $175\\,\\mu$m. The relationships between \\ion{H}{i} and $175\\,\\mu$m and total gas and $175\\,\\mu$m are close to linear, whereas that between CO and $175\\,\\mu$m is a power law with exponent 1.6 due to a possible non-linearity between CO and \\ion{H}{i}. In the inner part of M\\,31, $R < 35\\arcmin$ (but outside the nuclear area), only total gas and FIR($175\\,\\mu$m) are reasonably well correlated (see Table~\\ref{tab:correl}).  The total molecular mass of M\\,31 within a radius of 18~kpc is $\\rm 3.6\\times 10^8\\, M_{\\sun}$, using the above-mentioned value of $X_\\mathrm{CO}$. As the total \\ion{H}{i} mass (without correction for opacity) is $\\rm 4.86\\times 10^9\\, M_{\\sun}$, the total mass of the neutral gas is $\\rm 5.2\\times 10^9\\, M_{\\sun}$. The total mass of the cold dust is $\\rm (1.3-3.8)\\times 10^9\\, M_{\\sun}$, hence the total gas--to--dust mass ratio is 410--137. The lower value is in agreement with the Galactic one.   The wealth of information contained in our new $^{12}$CO(1--0) survey of M\\,31 enables a number of new investigations into the physical relationships between molecular gas, atomic hydrogen gas, cold and warm dust, ionized gas, relativistic electrons and magnetic fields. Such studies will be the subject of forthcoming papers.  The data shown in Fig.~\\ref{fig:co-map} can be obtained from M.~Gu\\'elin ({\\tt guelin@iram.fr})."
    },
    "0512/astro-ph0512080_arXiv.txt": {
        "abstract": "We present the discovery of seven quasars at $z>5.7$, selected from $\\sim 2000$ deg$^2$ of multicolor imaging data of the Sloan Digital Sky Survey (SDSS). The new quasars have redshifts $z$ from 5.79  to 6.13. Five are selected as part of a complete flux-limited sample in the SDSS Northern Galactic Cap; two have larger photometric errors  and are not part of the complete sample. One of the new quasars, SDSS J1335+3533 ($z=5.93$), exhibits no emission lines; the 3-$\\sigma$ limit on the rest-frame equivalent width of Ly$\\alpha$+NV line is 5 \\AA. It is the highest redshift lineless quasar known, and could be a gravitational lensed galaxy, a BL Lac object or a new type of quasar. Two new $z>6$ quasars, SDSS 1250+3130 ($z=6.13$) and SDSS J1137+3549 ($z=6.01$), show deep Gunn-Peterson absorption gaps in Ly$\\alpha$. These gaps are  narrower the complete Gunn-Peterson absorption troughs observed among quasars at $z>6.2$ and do not have complete Ly$\\beta$ absorption. ",
        "introduction": "This paper is the fourth in a series presenting $i$-dropout $z\\gtrsim 5.7$ quasars selected from the multicolor imaging data of the Sloan Digital Sky Survey (SDSS; \\cite{York00}, Stoughton et al. 2002). In \\cite{z58} and in the first three papers of this series (\\cite{PaperI}, Paper I; \\cite{PaperII}, Paper II; \\cite{PaperIII}, Paper III), we presented the discovery of twelve luminous quasars at $z=5.74 - 6.42$, selected from $\\sim 4600$ deg$^2$ of SDSS imaging in the Northern Galactic Cap. In this paper, we describe the discovery of seven new quasars at $z=5.79 - 6.13$, selected from $\\sim 2000$ deg$^2$ of new SDSS imaging data. The scientific objectives, photometric data reduction, candidate selection and additional photometric and spectroscopic observation procedures are described in detail in Paper I and will not be repeated here. We present the photometric observations of the $i$-dropout candidates in the new area in \\S 2. The spectroscopic observations and the photometric and spectroscopic properties of the newly-discovered quasars are described in \\S 3. SDSS J133550.81+353315.8\\footnote{The IAU naming convention for SDSS sources is SDSS JHHMMSS.SS$\\pm$DDMMSS.S, and the positions are expressed in J2000.0 coordinates. The astrometry is accurate to better than $0.1''$ in each coordinate.} ($z=5.93$, SDSS J1335+3533 for brevity) is a quasar without detectable emission lines; we discuss the properties of this unusual object in \\S4.  Following the previous papers in this series, we use two cosmologies to present our results: (1) $\\rm H_0 = 50\\ km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\Lambda} = 0$ and $\\Omega_{M} = 1$ ($\\Omega$-model); (2) $\\rm H_0 = 71\\ km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\Lambda} = 0.73$ and $\\Omega_{M} = 0.27$ ($\\Lambda$-model, \\cite{Spergel03}). ",
        "conclusions": ""
    },
    "0512/astro-ph0512049_arXiv.txt": {
        "abstract": "We model the core helium flash in a low-mass red giant using Djehuty, a fully three-dimensional (3D) code. The 3D structures were generated from converged models obtained during the 1D evolutionary calculation of a 1$\\Msun$ star.  Independently of which starting point we adopted, we found that after some transient relaxation the 3D model settled down with a briskly convecting He-burning shell that was not very different from what the 1D model predicted. ",
        "introduction": "The core helium flash is an important event in the life of most stars with a zero-age mass between about 1 and 2 $\\Msun$; the minimum masses are a little lower for metal-poor stars. Since the work of Mestel (1952) and Schwarschild \\& H{\\\"a}rm (1962) it has been clear that such stars ignite helium in a thermonuclear runaway situation, the helium flash, because the helium core is electron-degenerate at the time of ignition. Empirically, it is clear that this runaway is (usually) not a catastrophic affair, like a supernova explosion, because a whole class of stars, the horizontal-branch stars of globular clusters, is well explained by the survival of helium-flash stars in a long-lived state of core helium burning (Faulkner 1966). Nevertheless, attempts to compute the evolution during the flash have a confusing history: some calculations (both 1D and 2D) have predicted a rather severe explosion, and others (both 1D and 2D) a relatively benign though rapid ignition.  Most calculations until fairly recently have been 1-dimensional (1D) simulations, in which turbulent convection has been treated by a spherical averaging process based largely on the mixing-length concept. Deupree (1996), who gives a nice summary of earlier work, performed some 2D (axially symmetric) simulations. He found that the 2D estimates were critically dependent on approximations made regarding eddy viscosity. We expect that by working in 3D we will not need to make such approximations, and we suggest that our results bear out this expectation.  The Djehuty project of the Lawrence Livermore National Laboratory is an effort to model stars in 3D. Our ultimate aim is to be able to model an entire star, up to and including the photosphere; and indeed to generalize this to binary stars, including gas flows between them in, for example, a Roche-lobe-overflow situation. We are approaching this goal, but it is fairly easy to see that a star like the Sun, for instance, would require at least $10^{12}$ nodes inside it if there is to be adequate resolution near the surface. As computer power continues to increase this will no doubt become possible, but for the present we limit ourselves to about $10^8$ nodes. We therefore content ourselves with a simulation of the He flash that includes only the He core and the radiative portion of the envelope; we ignore the deep surface convection zone.  Apart from the intrinsic interest of the He flash, our other reason for pursuing this particular problem is that it potentially is a very good test for the stability and accuracy of a hydrodynamic code. This is because we have, as mentioned above, a rather good reason to anticipate what the outcome should be. We do not expect it to become a violent supernova-like event. In following an explosion, it is not easy to look at the outcome and say `that is clearly what should have been expected', even if it is what we expected. But in a non-explosive situation it is not difficult to compare, for instance, the heat flux actually carried by turbulent convection across a spherical shell with the expectation from a simple mixing-length model.  In this paper we consider only non-rotating and non-magnetic cores, but we believe that both these processes could be important and we hope to address them in a later paper. In Section 2 we briefly outline the code; in Section 3 we describe (a) our 1D input models and (b) our 3D output models. In Section 4 we describe some issues and numerical tests regarding the stability of the calculation. ",
        "conclusions": "1) The 3D simulations were robustly stable, and, apart from the convective shell itself, the behavior of the star was consistent with hydrostatic modeling, even at the peak of the helium flash. Although exact spherical symmetry is obviously required in the 1D code, the 3D models seem to retain more-or-less spherical symmetry: there is no tendency for one hot spot to erupt and then dominate the shell, rendering it very asymmetric.  2) Convection is a critical element in determining the evolution through the helium flash.  In 1D hydrostatic modeling, convection is an approximation with effectively no information on the complex process by which hot spots develop and relax themselves.  3) In all of our models, the convection approached but never exceeded the outer boundary of convection as determined from a stability criterion in the 1D code.  However, we cannot claim that overshoot will not occur in longer runs.  In 1D, the inner boundary of the convective shell is nearly coincident with the peak energy producing shell. Our simulations do show a slight and potentially significant mixing below the convective shell from downward overshoot.  This leads to erosion of the non-burning central core, and if it continues could reduce or eliminate the mini-flashes that occurred in the 1D simulation. We intend to explore this further.  4) In the future, we intend to address all of the following: rotation, magnetic fields, and low metallicity. We shall also pursue further the possibility that some slow mixing outside the hydrogen-burning shell during First-Giant-Branch evolution might affect the surface abundances."
    },
    "0512/astro-ph0512238_arXiv.txt": {
        "abstract": "We consider geometries and possible physical models for weak low ionization absorbers based on the relative incidence of low and high ionization absorption systems.  To facilitate this, we present a survey of weak low ionization absorption systems ($W_r(2796){\\leq}0.3${\\AA}) in 35 quasar spectra from the archive of high resolution, ultra-violet {\\it HST}/STIS data.  When possible, we supplemented these spectra with Keck/HIRES and {\\it HST}/FOS data that cover more transitions over a larger range of wavelengths.  We found a total of 16 metal-line systems, with low and/or high ionization absorption detected.  It is known that the weak low ionization absorbers (which we probe with {\\MgIIdblt} or through the combination of {\\CII}$\\lambda$1335 and {\\SiII}$\\lambda$1260) trace an abundant population of metal-enriched regions (close to solar metallicity). Generally, models show that these systems have a $\\sim$10pc region of higher density gas and a $\\sim$1kpc region that represents a lower density phase of higher ionization absorption.  The goal of our survey was to compare absorption systems detected in low and/or high ionization gas (the latter traced with {\\CIVdblt} absorption).  We find the following: 1. All but 1 of the 10 weak low ionization systems have a related high ionization phase.  In 3 cases the high ionization gas has only a single component, kinematically centered on the low ionization absorption, and in the other 6 cases there are additional high ionization components offset in velocity.  There is one system, toward quasar 3C~273, without a high ionization cloud; 2. There are just 6 systems with only a high ionization phase as compared to the 9 systems with both low and high ionization phases; 3. The high ionization absorption in weak low ionization systems is, on average, stronger than in systems with only high ionization absorption; 4. The high ionization absorption in weak low ionization systems has similar kinematic structure to that in high ionization only systems.  We find that filamentary and sheetlike geometries are favored, due to the relatively small observed cross-section of high ionization only systems. Our statistical arguments suggest that, although low ionization absorbers are not closely associated with luminous galaxies, they arise in their immediate environments within the cosmic web. ",
        "introduction": "\\label{sec:intro}  Quasar absorption line spectroscopy is a powerful probe for the investigation of gas in the cosmic web.  It offers a high sensitivity that makes it possible to trace different gas phases over a large range of redshifts, exploring both the intergalactic medium and a variety of morphologies and evolutionary stages of galaxies.  The {\\MgII} resonant doublet has been extensively used to trace metals in galaxies at $0.3 < z < 2.2$, the redshift range for which {\\MgII} lies in the optical.  It is found that these low ionization systems, with $W_r(2796) {\\geq} 0.3$~{\\AA}, produce Lyman limit breaks \\citep{archiveII}.  The gas kinematics, evident in high resolution absorption profiles, are consistent with material in both the disks and extended halos of the host galaxies \\citep{kinmod,steidel02}.  Almost all of these ``strong'' {\\MgII} systems are found within $38h^{-1}$~kpc of $>0.05L_K^*$ galaxies \\citep{bb91,bergeron92,lebrun93,sdp94,s95,3c336}.  Furthermore, both observationally and theoretically, it is also expected that metal-rich absorbing gas can be found out to distances of at least $100h^{-1}$~kpc from luminous galaxies \\citep{chen, cirkovic}.  This gas can be detected through absorption in higher ionization species and/or as weak low ionization absorption.  Weak low ionization absorbers, with $0.02 {\\leq} W_r(2796) {\\leq} 0.3$~{\\AA}, constitute a significant fraction of the high {\\HI} column density regime ($15.8 < \\log N({\\HI}) < 16.8~{\\rm cm}^{-2}$) of the {\\Lya} forest at $z\\sim1$ \\citep{weak2}. With $dN/dz = 1.74\\pm0.10$ for $0.4 < z < 1.4$, they are twice as likely to be found along the quasar line of sight as strong low ionization absorbers \\citep{weak1}.  The majority of these systems arise in an optically thin, sub-Lyman limit environment \\citep{archiveII}.  Unlike strong absorbers, weak low ionization absorbers typically are not associated with bright galaxies (with $L{\\geq}0.05L_{\\star}$) to within a $50h^{-1}$~kpc impact parameter of the quasar line of sight (\\citet{weak1}; but see \\citet{churchill} for exceptions). This suggests that weak low ionization systems are a physically different population from strong absorbers, a result confirmed by the {\\MgII} equivalent width distribution measured from the Sloan Digital Sky Survey \\citep{nestor}.  Two thirds of weak low ionization absorbers have only a single narrow component with a Doppler $b$ parameter of just a few {\\kms} \\citep{weak1, weak2}.  Detection of {\\FeII} in some of these systems indicates a small ionization parameter and high density, which implies a small size, $<10$~pc \\citep{weak2}.  Based on photoionization models, at least some single-cloud weak low ionization absorbers without detected {\\FeII} are constrained to have similarly small sizes \\citep{weak1634}.  The absence of a Lyman limit break and the strength of the Lyman series lines provide constraints on the absorbers' metallicities.  Surprisingly, the derived metallicities are almost always greater than 10\\% solar, and are often as high as the solar value \\citep{weak2,weak1634}.  Those absorbers with a large iron to magnesium ratio must incorporate material from Type Ia supernovae, in which case the metals would be generated ``in-situ''.  Their large observed numbers along quasar lines of sight, in combination with the small derived sizes, suggests that single-cloud low ionization systems must be extremely abundant.  If they are assumed to have a spherical geometry there would be over a million such structures per bright galaxy. Most of the single-cloud weak low ionization absorbers have an associated higher ionization phase \\footnote{\\baselineskip = 0.5\\baselineskip A physically distinct region (or regions) that has a different density and/or temperature.}  that gives rise to {\\CIV} absorption \\citep{weak2}.  Photoionization modeling of three systems along the PG~$1634+706$ line of sight showed that low ionization absorption arises in a relatively high density region ($\\sim 0.01~{\\rm cm}^{-3}$) with a thickness of $0.1$--$100$~pc.  Higher ionization absorption arises in one or more lower density regions ($\\sim 10^{-3}~{\\rm cm}^{-3}$), with sizes on the order of a kiloparsec \\citep{weak1634}.  The other third of the weak low ionization absorbers have multiple, rather than single, low ionization clouds.  Modeling of one of these systems, toward PG~1634+704 led to a derived, low metallicity of $\\sim$3\\% solar in the low ionization gas \\citep{zonak}.  The kinematics of the low and high ionization gas suggested an origin in a pair of dwarf galaxies or in a symmetric wind from a starbursting dwarf \\citep{zonak}.  Also, \\citet{rosenberg} have suggested that local analogs to weak low ionization absorbers are related to winds from starbursting dwarfs.  In the case of the weak low ionization absorber at $z=0.0052$ toward 3C~273 \\citep{tripp3c273}, there is a post--starbursting dwarf galaxy, with a consistent redshift, at an impact parameter of $50 h^{-1}$~kpc \\citep{stocke}.  If winds from dwarfs typically produce multiple-cloud, weak low ionization absorption, then dwarf galaxies could be responsible for a large fraction of that absorber population.  On the other hand, there is evidence that some multiple-cloud weak low ionization absorbers are produced by lines of sight through the outskirts of luminous galaxies \\citep{masiero, ding05}.  This subclass would simply be an extension of the strong low ionization absorber population, which is directly produced by luminous galaxies \\citep{ding05}.  The single-cloud weak low ionization absorbers present a dilemma. They are not closely related to luminous galaxies like strong absorbers, yet they have high metallicities inconsistent with production in the traditional gas phase of dwarfs.  \\citet{weak2} suggested that they might be produced in Population III, pre-galactic star clusters or in supernovae remnants in very low luminosity galaxies, but these are only examples from possible classes of scenarios.  In this paper, we seek constraints on the nature of the physical structures that produce the low and high ionization phase absorption in single--cloud, weak low ionization absorption line systems.  {\\it Our strategy is to compare the relative numbers and kinematics of systems with detected absorption from both low and high ionization phases to those of systems without absorption detected from one or the other of those phases.}  For example, in principle, the low ionization absorption could be produced in a small spherical region embedded in a much larger spherical structure that produces {\\CIV} absorption.  In this case, we would expect that most systems have only high ionization absorption.  We use the archive of {\\it Hubble Space Telescope} ({\\it HST})/ Space Telescope Imaging Spectrograph (STIS) Echelle spectra ($R=30,000$ or $45,000$) \\citep{kimble} in order to make these comparisons, considering systems with coverage of {\\CIV} and at least some low ionization transitions ({\\SiII}$\\lambda$1260, {\\CII}$\\lambda$1335, and/or {\\MgIIdblt}).   In \\S~\\ref{sec:data}, we briefly describe the characteristics of our sample, and the search method. In \\S~\\ref{sec:systems_found} we describe the specific systems that we found.  In \\S~\\ref{sec:rez}, we present statistical results and Voigt profile fits. In the discussion of the paper, \\S~\\ref{sec:disc} and \\S~\\ref{sec:disc2}, we apply our results to simple thought experiments in order to constrain possible geometries of the systems, and finally to connect these models to physical interpretations. ",
        "conclusions": "\\label{sec:disc}  Considering the above results and the physical properties of weak low ionization absorption systems, derived from photoionization models, we will now discuss general constraints on their geometry.  Our goal is to better understand the nature of the structures in which the low and high ionization absorption arises.  Here we summarize briefly the important observational constraints that must be satisfied:  \\begin{enumerate}  \\item{There are just six high only systems found in our survey sample, which contains six single-cloud weak low ionization systems.  One additional system is found with single-cloud weak low ionization absorption, but with no detected high ionization absorption.  In comparison, three multiple cloud weak low ionization systems were found, each with detected high ionization absorption.}  \\item{For low + high ionization systems, the strongest {\\CIV} absorption is aligned with the strongest low ionization absorption.}  \\item{The {\\CIV} profiles for both the high only and low + high ionization systems have a similar distribution of kinematic spreads. For both categories, some systems have single components of high ionization absorption and some have multiple components.}  \\item{The {\\CIV} absorption tends to be similar, but somewhat weaker, for high only systems than for low + high ionization systems, both in the strongest {\\CIV} component and in outlying {\\CIV} components.}  \\item{The {\\Lya} absorption is very closely related to the kinematics of the {\\CIV} profiles, more so than it is to the low ionization absorption.}  \\end{enumerate}  Our discussion will focus on the origin of single-cloud weak low ionization systems with properties adopted from \\citet{weak2} and \\citet{weak1634}.  We choose this subset because many multiple-cloud weak low ionization systems may have a different physical origin, as described in \\S~\\ref{sec:intro}.  The structures that produce single-cloud weak low ionization absorption were found to have thicknesses ranging from $\\sim 1$~pc -- $100$~pc, and densities of $\\sim0.01$~{\\cc}.  The narrow line profiles imply a velocity dispersion of $\\sim 5$~{\\kms}, which corresponds to a virial mass of $\\sim 3 \\times 10^4$~{M$\\odot$} for a $10$~pc structure.  This is a considerably larger mass than the gas mass in a spherical structure with density $0.01$~{\\cc} and radius $10$~pc.  This suggests confinement by dark matter mini-halos or by a stellar structure such as a galaxy or star cluster.  The high ionization absorption related to the same systems is produced in larger structures, with larger velocity dispersions and with sizes of $\\sim 1$~kpc, which could be related to confinement.  In our survey, we found both low+high ionization systems and high only systems.  Point 4 above implies that there are real differences, either in degree or kind, between the low + high systems and the high only systems. There appear to be two different kinds of {\\CIV} structures, those with low ionization regions of higher density covering most of their area, and those without these higher density regions along the line of sight. In this way we can explain the high only systems and the weaker, offset components in low + high ionization systems as a different or less extreme population than the central component of a low + high ionization system.  In this population, eg., the total hydrogen column density might have been lower so that a higher density substructure did not collapse in the region.  The data would also be consistent with a {\\CIV} column density, for clustered clouds, that tends to decrease outward from a central cloud in which {\\MgII} absorption arises.  Without such a fall-off, a model could not explain the observed difference between the $N({\\CIV})$ of the component centered on the low ionization absorption and the offset {\\CIV} and high only system components.  Guided by these general principles, we now consider the most basic ``toy--model'' scenarios for the origins of the low and high ionization absorption in single-cloud weak low ionization absorbers:  \\begin{enumerate}  \\item{{\\it Single spherical low ionization region inside spherical high ionization region}--- The first, simplest model to consider, based on derived sizes, is one in which a $\\sim 10$~pc low ionization cloud is embedded in a $\\sim 1$~kpc higher ionization halo.  This model would give rise to $\\sim 10^4\\times$ more high only systems than single-cloud low + high systems.  We observe a ratio of only $6/6$, such a severe discrepancy that this model is easily ruled out.}  \\item{{\\it Multiple spherical low ionization regions inside spherical high ionization region}---  The previous model can be improved by using multiple spherical low ionization clouds inside a high ionization halo.  In order to produce the derived ratio of high only to low + high ionization systems, the covering factor should be $C_f \\sim 0.5$.  This would imply a probability of $\\sim C_f^2 = 0.25$ of passing through two of the low ionization clouds along a given line of sight.  This is not inconsistent with the data, since some fraction of the $\\sim 33$\\% of weak low ionization systems that have multiple clouds could be produced in this way.  However, in order to produce $C_f \\sim 0.5$, there would need to be $\\sim 5000$ low ionization structures with a size of $\\sim 10$~pc in each $\\sim 1$~kpc high ionization halo.  If the ratio of sizes of low to high ionization structures was instead $1/1000$, there would have to be a factor of $100\\times$ more small clouds per large halo.  The average spacing between the small clouds in the halo would be $94$~pc for the first choice of size ratio, and $20$~pc for the second.  This is not much larger than the cloud size, and seems a very contrived situation.  Without a physical model to support such a scenario, we consider this model unlikely, though it is not strictly ruled out.  Also, this simple model does not produce offset high ionization components as are found in our survey.  We can modify this model to attempt to explain the observed offset {\\CIV} components.  In this modified model, separate high ionization halos would give rise to these components.  These halos would have lower covering factors for low ionization absorption than the ``main'' halo (the one giving rise to low ionization absorption).  Since we are constrained to have an average of one offset component per low + high system, roughly half of the sky, looking out from the main halo, would need to be covered by the separate high ionization halos.  In this case, looking from a large distance at the clustered group of halos, an observer would see a covering factor for the separate high ionization halos that greatly exceeds that from the main halo.  Adding the separate high ionization halos would therefore produce more high only systems than we observe, so that in fact we have not improved on the model afterall.}  \\item{{\\it Low ionization shell surrounding high ionization shell}---  Alternatively, we now consider a situation where a low ionization shell surrounds an interior shell of high ionization gas.  In order not to exceed the number of multiple-cloud weak low ionization absorbers (relative to single-cloud), the low ionization shell should be fragmented.  This fragmented shell must cover roughly $33$\\% of the surface area of the bubble (assuming small fragments), so that it is more typical to pass through low ionization absorption on only one side of the interior region.  With only a single such ``bubble'', there is no clear explanation for offset high ionization components for this model. We might expect either one or two high ionization components, depending on the covering factor of the high ionization shell.  It is not clear that a component without detected low ionization absorption would have a smaller {\\CIV} column density than one which did, since it would come from a different layer. To account for the offset components, we can consider clustering of ``bubbles''.  However, such a model suffers from the same problems as our similar attempts to modify Model 2, in that the additional bubbles will produce too large a ratio of high only to low + high ionization systems.}  \\item{{\\it Network of filaments/sheets which gives rise to high ionization absorption, with embedded low ionization condensations}---  We finally consider a model in which low ionization condensations are embedded in filaments and/or sheets which give rise to high ionization absorption.  The overall covering factor of the low ionization regions (relative to high ionization regions) must be $C_f \\sim 0.5$ to explain the observed ratio of high only to low + high ionization systems.  In view of the small low ionization cloud thicknesses, a filamentary/sheetlike geometry is a straightforward way to produce the observed covering factor, without requiring huge numbers of separate spherical clouds. The lines of sight through filaments/sheets that do not give rise to low ionization absorption are typically characterized by smaller $N({\\CIV})$ values than those that do.  These smaller $N({\\CIV})$ structures would give rise both to high only systems and to the offset high ionization components of low + high systems.  In such a model, we would expect the number of high ionization components to be similar for the two categories of systems, as is observed.  We favor this model because it naturally produces all of the results derived from our survey.}  \\end{enumerate}"
    },
    "0512/astro-ph0512462_arXiv.txt": {
        "abstract": "We present stellar velocity dispersion (\\sigmastar) measurements for a significant sample of 40 broad-line (Type 1) active galaxies for use in testing the well-known relation between black hole mass and stellar velocity dispersion.  The objects are selected to contain Ca~{\\sc II} triplet, \\mgb\\ triplet, and Ca H+K stellar absorption features in their optical spectra so that we may use them to perform extensive tests of the systematic biases introduced by both template mismatch and contamination from the active galactic nucleus (AGN).  We use the Ca~{\\sc II} triplet as a benchmark to evaluate the utility of the other spectral regions in the presence of AGN contamination.  Broad \\feii\\ emission, extending from $\\sim 5050-5520$ \\AA, in combination with narrow coronal emission lines, can seriously bias \\sigmastar\\ measurements from the \\mgb\\ region, highlighting the need for extreme caution in its use.  However, we argue that at luminosities constituting a moderate fraction of the Eddington limit, when the \\feii\\ lines are both weak and smooth relative to the stellar lines, it is possible to derive meaningful measurements with careful selection of the fitting region.  In particular, to avoid the contamination of coronal lines, we advocate the use of the region 5250--5820 \\AA, which is rich in Fe absorption features.  At higher AGN contaminations, the \\chk\\ region may provide the only recourse for estimating \\sigmastar.  These features are notoriously unreliable, due to a strong dependence on spectral type, a steep local continuum, and large intrinsic broadening.  Indeed, we find a strong systematic trend in comparisons of \\chk\\ with other spectral regions.  Luckily the offset is well-described by a simple linear fit as a function of \\sigmastar, which enables us to remove the bias, and thus extract unbiased \\sigmastar\\ measurements from this region.  We lay the groundwork for an extensive comparison between black hole mass and bulge velocity dispersion in active galaxies, as described in a companion paper by Greene \\& Ho. ",
        "introduction": "It has long been known that some galaxies harbor supermassive black holes (BHs) at their centers, whose accretion-powered luminosity may outshine their entire host galaxy (e.g.,~Lynden-Bell 1969).  More recently, we have learned that most, if not all, galaxies with a bulge contain central BHs (Kormendy \\& Richstone 1995; Magorrian \\etal\\ 1998; Ho 1999), whose masses are tightly correlated with the stellar velocity dispersion (\\sigmastar) of the bulge (the \\msigma\\ relation: Ferrarese \\& Merritt 2000; Gebhardt \\etal\\ 2000a; Tremaine \\etal\\ 2002).  It is likely that the \\msigma\\ relation is established during the active galactic nucleus (AGN) phase of a galaxy's life-cycle, since energy emitted by the BH may simultaneously limit the gas supply for building both the bulge and the BH itself (e.g.,~Silk \\& Rees 1998).  If we are to understand this critical stage in galaxy evolution, we must have robust methods to estimate both BH masses and \\sigmastar\\ in active galaxies.  Reverberation mapping (Blandford \\& McKee 1982) is currently the most direct way to obtain BH masses in AGNs, although the masses are uncertain by a factor that depends on the poorly constrained geometry of the broad-line region (BLR).  Reverberation mapping provides an estimate of the size of the BLR from the lag between the variability in the photoionizing continuum and the broad emission lines.  With a BLR size at hand, one can infer a virial mass for the very central region of the AGN, and hence for the BH (Ho 1999; Wandel \\etal\\ 1999; Kaspi et al. 2000; Peterson et al. 2004): \\mbh\\ = $fR_{\\mathrm{BLR}} v^2/G$, where $v$ is the velocity dispersion of the BLR gas and the factor $f$ accounts for the geometry of the BLR (e.g.,~Onken \\etal\\ 2004; Kaspi \\etal\\ 2005).  For instance, a spherical BLR has an $f$ value of 0.75 when $v$ is measured using the full width at half-maximum (FWHM) of the broad line (e.g.,~Netzer 1990).  Reverberation mapping has further been used to derive an empirical relation between AGN luminosity and BLR radius (the radius-luminosity relation): \\rblr\\ $\\propto L_{5100 \\AA}^{0.6-0.7}$ (Kaspi \\etal\\ 2000, 2005; Greene \\& Ho 2005b). Since reverberation mapping data are time-consuming to obtain, and are currently available only for a small number of objects, practical applications of the virial technique to estimate BH masses have relied on the radius-luminosity relation and BLR line widths measured from single-epoch spectra (e.g.,~McLure \\& Dunlop 2001; Vestergaard 2002; Greene \\& Ho 2004).  Given the many uncertainties inherent in the virial technique (see Krolik 2001 for a detailed discussion), in particular the unknown geometry of the BLR, one may justifiably question its reliability.  On the other hand, Gebhardt \\etal\\ (2000b) and Ferrarese \\etal\\ (2001) showed that virial BH mass estimates agree surprisingly well with masses inferred from the \\msigma\\ relation, at least for a handful of objects.  Subsequent work (Nelson et al. 2004; Onken \\etal\\ 2004) has increased the sample size, still considering a total of only 17 objects between the two samples.  Onken \\etal\\ find that reverberation mapped masses, using the standard assumption of a spherical BLR and the FWHM, are $\\sim 0.26$ dex below the masses inferred from the Tremaine \\etal\\ (2002) fit to the \\msigma\\ relation; Nelson et al. obtain essentially the same result, concluding that the AGN sample lie systematically below the sample of inactive galaxies by 0.21 dex.  However, as noted by Nelson \\etal, the significance of the measured offset is low, given the large scatter (0.46 dex) in the data.  Since this value represents the zeropoint in the BH mass scale for AGNs, many more \\sigmastar\\ measurements are required to improve the calibration.  An enlarged sample would allow us to resolve outstanding technical questions, such as whether the FWHM or the actual second moment of the broad-line profile is a better measure of BLR velocity dispersion (e.g.,~Vestergaard 2002; Peterson \\etal\\ 2004), as yet unaddressed for virial masses obtained with single-epoch spectra.  Furthermore, we would be able to explore the importance of BLR geometry in a statistical way, both by addressing the possible importance of inclination (e.g.,~McLure \\& Dunlop 2001) and by obtaining a more secure empirical measure of $f$.  The BLR geometry may even depend on properties of the system such as BH mass and accretion rate.  With such calibrations in hand, we may begin to address other fundamental questions, such as whether the \\msigma\\ relation varies as a function of BH mass (e.g.,~Robertson \\etal\\ 2005) or evolves with time (e.g.,~Shields \\etal\\ 2003).  The primary hurdle in this endeavor lies with the difficulty of obtaining reliable \\sigmastar\\ measurements in active galaxies.  AGNs are bright, and thus detectable to cosmological distances, but the strong continuum of Type 1 sources dilutes the starlight while their rich emission-line spectrum confuses and distorts the shape of the stellar absorption features.  This paper presents a comprehensive discussion of how to tackle this problem.  Using a relatively large sample of AGNs for which we can measure \\sigmastar\\ reliably, we discuss the relative merits and complications of measuring \\sigmastar\\ from different spectral regions.  The sample is described in \\S2. Section 3 introduces the direct-fitting code we have developed for measuring \\sigmastar\\ and discusses the battery of tests to which we have subjected it to evaluate its robustness and limitations.   The many challenges inherent in dealing with AGN spectra are outlined in \\S4, where we devote considerable attention to finding the optimal spectral regions for measuring \\sigmastar\\ under realistic conditions encountered in AGNs.  We end with some practical suggestions to serve as a guide for other researchers (\\S5), followed by a summary (\\S6).  A companion paper (Greene \\& Ho 2005c) uses the final measurements obtained from this analysis to provide a new appraisal of the \\msigma\\ relation of AGNs.  Throughout we assume the following cosmological parameters to calculate distances: $H_0 = 100~h = 71$~\\kms~Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.27$, and $\\Omega_{\\Lambda} = 0.75$ (Spergel \\etal\\ 2003). ",
        "conclusions": "We present a statistically significant sample of 40 AGNs with measurable velocity dispersion (\\sigmastar) using regions surrounding the \\ion{Ca}{2} triplet, the \\mgb\\ triplet, and \\chk\\ lines, for the purposes of intercomparison.  Using a newly developed direct-fitting code, we perform a comprehensive set of simulations and  \\psfig{file=dilution_71.ps,width=0.5\\textwidth,keepaspectratio=true,angle=0} \\vskip -5mm \\figcaption[]{ Schematic diagram to delineate the regimes in which the different spectral regions can be most effectively used to measure \\sigmastar. The solid curves mark the three key redshift intervals ($z$ = 0.05, 0.76, and 1.3) to which an $m_i=19.1$ mag AGN will be spectroscopically targeted by the SDSS, and the dashed curves identify, from bottom to top, \\lledd\\ = 0.01, 0.1, 0.5, and 1.  The fraction of light coming from the bulge was estimated using the \\mbh-$L_{\\mathrm{bulge}}$ relation of Marconi \\& Hunt (2003). \\label{dil}} \\vskip 5mm cross-checks to evaluate the merits and limitations of using each spectral region, with the aim of obtaining realistic estimates of the many systematic uncertainties that affect measurements of \\sigmastar\\ in AGNs.  We argue that the \\ion{Ca}{2} triplet is least susceptible to template mismatch and AGN contamination from emission lines, and so provides the most reliable measurements of \\sigmastar.  We therefore use these lines as a benchmark to test the other spectral regions.  We examine two types of AGN contamination: featureless continuum dilution that effectively lowers the S/N of the absorption features, and emission lines, both narrow and broad, that fill in the stellar absorption lines and bias the line profiles in subtle ways.  In terms of dilution, we find that \\sigmastar\\ is measurable for AGN fractions \\lax 71\\%, \\lax 85\\%, and \\lax 90\\% using \\ion{Ca}{2} triplet, \\mgb, and \\chk, respectively.  As for AGN emission-line contamination, we find that measurements around \\mgb\\ are very sensitive both to broad \\feii\\ emission and narrow emission from [\\ion{Fe}{6}], [\\ion{Fe}{7}], and [\\ion{N}{1}].  The narrow lines can be avoided by using the spectral region 5250--5820 \\AA.  Broad \\feii\\ contamination can bias the fits severely when the \\feii\\ width is narrowest.  For a given \\mbh\\ this will occur at the highest luminosity, and thus is worst when the object radiates at a large fraction of its Eddington luminosity.  In such cases, and for the highest levels of AGN dilution, the \\chk\\ spectral region (and particularly the Ca~K line) is the best remaining option.  While there is a systematic offset between $\\sigma$(Ca~K) and \\sigmastar\\ derived from other spectral regions, in Type 2 AGNs it is well fit by a linear relation that allows us to derive unbiased \\sigmastar\\ estimates from Ca~K. Further work is required to ensure that this correction is generally applicable.  As a by-product of this study, we have derived reliable velocity dispersions for 35 Type 1 AGNs, and set out reasonable selection criteria to generate a much larger sample of AGNs with a broader range of $z$, black hole mass, and accretion rates.  Such a sample would enable us to examine the \\msigma\\ relation for active galaxies with true statistical power.  We will be able to calibrate virial masses for black holes with unprecedented accuracy, as well as search for second-order trends in the \\msigma\\ relation with mass, accretion rate, and redshift.  Greene \\& Ho (2005c) present a preliminary investigation of the \\msigma\\ relation using the sample analyzed in this study."
    },
    "0512/astro-ph0512148_arXiv.txt": {
        "abstract": "The physics of the intracluster medium, in particular the values for the thermal conductivity and the viscosity are largely unknown and subject to an ongoing debate. Here, we study the effect of viscosity on the thermal state of the intracluster medium using three-dimensional cosmological simulations of structure formation. It is shown that viscosity, provided it is not too far off from the unmagnetised Spitzer value, has a significant effect on cluster profiles. In particular, it aids in heating the cool cores of clusters. The central cooling time of the most massive clusters in our simulation is increased by more than an order of magnitude. In large clusters, viscous heating may help to establish an entropy floor and to prevent a cooling catastrophe. ",
        "introduction": "Cooling by bremsstrahlung and line emission leads to a loss of pressure support in the centers of galaxy clusters, which, in the absence of non-gravitational heating, would cause a slow, subsonic inflow of gas towards the center of the gravitational well.  As a result, cluster cores should cool and accrete gas at rates of hundreds and more solar masses per year.  This scenario is in conflict with observational evidence that indicates that mass deposition rates are consistently lower than predicted. Moreover, the gas temperatures in cluster centers are maintained typically above $\\sim 2$ keV \\citep{peterson:01}.\\\\  More evidence for non-gravitational heating in clusters comes from cluster scaling relations. These relations show departures from self-similarity: In the absence of non-gravitational heating and radiative cooling, the entropy is expected to scale linearly with the mean cluster temperature. However, observations by Ponman, Sanderson \\& Finoguenov (2003), Pratt \\& Arnaud (2005) and Piffaretti et al. (2005) indicate a scaling of entropy roughly according to $T^{2/3}$. Moreover, they reveal a systematic excess of entropy in low-mass clusters (e.g., Ponman, Sanderson \\& Finoguenov 2003). \\cite{churazov:01,bruggen:02, bruggen:02a} have argued that heating by a central AGN can keep the ICM from cooling dramatically in the center.\\\\  Thermal conduction has been put forward to explain the absence of soft X-rays from galaxy clusters (\\cite{kim:03} and references therein). Cosmological simulations with thermal conduction have been performed by \\cite{dolag:04} and \\cite{jubelgas:04}. If thermal conduction is at work, other transport processes such as viscosity are bound to be important, too. Thermal conduction transports energy and is mediated mainly by the faster electrons. Viscosity, on the other hand, transports momentum and is mediated primarily by the more massive ions. It is not clear that the suppression factors of both transport processes should be the same. Whether they are the same may depend on the scale magnetic fluctuations extend to. This scale may be much larger than the gyroradii of electrons and ions (in which case suppression factors could be comparable) or it could be comparable to the ion gyroradius. The magnitude of the suppression factor is motivated by various theoretical arguments (e.g., given in \\citealt{narayan:01}). However, we note that the precise value of the suppression factor is highly uncertain and, depending on the nature of magnetic turbulence, may even exceed the Spitzer value \\citep{cho:03} or be supressed well below it.\\\\  Based on observations of the Perseus cluster, it has been suggested by \\cite{fabian:03} that viscosity may play an important role in dissipating energy injected by the central AGN. The case for this is based on the existence of long, straight H$\\alpha$-filaments that appear to rule out the presence of strong turbulence in the cores of galaxy clusters. The Reynolds number for a fluid flow whose viscosity is suppressed with a factor, $f$, with respect to the Spitzer value, is given by  \\begin{equation} {\\rm Re}\\sim 50\\ f_{0.3}^{-1}\\ n_{\\rm e,-3}\\ L_{\\rm kpc}\\ v_{500}\\ T_7^{-5/2} \\ , \\label{eq:1} \\end{equation} where $n_{\\rm e,-3}$ is the electron number density in units of $10^{-3}$ cm$^{-3}$, $L_{\\rm kpc}$ the typical size of an eddy in units of kpc, $v_{500}$ the associated velocity in units of 500 km s$^{-1}$ and $T_7$ the temperature of the fluid in $10^7$ K. As such values of Reynolds numbers are below the critical value that separates laminar and turbulent regimes, it was concluded that viscosity can play an appreciable role in the ICM, provided that viscosity is not heavily suppressed.\\\\  Subsequently, heating by viscous dissipation of AGN-induced motions has been simulated by \\cite{ruszkowski:04,ruszkowski:04b,bruggen:05}. It was concluded that, provided viscosity is not suppressed significantly with respect to its unmagnetised value, viscous dissipation of AGN-induced motions can balance the radiative losses in the ICM. In \\cite{reynolds:05} the effect of viscosity on the evolution of radio bubbles was studied, and it was found that viscosity had a stabilising effect on underdense bubbles. \\cite{fujita:04} have studied the dissipation of motions induced by acoustic-gravity waves in cluster cores. They find that, provided the wave amplitude is large enough, they can suppress the radiative cooling of the cores. \\cite{kim:05} have investigated the heating of clusters by dynamical friction. They concluded that friction can be an important supplier of heat but is unlikely to prevent the onset of cooling flows.\\\\  Meanwhile, cosmological simulations of galaxy clusters with, both, particle and grid-based methods have reached a fairly mature state. In particular, the addition of increasingly sophisticated recipes for radiative losses, star formation and stellar feedback have led to cluster models that can produce many observed features of galaxy clusters \\citep{loken:02}.\\\\  \\cite{motl:04} have simulated the formation of cool cores and observed that any ``cooling flow'' is overwhelmed by the velocity field inside the cluster, which has speeds of up to 2000 km s$^{-1}$. Nonetheless, such violent motions did not prevent the formation of cool cores. This shows that the formation of cool cores is inevitable unless some source of heating is present. Full 3D cluster collision simulations by \\cite{ritchie:02} and \\cite{ricker:01} show that mergers can disrupt cooling flows. It is a common feature of all cosmological simulations of galaxy clusters that the ICM shows a substantial velocity field with many motions being supersonic. The ICM shows a complex dynamics with cool fronts, filaments, shocks etc. New observations confirm this picture. Detailed observations of unprecedented resolution by the latest X-ray observatories have revealed a rich portfolio of substructure in galaxy clusters. For example, \\cite{schuecker:01} find substructure in the majority of clusters in their REFLEX+BCS cluster sample. The ubiquity of substructure points to a high frequency of mergers and other events that prevent the cluster from relaxing to a smooth, unperturbed state.\\\\  Both, observations and simulations suggest that the ICM is in violent motion and that relaxed clusters, in which the gas sits almost statically in its potential well, are very rare, if they exist at all. In the presence of viscosity, a fraction of the kinetic energy in these motions can be dissipated to heat the ICM. Thus, even in inactive phases of a central AGN, there can be heating in the form of viscous dissipation.  In this paper, we study the effect of viscosity on the intracluster medium in a cosmological simulation. In particular, we wish to compute to what extent viscous dissipation of random motions in a cluster can contribute to the heating of cluster cores. In the next section, we describe the technique and setup of our simulations. Finally, the results are discussed in Sec.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We produced two simulations: one without viscosity and one with a third of Spitzer viscosity. Both runs included radiative cooling and star formation. A density projection of the viscous simulation is shown in Fig.~\\ref{fig1}. The largest cluster is situated near the upper edge of this figure. \\\\  Halos were identified using the HOP-algorithm developed by \\cite{eisenstein:97}. The parameters of the most massive cluster in our simulation are summarised in table 1. The virial radius is calculated for an overdensity of $\\delta \\rho/\\rho =200$, and the virial mass is the total mass (dark matter + baryons) within the virial radius. $L_X$ is the total X-ray luminosity within $R_{\\rm vir}$ in the band from 0.1 - 2.4 keV. Finally, $\\mathrm{M_{stars}}$ is the total mass converted to stars. It is striking that, while the gross features of the cluster, such as virial radius, mass and temperature, are nearly identical between the runs, the X-ray luminosities differ substantially. The X-ray luminosity depends on the density squared and is thus dominated by the central portion of the clusters. As discussed in the next section, there are pronounced differences in the central densities, which do not affect the total mass, though. Also, the mass converted to stars differ between the two runs, with the viscous run producing less stars.\\\\  \\begin{deluxetable}{| l || c | c | c} \\tablewidth{0pt} \\tablenum{1} \\tablecolumns{3} \\tablecaption{Properties of the most massive cluster in our sample\\label{tab:clusters}} \\startdata \\hline & without visc &  with visc      \\\\ \\hline \\hline $\\mathrm{R_{virial}}$ & $1.11 \\: \\mathrm{Mpc}$ & $1.12 \\: \\mathrm{Mpc}$ \\\\ $\\mathrm{M_{virial}}$ & $2.2 \\times 10^{14} \\: \\mathrm{M_{\\odot}}$ & $2.3 \\times 10^{14} \\: \\mathrm{M_{\\odot}}$ \\\\ $\\mathrm{L_{x}}$ & $1.6 \\times 10^{46} \\: \\mathrm{erg} \\, \\mathrm{s^{-1}}$ & $2.2 \\times 10^{45} \\: \\mathrm{erg} \\, \\mathrm{s^{-1}}$ \\\\ $\\mathrm{T_{\\rm vir}}$ & $3.0 \\times 10^{7} \\: \\mathrm{K} $ & $3.1 \\times 10^{7} \\: \\mathrm{K} $ \\\\ $\\mathrm{M_{stars}}$ & $7.4\\times 10^{12} \\: \\mathrm{M_{\\odot}}$ & $6.5\\times 10^{12} \\: \\mathrm{M_{\\odot}}$ \\enddata \\end{deluxetable}      In Fig.~\\ref{fig2} - \\ref{fig5} we show the mass-weighted temperature, density, cooling time and entropy, respectively, as a function of radius in the most massive, non-merging cluster in our sample.  While a cool core forms in the run without viscosity, it is absent in the viscous run. As can be seen from the temperature profile, Fig.~\\ref{fig2}, viscosity has essentially removed the cool core and the mass-weighted temperature even rises slightly in the centre. The density profiles also show significant differences (see Fig.~\\ref{fig3}).  In the presence of viscosity, the density is nearly constant in the core, whereas, in the non-viscous run, it rises sharply in the inner 40 kpc.  With 1/3 Spitzer viscosity, the density is very flat over the inner hundred kiloparsecs. In summary, we find that the runs with viscous dissipation lead to a hotter and less dense core. Consequently, the cooling time increases in the center, with respect to the non-viscous runs (see Fig.~\\ref{fig4}). The central cooling time is about two orders of magnitude higher than in the non-viscous case and larger than the Hubble time. If viscous heating was this efficient, no other sources of heating would be required to prevent a cooling catastrophe. The entropy, which is shown in Fig.~\\ref{fig5}, displays a very extended floor and has no central dip as in the run without viscosity. A comparison of a statistically relevant sample of simulated clusters can be compared with observed samples, such as the one by \\cite{donahue:05}. Thus, one may be able to constrain the viscosity of the ICM from X-ray observations.\\\\  In Fig.~\\ref{fig6} we compare the X-ray luminosity from the cluster in the band from 0.1 - 2.4 keV. In the case with viscosity, it is apparent that the cluster is more extended and that it lacks a strong emission spike in the center.\\\\  The temperature-dependence of viscosity biases this mode of heating towards hotter, more massive clusters, and will, thus, affect cluster scaling relations. We find that the effect of viscosity becomes systematically less important with decreasing mass of the cluster. In Fig.~\\ref{fig7} - Fig.~\\ref{fig9}, we show the corresponding profiles of a smaller cluster with a mass of $9.7\\times 10^{13}\\ M_{\\odot}$. Evidently, viscosity only has a minor effect on the central density profile. A proper study of this effect requires a bigger simulation box that contains a large number of clusters, and is the subject of future work.\\\\  Note that in the simulations presented here, heat conduction has been neglected. As argued above, heat conduction and viscosity are intrinsically linked, and it would be interesting to study their joint effect on the cluster. Since the suppression factors of viscosity and conductivity are unknown, the consideration of heat conduction introduces another, essentially free, parameter into the problem. For this first simulation, we decided to study the isolated effect of viscosity only. \\\\  In order to validate our code modules, we have repeated a low-level test simulation in a small box with the FLASH code. The viscosity routines in FLASH have been tested in \\cite{ruszkowski:04,ruszkowski:04b}. Starting from similar initial conditions and including the same physics, we obtain very similar results between Enzo and FLASH.   As a result of numerical diffusion, the effective Reynolds number will be finite, even in the zero-viscosity case. The effective Reynolds numbers attainable in the simulation are proportional to the number of grid points across the fluctuation of interest to the power $n$, where $n=3$ is the order of the numerical scheme\\footnote{See, e.g., \\citet{bowers:91} for the definition of ``the order of the numerical scheme'', as it is different from the customary definition of accuracy of a perturbative calculation.}  \\citep{porter:94}. As expressed in Eq.~\\ref{eq:1}, the Reynolds number in the ICM is of the order of 50, assuming 0.3 of the Spitzer value. Test simulations with FLASH and Enzo suggest that the effective Reynolds number for the resolution chosen here is $>$ 1000. Both runs that we have presented here have identical spatial resolution. Hence, the differences between the runs are solely due to the physical viscosity.\\\\  Both, radiative losses and viscosity, are sensitive to the spatial resolution of the computational grid. We experimented with different refinement criteria and compared results from runs with different effective resolutions. Generally, the effect of physical viscosity becomes larger as the resolution increases. The same is true for radiative losses. The results presented here seem to be reasonably converged and differ only marginally from the run with a refinement level less.\\\\  Assuming hydrostatic equilibrium when inferring cluster masses from synthetic observations, leads to masses that are systematically lower (by 10-15 per cent) than the actual cluster masses in simulations. This affects cluster constraints on cosmological parameters (see e.g., \\citet{allen:04}). If the ICM is viscous, it is conceivable that the assumption of hydrostatic equilibrium is better met. Thus, gas viscosity may reduce systematic deviations of cluster masses inferred from X-ray observations from their true values, which, in turn, may have consequences for precision measurements of cosmological parameters."
    },
    "0512/astro-ph0512304_arXiv.txt": {
        "abstract": "We present the stellar and gas kinematics of a sample of 18 nearby late-type spiral galaxies (Hubble types ranging from Sb to Sd), observed with the integral-field spectrograph {\\tt SAURON} at the 4.2-m William Herschel Telescope. {\\tt SAURON} covers the spectral range 4800-5380 \\AA, allowing us to measure the H$\\beta$, Fe, Mg{\\textit{b}} absorption features and the emission in the H$\\beta$ line and the [OIII]$\\lambda\\lambda$4959,5007\\AA\\, and [NI]$\\lambda\\lambda$5198,5200\\AA\\, doublets over a 33$''$$\\times$ 41$''$ field of view. The maps cover the nuclear region of these late-type galaxies and  in all cases include the entire bulge. In many cases the stellar kinematics suggests the presence of a cold inner region, as visible from a central drop in the stellar velocity dispersion. The ionised gas is almost ubiquitous and behaves in a complicated fashion: the gas velocity fields often display more features than the stellar ones, including wiggles in the zero-velocity lines, irregular distributions, ring-like structures. The line ratio [OIII]/H$\\beta$ often takes on low values over most of the field, probably indicating a wide-spread star formation. ",
        "introduction": "From a theoretical point of view, we have a well-defined paradigm for the formation of disc galaxies within the Cold Dark Matter (CDM) hierarchical structure formation scenario (\\citealt{fall}, \\citealt{silk}): discs quietly settle and cool inside dark matter haloes, while bulges form through mergers of multiple haloes. However, some of the observed properties of spiral galaxies suggest a larger complexity in their formation history. The presence of bulges is not ubiquitous and their nature can be ambiguous. Evidence has accumulated in the past years showing that many bulges have a disc-like, sometimes exponential radial fall-off of the stellar density (\\citealt{andredakis}, Andredakis, Peletier \\& Balcells 1995, \\citealt{jong}, Courteau, de  Jong \\& Broeils 1996, Carollo \\& Stiavelli 1998, \\citealt{seigaretal}, MacArthur, Courteau \\& Holtzmann 2003). Numerical simulations seem to suggest that the dissolution of bars inside the discs may trigger the formation of three-dimensional stellar structures with roughly exponential profiles (\\citealt{pfenniger}, \\citealt{combes}, \\citealt{raha}, Norman, Sellwood, Hasan 1996); this could mean that some bulges form through the evolution of dynamical instabilities in the disc. Quite recently, the quality of imaging data made available through HST boosted the study of the inner regions of spiral galaxies, showing that they can host a variety of structures: bulges, nuclear star clusters, stellar discs, small bars, double bars, star-forming rings (\\citealt{marcella97}, \\citealt{marcella98b}, Carollo, Stiavelli \\& Mack 1998, \\citealt{marcella99}, \\citealt{perez}, \\citealt{marcella02}, \\citealt{boker}, \\citealt{laine}, Falc\\'on-Barroso et al. 2005, \\citealt{emma}), without there being an agreement about their origin and evolutionary pattern. Ongoing large projects like the panchromatic SINGS survey \\citep{sings} which makes use of observations at infrared, visible and ultraviolet wavelengths represent a very useful approach to building a comprehensive picture of galactic structure, but at the moment rely mostly on imaging. Looking at disc galaxies from a spectroscopic perspective would add kinematic information and insight into stellar populations which cannot come from imaging, and could help us tracing their star formation and mass assembly histories.\\\\ \\indent Contrary to the massive spheroids, the stellar populations and kinematics of late-type disc-dominated galaxies are poorly known, due to the difficulty of reliably measuring and interpreting such diagnostics in low surface brightness environments which are so full of dust, star formation and substructures: not much attention has been paid to the spectroscopic counterpart of all the mentioned imaging that has been carried out. There are a few exceptions to this statement: \\citet{boker01} started a project on STIS long-slit spectroscopy of 77 nearby late-type spiral galaxies previously imaged with HST/WFPC2; first results are discussed in \\citet{boker03}; \\citet{walcher} analysed UV slit-spectroscopy of the nuclei of nine late-type spirals; these studies are mainly focussed on the nature of the innermost components, in particular on the nuclear star clusters.\\\\ \\indent We are currently engaged in a study aimed at investigating the properties of the nuclear regions of very late-type galaxies. In such environments, long-slit spectra are too limited to be useful for modelling and interpretation and have generally been used only to discuss the properties of emission-lines (see for example \\citealt{gallag}, who measure the position-velocity curve of 21 extreme late-type spiral galaxies using the H$\\alpha$ emission-line). Here we present deep integral-field spectroscopy that not only makes it easier to study the kinematics and physical properties of stars and gas, but also allows to study and model the stellar populations.\\\\ \\indent We were granted 6 nights at the William Herschel Telescope (WHT) of the Observatorio del Roque de los Muchachos in La Palma, Spain, to obtain two-dimensional spectroscopy with the integral-field spectrograph {\\tt SAURON}, which was custom-built for a representative census of elliptical and lenticular galaxies, and Sa bulges (the so-called {\\tt SAURON} survey, see Bacon et al.\\ 2001, de Zeeuw et al.\\ 2002, hereafter, respectively, Paper I, Paper II). The present work can be regarded as a natural extension of the {\\tt SAURON} survey towards the end of the Hubble sequence. Our purpose was to use {\\tt SAURON} in order to map the stellar and gaseous (H$\\beta$, [OIII], [NI]) kinematics and the absorption line-strength distributions of the indices H$\\beta$, Mg, Fe, in the region 4800-5380\\AA. In this paper we present the observations and data reduction and the resulting kinematical maps for 18 Sb-Sd galaxies. The data and maps will be made available via the {\\tt SAURON} website (http://www.strw.leidenuniv.nl/sauron\\,).\\\\ \\indent The paper is structured as follows. Section \\ref{samplesec} describes the sample selection and characteristics; Section \\ref{observationsec} summarizes the observations and data reduction; Section \\ref{methodsec} describes the methods applied to calculate the stellar and gaseous kinematics from our spectra; Section \\ref{comparison} carries out a comparison with previous measurements; Section \\ref{mapssec} presents and discuss the kinematical maps and looks in particular at the behaviour of the stellar velocity dispersion. Finally, Section \\ref{conclusionsec} summarizes the results. Detailed modelling and interpretation of the data will come in future papers. ",
        "conclusions": "Two-dimensional kinematics and stellar population analysis of spiral galaxies toward the end of the Hubble sequence (Hubble types later than Sb) is still a relatively unexplored field: late-type spirals are very complex objects, often faint and full of substructures, as recently proved by analysis of HST images. They have been the target of a few photometric and long-slit optical spectroscopic observations, but measurements of their two-dimensional kinematics were still missing. We have started a project on a sample of 18 such objects using integral-field spectroscopic observations obtained with {\\tt SAURON}. This allowed us to measure the stellar kinematics, the flux and kinematics of the H$\\beta$ 4861\\AA\\, and [OIII]$\\lambda\\lambda$4959,5007\\AA\\, emission-lines and the strength of the H$\\beta$, Fe and Mg{\\textit{b}} absorption features over a two-dimensional area covering the central region of our galaxies.\\\\ \\indent In this paper we discussed the first results from this study, presenting the two-dimensional kinematics for stars and ionised gas. The majority of our galaxies is shown to be kinematically cold and to possess a considerable amount of ionised gas, covering in most cases a large part of the {\\tt SAURON} FoV and frequently following bar or spiral arm patterns in the spatial distribution. A quite common feature of our measured stellar kinematical maps is a central depression in the velocity dispersion, which assumes very often low values; we measured the velocity dispersion profiles and correlated their slopes with the morphological type: later-type galaxies tend to have velocity dispersion profiles which increase outwards. This implies small bulge/disc ratios and the presence of inner, occasionally star-forming, disc-like structures. We also qualitatively compared the characteristics of our maps with the galaxy's properties known from literature HST isophotal analysis: the main conclusion common to spectroscopy and photometry is that the kinematic detection of a cold inner region turns out to be often related to the lack of a classical stellar bulge and the presence of small-scale structures (nuclear star clusters, inner rings, inner bars). The gaseous component turns out to be almost ubiquitous and kinematically highly complex, displaying in many cases irregular velocity fields, with the kinematic axis twisting or bending or wiggling, or even without a clear sense of rotation, possibly because of the dust which strongly affects these objects. They also host intense star formation, often spread over the whole region we have observed, as suggested by the low values in the [OIII]/H$\\beta$ line ratio maps.\\\\ \\indent In follow-up papers we will model the observed kinematic fields in detail, present the line-strength maps for these same galaxies, consider the bulge-disc decomposition, and compare our results with those for the 24 Sa bulges in the {\\tt SAURON} survey."
    },
    "0512/physics0512049_arXiv.txt": {
        "abstract": "We have in earlier work (Basse N P 2005 {\\it Phys. Lett. A} {\\bf 340} 456) reported on intriguing similarities between density fluctuation power versus wavenumber on small (mm) and large (Mpc) scales.  In this paper we expand upon our previous studies of small and large scale measurements made in fusion plasmas and using cosmological data, respectively.  The measurements are compared to predictions from classical fluid turbulence theory. Both small and large scale data can be fitted to a functional form that is consistent with the dissipation range of turbulence.  The comparable dependency of density fluctuation power on wavenumber in fusion plasmas and cosmology might indicate a common origin of these fluctuations. ",
        "introduction": "\\label{sec:intro}  Transport of particles and energy across the confining magnetic field of fusion devices is anomalous \\cite{wootton}, i.e., it is much larger than the neoclassical transport level associated with binary collisions in a toroidal geometry \\cite{hinton}. It is thought that anomalous transport is caused by plasma turbulence, which in turn manifests itself as fluctuations in most plasma parameters. To understand anomalous transport, a two-pronged approach is being applied: (i) sophisticated diagnostics measure fluctuations and (ii) advanced simulations are being developed and compared to these measurements. Once our understanding of the relationship between fluctuation-induced anomalous transport and plasma confinement quality is more complete, we will be able to reduce transport due to the identified mechanism(s).  The fusion plasma measurements presented in this paper are of fluctuations in the electron density. Small-angle collective scattering \\cite{saffman,basse1} was used in the Wendelstein 7-AS (W7-AS) stellarator \\cite{renner} and phase-contrast imaging (PCI) \\cite{mazurenko} is being used in the Alcator C-Mod tokamak \\cite{hutch}.  We specifically study density fluctuation power versus wavenumber (also known as the wavenumber spectrum) in W7-AS and C-Mod. These wavenumber spectra characterize the nonlinear interaction between turbulent modes having different length scales.  The second part of our measurements, wavenumber spectra (i) of galaxies from the Sloan Digital Sky Survey (SDSS) \\cite{sdss} and (ii) from a variety of sources (including the SDSS data) are published in Ref. \\cite{tegmark1} and have been made available to us \\cite{tegmark2}.  The paper is organized as follows: In section \\ref{sec:rev} we review our initial results from Ref. \\cite{basse2}. Thereafter we analyze our expanded data set in section \\ref{sec:add} and in response to the results revise our treatment of the original W7-AS measurements in section \\ref{sec:anal}. A discussion follows in section \\ref{sec:disc} and we conclude in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have in this paper reported on suggestive similarities between density fluctuation power versus wavenumber on small (mm) and large (Mpc) scales.  The small scale measurements were made in fusion plasmas and compared to predictions from turbulence theory. The data sets fit Eq. (\\ref{eq:exp_pow_decay}), which has a functional form that can be explained as dissipation by turbulence theory. The large scale cosmological measurements can also be described by Eq. (\\ref{eq:exp_pow_decay}). In general, two wavenumber ranges separated by a transitional region are identified.  The similar dependency of density fluctuation power on wavenumber might indicate a common origin of these fluctuations, perhaps from fluctuations in QGPs at early stages in the formation of the universe. The value of $\\alpha$ is almost identical for both fusion plasma and cosmological measurements at wavenumbers close to but above the peak of the spectra.  To progress further, it is essential that the quantity of wavenumber-resolved fusion plasma turbulence measurements is vastly increased.  \\ack This work was supported at MIT by the Department of Energy, Cooperative Grant No. DE-FC02-99ER54512. We thank M Tegmark for providing all cosmological measurements analyzed in this paper.  \\newpage"
    },
    "0512/astro-ph0512418_arXiv.txt": {
        "abstract": "We examine the black hole mass - galaxy bulge relationship in high-redshift QSOs.  Black hole masses are derived from the broad emission lines, and the host galaxy stellar velocity dispersion \\sigstar\\ is estimated from the widths of the radio CO emission lines.  At redshifts $z > 3$, the CO line widths are narrower than expected for the black hole mass, indicating that these giant black holes reside in undersized bulges by an order of magnitude or more.  The largest black holes ($\\mbh > 10^9$~\\msun) evidently grow rapidly in the early universe without commensurate growth of their host galaxies.  CO linewidths offer a unique opportunity to study AGN host galaxy dynamics at high redshift. ",
        "introduction": "The evolution of galaxies and their central supermassive black holes (SMBH) is a major topic of current interest (see review by Combes 2005).  The tight correlation of the black hole mass \\mbh\\ and the luminosity and the velocity dispersion \\sigstar\\ of the host galaxy's bulge (Gebhardt et al. 2000a; Ferrarese \\& Merritt 2000; review by Kormendy \\& Gebhardt 2001) points to a close evolutionary relationship.   This has inspired a number of theoretical investigations involving growth of the black hole by accretion of gas. When the luminosity of the active galactic nucleus (AGN) becomes large enough, it drives the remaining gas from the nucleus of the host galaxy, ending black hole growth as well as star formation (Silk \\& Rees 1998;  Di Matteo et al. 2005, and references therein).  Numerical simulations by Di Matteo et al. find a tight \\mbhsigstar\\ relationship from simple assumptions about heating by the AGN luminosity leading to evacuation of residual gas.  The \\mbhsigstar\\ relationship for nearby galaxies  is given by Tremaine et al. (2002) as \\begin{equation} \\label{e:tremaine} \\mbh = (10^{8.13}~\\msun)(\\sigstar/200~\\kmps)^{4.02}, \\end{equation} with black holes spanning the range roughly $10^5-10^9$ \\msun.  One clue to the origin of this relationship is its evolution over cosmic time. Shields et al. (2003, hereinafter S03) investigated the \\mbhsigstar\\ relationship in QSOs at high redshift by estimating \\mbh\\ and \\sigstar\\ from the \\hbeta\\ and \\oiii\\ emission lines, respectively.  They found that the local \\mbhsigstar\\ relationship is obeyed, in the mean, up to masses approaching $10^{10}~\\msun$ and at redshifts up to $z = 3.3$. Salviander et al.  (2005) similarly find little evolution in the \\mbhsigstar\\ relationship in QSOs at redshifts up to $z = 1.1$.    However, the observation of luminous quasars up to redshifts $z = 6.4$ (Fan et al. 2001) shows that large black holes ($\\sim 10^9~\\msun$) can grow rapidly in the early universe (Haiman \\& Loeb 2001; Volonteri \\& Rees 2005).   This raises the question of whether, even at such early times, these massive black holes reside in commensurate galaxies so as to obey the local \\mbhsigstar\\ relationship.  The use of \\oiii\\ as a surrogate for \\sigstar\\ (see below) has a practical limit of $z \\approx 3$, as the \\oiii\\ lines are shifted beyond the infrared K-band window for higher redshifts.  However, the radio CO emission lines have been observed in a number of high redshift QSOs and radio galaxies (Solomon \\& Vanden Bout 2005, hereinafter SV05).  To the extent that the CO line widths reflect orbital motion in the gravitational potential of the host galaxy, this affords an opportunity to assess the \\mbhsigstar\\ relationship at a cosmic time of only one billion years.  In this paper, we investigate the relationship between \\mbh\\ and CO linewidth for high redshift QSOs as a means of assessing the \\mbhsigstar\\ relationship at early times.  All values of luminosity used in this study are calculated using the cosmological parameters $\\hnot = 70~\\kmpspmpc, \\Omega_{\\rm M} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "These results suggest that, at least for the very large black holes involved in the highest redshift QSOs, rapid black hole growth has occurred at $z > 4$ without the formation of a proportionally massive host galaxy.    Taking into account the possible underestimation of \\sigco\\ with respect to \\sigstar\\ (see section 2.3), \\deltalogmbh\\ should be renormalized down by $\\sim0.5$. This gives $\\deltalogmbh \\approx 0.5$ and 1.5  for groups 1 and 2,  respectively, much larger than any systematic deviations in \\mbhsigstar\\ relationship for local galaxies.  Two uncertainties in the interpretation of the CO line widths involve mergers and disk orientation.  The CO profiles often show double peaks that could be interpreted as mergers or the characteristic profile of a rotating disk.  However, for a merger, the individual components will have \\wco\\ even narrower than the observed profile.  One of the components presumably corresponds to the observed black hole, and \\sigco\\ for that component alone will be smaller than the \\sigco\\ we have used, based on the full line profile.  This makes the mismatch of \\mbh\\ and \\sigco\\ even worse.  For a single disk model, face-on orientations might be favored for the high redshift QSOs, perhaps to a greater degree than occurs for the PG quasars discussed above.   However, the CO maps are frequently elongated (see discussion of SDSS J1148 above), and the frequent double peaked profiles might not be observed for closely face-on orientations if there is a substantial turbulent component to the orbital motion. Moreover, we see an order-of-magnitude increase in \\deltalogmbh\\ between groups 1 and 2, despite abundant molecular gas  in both groups.  If we go so far as to renormalize the CO widths to give $\\deltalogmbh\\  = 0$ for group 1, this still leaves group 2 with $\\deltalogmbh \\approx 1$.   The CO results here differ from the findings of SO3 for objects in the redshift range $z = 1$ to 3 using \\sigthree.  S03 find that, on average, QSOs at  $z = 1$ to 3 obey  equation \\ref{e:tremaine}. The black hole masses are similar to those of our high redshift CO QSOs.   The \\oiii\\ emission comes from the narrow line region (NLR) whose radius in such luminous QSOs should be similar to that of the molecular gas in our CO QSOs.  There is substantial scatter;  Bonning et al. (2005) find an rms scatter of 0.13 dex in the use of \\sigthree\\ as a surrogate for \\sigstar, corresponding to 0.5 in \\deltalogmbh. This seems inadequate to explain the discrepancy between the CO and [O~III] results, since a number of objects are involved in each sample.  Salviander et al. (2005) discuss \\sigthree\\ and \\mbh\\ for five radio-quiet QSOs from Sulentic et al. (2004).  Salviander et al. measured  [O~III] widths from the Sulentic et al. spectra, which have better spectral resolution than the high redshift data used by S03. These objects, with redshifts of 0.8 to 2.4, have a mean \\deltalogmbh\\ of 0.5, somewhat larger than in S03.  On the other hand, Bonning et al. (2005) find a mean \\deltalogmbh\\ of -0.4 for 6 radio quiet QSOs in the redshift range 2 to 3 from Shemmer et al. (2004) and Netzer et al. (2004).  Overall, these results are consistent with the findings of S03. The apparent discrepancy between the CO and \\oiii\\ results underscores the need for a better understanding of both the CO and \\oiii\\ line widths in the most luminous QSOs.  Our results for high redshift QSOs contrast with the results of Borys et al. (2005), who compare the stellar mass in submillimeter galaxies (SMG's) at $z \\approx 2$ with the central black hole masses.  Using population synthesis models to interpret the infrared luminosity of the SMG's, Borys et al. derive a typical stellar mass $M_* \\approx 10^{11.4}~\\msun$.  The black hole masses are derived from the X-ray luminosity on the assumption that the bolometric AGN luminosity is close to the Eddington limit (Alexander et al. 2005).  The black hole masses are $\\sim50$ times {\\em smaller} than expected for the galaxy mass and the typical ratio of $\\mbh/\\mbulge$ for nearby inactive galaxies.  Borys et al. argue that their small black holes are consistent with a scenario in which a gas-rich merger triggers a massive starburst and builds a large bulge quickly.  An optical QSO is only visible for a brief final phase when the black hole has grown large, supporting a large AGN luminosity that dispels the residual gas and reveals the nucleus.  Our CO QSOs  violate the canonical \\mbhsigstar\\ relationship in the opposite sense of having black holes too large for their host galaxies.  Do the giant black holes observed at high redshift acquire commensurate host galaxies through later mergers, gas accretion, and star formation? The local galaxy luminosity function does not afford a sufficient number of commensurately large galaxies to host the largest black holes observed in QSOs (Netzer 2003; Shields \\& Gebhardt 2004).   Some of these black holes evidently remain to this day in comparatively modest galaxies.  CO emission lines are currently the only practical means of testing the applicability of the \\mbhsigstar\\ relation to galaxies with redshifts larger than $z ~\\sim3$.  The present capability for such studies is limited.  The number of high-z galaxies with detected CO emission is small and not likely to grow dramatically with present telescopes.  Furthermore, the known examples are a flux limited sample, mostly aided by gravitational lensing.  New facilities such as the Atacama Large Millimeter Array (ALMA) will allow larger, more statistically significant and unbiased studies (Carilli 2005)."
    },
    "0512/astro-ph0512132_arXiv.txt": {
        "abstract": "We have observed ten red giant stars in four old Large Magellanic Cloud globular clusters with the high-resolution spectrograph MIKE on the Magellan Landon Clay 6.5-m telescope. The stars in our sample have up to 20 elemental abundance determinations for the $\\alpha$-, iron-peak, and neutron-capture element groups.  We have also derived abundances for the light odd-Z elements Na and Al.  We find NGC 2005 and NGC 2019 to be more metal-rich than previous estimates from the Ca\\thinspace{\\sc ii} triplet, and we derive [Fe/H] values closer to those obtained from the slope of the red giant branch.  However, we confirm previous determinations for Hodge 11 and NGC 1898 to within 0.2~dex. The LMC cluster \\lbrack Mg/Fe\\rbrack{} and \\lbrack Si/Fe\\rbrack{} ratios are comparable to the values observed in old Galactic globular cluster stars, as are the abundances [Y/Fe], [Ba/Fe], and [Eu/Fe].  The LMC clusters do not share the low-Y behavior observed in some dwarf spheroidal galaxies. \\lbrack Ca/Fe\\rbrack, \\lbrack Ti/Fe\\rbrack, and \\lbrack V/Fe\\rbrack{} in the LMC, however, are {\\em significantly} lower than what is seen in the Galactic globular cluster system. Neither does the behavior of \\lbrack Cu/Fe\\rbrack{} as a function of [Fe/H] in our LMC clusters match the trend seen in the Galaxy, staying instead at a constant value of $\\sim$$-$0.8. Because not all \\lbrack$\\alpha$/Fe\\rbrack{} ratios are suppressed, these abundance ratios cannot be attributed solely to the injection of Type Ia SNe material, and instead reflect the differences in  star formation history of the LMC vs.\\ the Milky Way. An extensive numerical experimental study was performed, varying both input parameters and stellar atmosphere models, to verify that the unusual abundance ratios derived in this study are not the result of the adopted atomic parameters, stellar atmospheres or stellar parameters. We conclude that many of the abundances in the LMC globular clusters we observed are distinct from those observed in the Milky Way, and these differences are intrinsic to the stars in those systems. ",
        "introduction": "Globular clusters have been key to gaining insights into the early epoch of formation and evolution for galaxies in general and for the Galaxy in particular. Because of the proximity of Galactic globular clusters (GGC) to us, we can obtain color-magnitude diagrams and high-resolution spectra of individual stars, which has allowed us to measure ages and abundances with unique accuracy. These data show a complex and interesting picture for the GGC, including a dispersion in abundance ratios, trends in ratios with kinematics, and the possibility of the capture of clusters from other galaxies. We can now observe clusters in other galaxies of the Local Group with the same techniques to compare their cluster systems with the GGC and determine the variation in globular cluster systems from galaxy to galaxy and the possible contributions of other galaxies to the Milky Way system.  The Magellanic Clouds, less distant than some GGCs, provide an excellent opportunity to observe abundance patterns in another globular cluster system in detail. The Large Magellanic Cloud (LMC) has long been known to harbour clusters of similar age, mass and metallicity to the GGCs (Searle, Wilkinson, \\& Bagnuolo 1980).  Testa \\etal\\ (1995) and Brocato \\etal\\ (1996) provided the first ages based on main-sequence turnoff measurements of the oldest clusters in the LMC.  The main-sequence turnoffs in a large number of old clusters in the LMC have subsequently been observed with the Hubble Space Telescope (HST). Some clusters in the LMC are coeval with nearby GGCs, such as M5, M4, and M92 (Olsen \\etal{} 1998, LMC-O98, hereafter; Johnson \\etal{} 1999, LMC-J99, hereafter).  The ages, kinematics, metallicities and abundance ratios of the GGCs have provided much insight into the formation of the Galaxy. Searle \\& Zinn (1978) argued that the outer halo clusters were younger than the inner halo clusters and that implied that a slow, chaotic buildup of the outer parts of the Galaxy had occurred. That mergers have contributed to the formation of the Galaxy was clearly shown with the discovery that the Sagittarius dwarf spheroidal galaxy (dSph) is currently being subsumed by the Milky Way (Ibata \\etal\\ 1994). The positions of GGCs on great circle orbits (Buonanno \\etal{} 1994) that sometimes include other satellites of the Milky Way (Fusi Pecci \\etal\\ 1995) hint at past accretion events. Lin \\& Richer (1992) argued that the positions and radial velocities of Rup 106 and possibly Pal 12 suggest that they had been tidally captured from the Magellanic Clouds (MC). By incorporating proper motions in the analysis, Dinescu \\etal\\ (2000) suggested that it was more likely that Sagittarius was the original host galaxy of Pal 12. Bellazzini, Ferraro, \\& Ibata (2003) extended the analysis to conclude that at least four outer halo GGCs belonged to Sagittarius, in addition to the four clusters whose positions lie near the main body of Sagittarius (Ibata \\etal\\ 1995)  Information about the history of the Galaxy is also contained in the chemical abundance ratios of old stars.  In a seminal paper, Tinsley (1979) argued that enhanced \\afe{}-ratios{\\footnote{We adopt the usual spectroscopic notation that [A/B] $\\equiv$ {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\star}$ -- {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\odot}$, and log~$\\epsilon$(A) $\\equiv$ {\\rm log}$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm H}$) + 12.0, for elements A and B.  Also, in this paper, except for instances where [m/H] or Z are specifically stated, we define metallicity as the stellar [Fe/H].}  in metal-poor stars were a consequence of the different timescales for the production of the $\\alpha$-elements (e.g., O, Ne, Mg, Si, Ca, and sometimes Ti) in core-collapse supernovae (Type II SNe) vs.\\ the Fe produced by both SNe Type Ia and Type II.  SNe Type II progenitors are short-lived (1Myr--100Myr) massive stars whereas progenitors of SNe Type Ia (mass-exchange binary systems including a white dwarf star) require longer to evolve and do not contribute to the chemical evolution of the Galaxy until $\\ge 1$Gyr subsequent to the formation of the binary system (Timmes, Woosley, \\& Weaver 1995; Matteucci \\& Recchi 2001).  Hence, the ratio of Type Ia/Type II SNe events determines the \\afe{}.  Systems that have recently started forming stars and have only had the contributions from massive stars to the interstellar medium would then be predicted to possess low Type Ia/II ratios and \\afe{}\\ $>0$. Old GCCs possess high \\afe{} ratios (Pilachowski, Sneden, \\& Wallerstein 1983; and references therein), in accord with the idea that GGCs were among the first surviving Galactic objects to have formed.  Nissen \\& Schuster (1997; hereafter, NS97) discovered a sample of moderately metal-poor field stars with low [$\\alpha$/Fe] ratios and suggested that they could have accreted from dwarf galaxies with a chemical evolution history different than that of the solar neighborhood, allowing the material out of which they formed to include the ejecta from Type Ia SNe while they were still relatively metal-poor. Also exhibiting low \\afe{} ratios with respect to the general GGC population are Rup 106 and Pal 12 (Brown, Wallerstein, \\& Zucker 1997), which are 2-3 Gyr younger than other GGCs (Buonanno \\etal\\ 1990; Stetson \\etal\\ 1989). Brown \\etal\\ interpreted the solar \\afe{}-ratios to be the result of the cluster being formed long enough after star formation had begun in the surrounding region to have its abundance ratios substantially affected by contributions of iron from Type Ia SNe.  The characteristics of the production sites of other elements may also provide insight into timescales, initial mass functions, and other properties of clusters. At this time, both the predictions from theory and the data from globular clusters are murkier than the case of [$\\alpha$/Fe]. For example, the site of the rapid neutron-capture process ($r$-process) is uncertain, but it is clear that the r-processed material appears before the slow-neutron-capture process ($s$-process) in asymptotic giant branch stars begins to contribute much to Galactic chemical evolution (Truran 1981). The r-process produces some heavy elements, such as Eu, more readily than others, such as Ba and La, while the s-process does the opposite. Therefore, ratios such as [Ba/Fe] and [Ba/Eu] contain information about when clusters formed from a chemical evolution standpoint. Other element ratios are also sensitive to the mix of stars that polluted the ISM. First, the metallicity of the progenitor of a Type II SNe is important in the synthesis of such elements as Na, Al, and Cu (e.g., Arnett 1971, Woosley \\& Weaver 1995). Second, the mass of the SN affects the ratio of the $\\alpha$ elements produced (e.g. Woosley \\& Weaver 1995; McWilliam 1997). Less massive stars make lower ratios of [Mg/Ca] and [Mg/Si], for example. [Si/Ti] should be highest for a 20 M$_{\\odot}$ star according to the Woosley \\& Weaver yields.  Recent efforts to measure many elements in globular clusters have shown that other abundance ratios, such as the ones listed above, vary between GGCs. Ivans et al.\\ (1999; 2001; hereafter M4-I99 and M5-I01), measured abundance ratios of 14 elements in 36 giants in each of M4 and M5, two GGCs with similar [Fe/H] and ages. Within either cluster, the stars possess comparable abundance ratios for elements not sensitive to proton-capture nucleosynthesis, but the same is not true of a comparison between clusters. Confirming and expanding upon the earlier results of Brown \\& Wallerstein (1992), M4-I99 found that the mean [Si/Fe] ratio for the M4 stars is 3-$\\sigma$ greater than in M5 stars.  The abundances of [Al/Fe], [Ba/Fe], and [La/Fe] are also significantly higher in M4 stars. Interestingly, these clusters also differ in their apogalactic distances, with apogalactocentric radii of 5.9 and 35.4 kpc for M4 and M5, respectively (Dinescu et al. 1999). This same apparent trend with apogalactic distance and some abundance ratios may also be reflected in halo field and cluster stars (NS97; Hanson \\etal\\ 1998; Stephens 1999; Fulbright 2002; Lee \\& Carney 2002; Fulbright 2004).  However, employing an extensive sample from the literature which included the results incorporated in Stephens (1999) and Fulbright (2004), Venn \\etal{} (2004) argue that the low [$\\alpha$/Fe] ratios in halo stars are, if anything, correlated with extreme retrograde orbits and statistically not correlated with apogalactic distance.  Additional cluster-to-cluster abundance variations have been found in other studies.  Pursuing the investigation of the \\afe\\ trends to the inner halo, Lee \\& Carney (2002) report high [Si/Fe] and low [Ti/Fe] in NGC 6287, NGC6293, and NGC6541, three metal-poor GCC (--1.8 $\\le$ [Fe/H] $\\le$ --2.0). Lee, Carney, \\& Habgood (2004) report high [Si/Fe] and low [Ti/Fe] in M68 stars as well. Pal 12 stars, in addition to low \\afe{} compared with other GGCs, also possess subsolar values of [Na/Fe] and [Ni/Fe] (Brown et al. 1997; Cohen 2004). Ter 7 stars show low [Ni/Fe] (Tautvai{\\v s}ien{\\.e} et al.\\ 2004).  As these examples illustrate, the evidence for abundance variations between GGCs has been firmly established  Within an individual GGC, star-to-star variations are observed among the light elements sensitive to proton-capture nucleosynthesis (e.g., C, N, O, Na, Mg, and Al).  Early detections of CN variations among giant stars by Lindblad (1922) and Popper (1947) were expanded to higher resolution  (see e.g., Osborn 1971; Peterson 1980) and, for some GGCs, to stars on the main sequence (see Hesser 1978; Hesser \\& Bell 1980).  Also observed in globular cluster populations (but absent in the field star population) are anti-correlations in the abundances of [O/Fe] with [Na/Fe] and [Al/Fe] (see Gratton, Sneden, \\& Carretta 2004 for a recent review).  One possible explanation for the light-element abundance patterns is ``deep'' mixing in red giants (e.g., Sweigart \\& Mengel 1979; Denissenkov \\& Weiss 1996, Denissenkov \\& VandenBerg 2003), dredging up the products of proton-capture nucleosynthesis from the interior out to the photosphere.  It remains unclear, however, how the temperatures of the interiors of the red giant stars can even get hot enough to convert Mg to Al (Langer, Hoffman, \\& Zaidins 1997; Messenger \\& Lattanzio 2002).  For some time, it had been thought that intermediate-mass asymptotic giant branch (AGB) stars could be responsible for producing the abundance patterns (Cottrell \\& Da Costa 1981), but recent work by Denissenkov \\& Herwig (2003) and Fenner \\etal\\ (2004) show that the observed abundance correlations are not replicated in model yields of AGB stars.  The presence of the abundance correlations at or below the main-sequence turnoff suggests that the variations may be primordial or the result of pollution by more evolved stars.  It is likely that some combination of effects are at work.  In M13, for example, the abundance patterns are also correlated with the evolutionary state of the stars (Kraft \\etal\\ 1993; Sneden \\etal\\ 2004; Johnson \\etal\\ 2005). While all clusters that have been examined for deep mixing effects show the associated abundance anomalies, some clusters appear to be more affected than others. The classic example is M3 and M13 (Kraft \\etal\\ 1992), where stars in M13 have [O/Fe] down to values of $-$0.87, while the most oxygen-poor stars M3 stars have [O/Fe] = $-0.25$. Other abundance ratios of light elements sensitive to proton-capture nucleosynthesis are similarly extreme in M13 but not in M3 (see e.g., Johnson \\etal\\ 2005 and references therein).  To summarize the situation in the GCC system, the majority of clusters exhibit super-solar \\afe{} ratios, abundance ratio trends with apogalactic distance or prograde/retrograde orbits, solar iron-peak element ratios, and evidence of abundance correlations in the light elements, possibly due to deep mixing. {\\it Are these universal properties of old globular cluster systems, or do they represent the unique history of the Milky Way?}   Low-dispersion spectra of individual giants by Cowley \\& Hartwick (1982) were used to measure spectral indices for nine old LMC clusters, including Hodge 11.  Subsequently, Olszewski \\etal\\ (1991, LMC-O91) performed a comprehensive study to measure the metallicities of the LMC clusters using low-dispersion spectra of the \\ion{Ca}{2} infrared triplet lines in individual giants.  These measurements have been extremely useful in tracing the age-metallicity relationship in the LMC and in providing estimates of the overall metallicity.  However, for some Magellanic Cloud clusters (e.g., NGC~2019 and NGC~2005), there is a disagreement between the metallicities from LMC-O91 and the slopes of the red giant branches measured by LMC-O98.  Abundance ratio questions could not be addressed, however, until high-resolution studies of LMC stars became available.  The study by Hill \\etal\\ (2000) included LMC clusters of a range of ages, including one old cluster: NGC 2210 ([Fe/H] = $-$1.75). The three stars observed in NGC 2210 have [O/Fe] values of 0.02, 0.19 and 0.21 dex, lower than are typical in GGC red giant stars. Smith \\etal\\ (2002) observed one star in NGC 1898 in the near-IR with PHOENIX on Gemini. [Ti/Fe] is low in this star, illustrating that, for this cluster at least, LMC clusters do not always exhibit the high \\afe{} abundances that typically belong to old GGCs.  In this paper, we report on abundances in four clusters in the Large Magellanic Cloud (LMC) observed with the MIKE spectrograph on Magellan. These clusters, NGC 1898, NGC 2005, NGC 2019 and Hodge 11 are globular clusters with ages from color-magnitude diagrams that are as old as the majority of GGCs (LMC-O98, LMC-J99) and have metallicities ranging from $-2.0 $ to $-1.0$. Therefore, they are similar to the kinds of GGC that have been important in deciphering the history of the Galactic spheroid. ",
        "conclusions": "Employing a linelist of $>$300 lines with laboratory-based gf-values, we have derived elemental abundances for the $\\alpha$-, iron-peak, and neutron-capture element groups of ten giant stars in four old globular clusters of the LMC.  In deriving the abundances of Sc, V, Mn, Co, Cu, Ba, La, and Eu, we took into account HFS in all of the lines employed.  Extensive numerical experiments were performed to elucidate the effects of differing choices of gf-values, stellar atmospheres, and stellar parameters on the abundances we derived.  While we find that many abundance similarities exist between the globular cluster stars in the LMC and our Galaxy (e.g., the ratios of [O/Fe], [Na/Fe], [Al/Fe], [Mg/Fe], and [Si/Fe]), the same is not true of all of the elements we studied.  In particular, we find differences {\\em within both} the $\\alpha$-element {\\em and} iron-group abundances.  We find lower-than-MWG-average values, and indeed, in some cases, clearly sub-solar values of [Ca/Fe], [Ti/Fe], [V/Fe], [Ni/Fe]. The LMC cluster giant star abundances of [Co/Fe], [Cr/Fe] and [Mn/Fe] may indicate additional offsets.  However, the available literature on the Galactic abundances of this trio of iron-peak elements is sparse (in one or the other of the globular cluster or field star populations) and further data are required to determine with greater certainty whether the differences in the abundances derived for the different groups are significant.  In the case of another iron-peak element, the behavior of [Cu/Fe] in our LMC clusters with respect to [Fe/H] appears to be constant with a value of $\\sim$$-$0.8.  While this is in marked contrast to the abundances observed in other MWG halo field and cluster stars, it does resemble the trend observed in \\wcen.  More relatively metal-rich LMC stars need to be observed in order to discern whether or not a rise in the Cu abundance exists with respect to the Fe abundance  With regards to the neutron-capture elemental abundance ratios of [Y/Fe], [Ba/Fe], and [Eu/Fe], we find the LMC star results to be similar to  the MWG-average values.  We compared the abundances derived for NGC1898 and NGC2019 against predictions of the scaled solar system contributions of the $r$- and $s$-process by Arlandini et al.\\ (1999) and Burris et al.\\ (2000).  We find that the abundances of neutron-capture elements Y, Zr, Ba, La, Nd, and Eu can largely be accounted for by the $r$-process, with, at most, a 20\\% contribution from the $s$-process to improve the agreement with Y, Ba, and Nd.  The abundance ratio distributions observed in red giant stars in the LMC globular clusters are markedly different from those found in the GGC red giants, the halo field red giants, and the red giant stars of the dSph systems.  Since the $\\alpha$ elements in the LMC clusters are not universally suppressed, and the ages of these clusters are old, we do not favor contributions by Type Ia SNe, but rather a unique star formation history that produced smaller amounts of Ca, Ti, V, Ni and Cu than in the Milky Way.  Possible explanations include a bias in the mass function of SNe that either exploded or whose ejecta were retained, or stars in the LMC being formed from material resulting from contributions by lower metallicity SNe than in our Galaxy. The cause of the low [Y/Fe] values seen in the stars of dSph systems does not operate in the LMC clusters, and marks another difference, in addition to the [Mg/Fe] and [Si/Fe] values, between the LMC and the dSph systems.  There do not appear to be universal trends among the satellite galaxies of the Galaxy. We conclude that many of the abundances in the LMC globular clusters we observed are distinct from those observed in the Milky Way, and these differences are intrinsic to the stars in those systems."
    },
    "0512/astro-ph0512074_arXiv.txt": {
        "abstract": "We present a technique to adaptively bin sparse data using weighted Voronoi tesselations (WVTs). WVT binning is a generalisation of Cappellari \\& Copin's (2001) Voronoi binning algorithm, developed for integral field spectroscopy. WVT binning is applicable to many types of data and creates unbiased binning structures with compact bins that do not lead the eye. We apply the algorithm to simulated data, as well as several X-ray data sets, to create adaptively binned intensity images, hardness ratio maps and temperature maps with constant signal-to-noise ratio per bin. We also illustrate the separation of diffuse gas emission from contributions of unresolved point sources in elliptical galaxies. We compare the performance of WVT binning with other adaptive binning and adaptive smoothing techniques. We find that the CIAO tool {\\it csmooth} creates serious artefacts and advise against its use to interpret diffuse X-ray emission. ",
        "introduction": "\\label{Introduction}  X-ray data are generally very sparse in nature. To deal with this problem, astronomers are often forced to either bin or smooth their data. The most commonly used techniques are simply binning to square blocks of a fixed size or convolving with a fixed kernel. However, due to the large dynamic range in many extended objects, ordinary binning and smoothing techniques are never able to capture structure on large scales without masking detail on smaller scales. This deficiency is the motivation for spatially adaptive algorithms.  With the advent of the two major X-ray satellites, {\\it Chandra} and {\\it XMM-Newton}, it is now possible to resolve fine morphological structures in extended X-ray emitting sources, such as galaxies, clusters, or supernova remnants.  This calls for new techniques to reliably extract spatial information. \\citet[][hereafter SF01]{Sanders} were the first to answer with a 2-dimensional adaptive binning algorithm, applicable to background-corrected intensity images and hardness ratio maps.  However, this algorithm is restricted to a limited set of bin sizes, which prevents it from being fully adaptive and from adjusting its resolution so as to keep the signal-to-noise ratio (S/N) constant. This creates jumps in S/N of a factor of $\\sim 2$, which, along with its quadrilateral bin shapes, can lead the eye and suggest structure that is not there.  Motivated by the different problem of analysing 2-dimensional optical integral field spectroscopic data, \\citet[][hereafter CC03]{Cappellari} developed an innovative adaptive binning technique using Voronoi tesselations. Their algorithm is able to smoothly adjust the bin size to the local S/N requirements and does not impose a Cartesian geometry on the image. Unfortunately, it can be used only with strictly positive, Poissonian or optimally weighted data whose S/N is guaranteed to add in quadrature. This prevents it from being useful in even simple situations in X-ray astronomy, involving data corrected for exposure map effects or background, or in creating hardness ratio maps.   In this paper, we generalise CC03's Voronoi binning technique so that it can be used with any type of data. The generalised algorithm makes use of Weighted Voronoi Tesselations (WVT), and combines the virtues of both CC03's and SF01's techniques. It is as robust as, and even more versatile than, SF01's code, yet retains the advantage of CC03's flexible bin sizes. The algorithm produces smoothly varying binning structures that are geometrically unbiased and do not lead the eye.  In section \\ref{ExistingAbinning} of this paper, we review the two binning techniques of SF01 and CC03 in more detail, pointing out their advantages and drawbacks. In \\S\\ref{WVTbinning}, we explain the functionality of the generalised WVT binning technique, and compare its performance to the two older algorithms in section \\ref{Performance}. Section \\ref{Applications} then demonstrates the utility of WVT binning in creating X-ray intensity images, hardness ratio maps, and temperature maps, and in disentangling the diffuse gas emission in elliptical galaxies from the contribution of unresolved point sources. Finally, \\S\\ref{Asmoothcomparison} quantitatively compares WVT binning to commonly used adaptive smoothing algorithms, before commenting on the availability of the code in section \\ref{Availability} and ending with conclusions in \\S\\ref{Conclusions}. ",
        "conclusions": "We have presented a generalisation of the Voronoi adaptive binning technique by \\citet{Cappellari}, broadly applicable to X-ray and other data. The generalised algorithm exploits the properties of weighted Voronoi tesselations, rather than the overly restrictive Gersho conjecture.  WVT binning is applicable to any type of data as long as there is a way to robustly calculate the S/N, and the S/N distribution changes smoothly over the size of a bin.  We have demonstrated the capabilities of WVT binning on exposure- and background-corrected X-ray intensity images, colour and temperature maps, and in isolating the diffuse gas emission in elliptical galaxies.  WVT binning overcomes the shortcomings of both Voronoi and quadtree binning, the latter of which is in growing use in X-ray astronomy. \\citet{Sanders_contourbinning} have recently published results using a ``contour binning'' algorithm, in which the bin boundaries follow the isophotes of an adaptively smoothed image. This creates very irregular and elongated bins, which lead the eye and introduce a shell-like appearance. We are unable to make a rigorous quantitative comparison with this technique, as the details are still unpublished.  However, we reemphasise that our WVT binning produces an unbiased distribution of compact bins, and does not lead the eye.  We have also demonstrated the pitfalls of adaptive smoothing, and regretfully advise against the use of the CIAO tool {\\it csmooth} for images of diffuse emission, as it creates very serious artefacts. If an adaptive smoothing technique has to be used, we recommend the XMMSAS tool {\\it asmooth} instead. However we urge that adaptively smoothed images be published only in conjunction with the smoothing scale map or an equivalent WVT binned image to facilitate the identification of real structures."
    },
    "0512/astro-ph0512597_arXiv.txt": {
        "abstract": "We present two new source extraction methods, based on Bayesian model selection and using the Bayesian Information Criterion (BIC).  The first is a source detection filter, able to simultaneously detect point sources and estimate the image background.  The second is an advanced photometry technique, which measures the flux, position (to sub-pixel accuracy), local background and point spread function.  We apply the source detection filter to simulated Herschel-SPIRE data and show the filter's ability to both detect point sources and also simultaneously estimate the image background.  We use the photometry method to analyse a simple simulated image containing a source of unknown flux, position and point spread function; we not only accurately measure these parameters, but also determine their uncertainties (using Markov-Chain Monte Carlo sampling).  The method also characterises  the nature of the source (distinguishing between a point source and extended source).  We demonstrate the effect of including additional prior knowledge. Prior knowledge of the point spread function increase the precision of the flux measurement, while prior knowledge of the background has only a small impact.  In the presence of higher noise levels, we show that prior positional knowledge (such as might arise from a strong detection in another waveband) allows us to accurately measure the source flux even when the source is too faint to be detected directly.  These methods are incorporated in SUSSEXtractor, the source extraction pipeline for the forthcoming Akari FIS far-infrared all-sky survey. They are also implemented in a stand-alone, beta-version public tool that can be obtained at http://astronomy.sussex.ac.uk/$\\sim$rss23/sourceMiner\\_v0.1.2.0.tar.gz ",
        "introduction": "\\label{introduction} Source extraction is close to ubiquitous in modern observational astrophysics.  The ability to identify and accurately quantify objects of interest in astronomical observations, in particular with reliable automated methods, is becoming ever more important with the advent of modern, large-area surveys.  It is crucial that we are able to ask precise, statistical questions of the data from these surveys.  Is there a source at a given location in the sky?  Is it point-like or extended?  And what set of parameters can define it?   Any science derived from the study of astronomical objects proceeds directly from accurate source extraction.  In order to extract sources from astronomical data, we typically face a number of challenges.  Firstly, there is instrumental noise.  It is often possible to measure this instrumental/observational characteristic and use this information to partially offset the effects. More problematic are any so-called 'backgrounds' to the observation.  These can be due to galactic emission, cosmological backgrounds, faint source confusion, or even simply emission from parts of the telescope itself.  These are often much harder to account for and often constitute an in-depth study in themselves.  A prime example is the extraction of sources from Cosmic Microwave Background (CMB) data \\citep[see e.g.][]{Vielva-01}.  We may also have to contend with systematic effects such as glitches that can be caused by cosmic ray hits on the detectors of space telescopes.  Because of these challenges and also because it is critical to exact the utmost precision from our (often very expensive to gather) data, we must strive to use all the available information when extracting sources.  This means not only using all available data samples, but also accurate noise estimates, measurements of the point spread function and also inclusion of any other prior knowledge we may have.  Over the years, a number of methods have been created in order to use various sets of information to obtain 'optimal' (subject to certain sets of assumptions) source extraction methods.  There are many techniques based on the concept of filtering data to enhance relatively the signal due to objects of a certain set of characteristics.  Examples of these include the matched, scale adaptive and wavelet filters \\citep[see e.g.][]{Vio-02, Barnard-04, Lopez-05, Barreiro-05}.  More recently, \\citet[][]{makovoz-05} have derived a filter of this type using the Bayesian formalism, thus allowing for the explicit inclusion of prior knowledge.  Fitting of the point spread function to image data has also been used as a way of accurately determining the position and flux of a (point) source, \\citep[see e.g.][]{Scott-02} The model-fitting methodology has been given a much more general grounding in statistical theory by \\citet[][]{Hobson-03} who have detailed a very general (and powerful) Bayesian framework for the extraction of sources.  Bayesian methodology has also used in the Poisson noise regime \\citep[see e.g.][]{Guglielmetti-2004}. There are a number of publicly available source extraction packages, which use a variety of the above methods (plus some other measures) in order to accurately extract sources.  These include, for example, DAOPHOT \\citep[][]{DAOPHOT-87} and Source Extractor \\citep[][]{SExtractor-96}.  Perhaps the most flexible of these approaches is that of using Bayesian statistics \\citep[see e.g.][]{JaynesStats-book, MackayInfoTheory-book}, as it allows one to ask very precise statistical questions of the data. This framework is also highly general, allowing the inclusion of all pertinent information. In this paper, we explore the use of Bayesian statistics for source extraction.  We present a pair of new methods based on this formalism, one for simultaneous source detection and background estimation/subtraction, and the other for an advanced form of source photometry that also allows the determination of the nature (point-like, extended etc) of the source.  The contents of this paper are therefore as follows. In \\emph{Section \\ref{section:methods}}, we present a general formalism for performing Bayesian source extraction.  We also detail two specific implementations. In \\emph{Section \\ref{section:results}} we apply these methods to a simulated data sets, in order to demonstrate their abilities. Finally, our conclusions are presented in \\emph{Section \\ref{section:conclusions}}. \\section[]{Methods} \\label{section:methods} In this section, we present a general formalism for performing Bayesian source extraction.  We then apply this formalism to derive two specific source extraction methods, with an eye to the analysis of modern, large photometric astronomical surveys (although their applicability is more general).  For this reason, both methods will address 2D (i.e. photometric image) data, although we note that the formalism extends to an arbitrary number of data dimensions.  Classic source extraction methodology divides the overall task into two distinct stages, source \\emph{detection} and source \\emph{photometry}. While the Bayesian paradigm allows for the possibility of a single, combined approach, the nature of the data we are considering dictates that we resist this. Modern photometric surveys are often large enough that such a combined approach is likely to be computationally prohibitive.  The methods we present below retain the two-stage approach, thereby proving computationally much quicker to use.  We note that in the following subsections, we will assume throughout that the noise on each image pixel is Gaussian, of known variance, and uncorrelated from pixel to pixel.  Additionally, when we are summing over pixels, we will always choose a subset of the image pixels that are local to the centre location we are considering.   A method for determining optimally such subsets is given in \\ref{subsection:dataSubset}.  The assumption of Gaussian noise warrants some discussion. In many real applications the noise distribution will naturally be close to Gaussian, e.g. when the dominant noise comes from well behaved instrumental noise. In other cases a Gaussian distribution might be inappropriate, e.g. in an context where the data are strictly non-negative.  In some such cases a Poisson distribution might provide a more natural description, when the photon statistics dominate. However, if the photon numbers are sufficiently high then a Gaussian model is an adequate approximation to the Poisson distribution.  This condition arises often in astronomy, e.g. when the sky background dominates.  In the case-study we are considering, observations with Herschel, the noise is dominated by the thermal background of the warm telescope primary and the Gaussian approximation is reasonable. It would be possible to generalise the method to include non-Gaussian noise distributions, including Poisson or log-Normal distributions but that investigation is beyond the scope of this paper.  \\subsection{General formalism} The essence of Bayesian data analysis is to create a reasonable parameterised model of the data.  These parameters can then be constrained by the data themselves, along with any available prior knowledge.  We begin with Bayes theorem.  \\begin{equation} P(\\theta|D,H)=\\frac{P(D|\\theta,H)P(\\theta,H)}{P(D|H)}\\,, \\label{BayesTheorem} \\end{equation}  Where $P(\\theta|D,H)$ is the posterior probability of the model parameters ($\\theta$), given the data $D$ and a hypothesis $H$. $P(D|\\theta,H)$ is the likelihood of the data (henceforth referred to as $\\cal{L}$, for simplicity) given a set of model parameters, $P(\\theta,H)$ represents any prior knowledge we may have about the likely values of the parameters, and $P(D|H)$ is the Bayesian Evidence. Bayes theorem provides the framework for our work.  We start with the likelihood.  If we are able to assess this, then (after applying a prior), we will have the posterior probability distribution, which is the result we require.  Following the normal route for uncorrelated, Gaussian noise, we have the following:  \\begin{equation} \\label{eqn:likelihood} {\\cal{L}} \\propto \\exp\\left(-{\\chi^2\\over2}\\right) \\end{equation}  \\begin{equation} \\label{eqn:chiSquared} \\chi^2=\\sum_{i=1}^{N_{pixels}}\\left({{d_i-m(\\theta)_i}\\over\\sigma_i}\\right)^2 \\end{equation}  and $d_i$ is the value of the $i^{th}$ data pixel of the subset of image pixels under consideration, $m_i$ is the corresponding value from a (parameterised) model of the signal and $\\sigma_i$ is the standard deviation of the (Gaussian) noise associated with that pixel.  Calculation of the likelihood function therefore depends on the parameterised model of the signal that we are considering.  In this case, the model will contain a source (point-like or extended).  It will also contain a representation of the astronomical/instrumental background, as well as possibly containing parameters describing instrumental characteristics For example, if not well-defined by independent measurements, the point spread function could be parameterised, and hence simultaneously measured by the model-fitting procedure.  Once the likelihood has been constructed, any prior knowledge that we have about the parameters can be included, in the form of the prior probability (a density function spanning the same parameter space as the likelihood).  This function might typically include information such as prior knowledge of the positions of sources etc, although it is perfectly acceptable to use an uninformative flat prior (i.e. equal valued at all points in parameter space), if one has no relevant prior knowledge (we note that this is the implicit assumption in maximum likelihood methods).  As the Evidence is a constant, normalising term, we now have the (unnormalised) posterior distribution.  We can map this distribution by calculating posterior values over a hypercube of parameter-space points or Markov-Chain Monte Carlo (MCMC) sampling. The peak of this distribution is our most likely solution, and (once normalised) the distribution as a whole provides the statistical confidence regions.  The posterior probability distributions of individual parameters can be obtained by marginalising over the other parameters \\citep[see e.g.][]{SiviaBayesian-book}.  This can be done in a number of ways. If MCMC sampling has been used to map the posterior probability distribution, then simply making a histogram of the samples using the values of a single parameter automatically gives the corresponding 1D marginalised distribution (a well-known and highly useful feature of sampling from the posterior).  If one were considering only a small number (three or fewer, say) of parameters then it may be feasible to calculate posterior values over a hypercube of parameter points and then marginalise numerically (although this is a very brute-force approach).  Or one can assume a functional form for the posterior and perform the marginalisation analytically.  One common choice for the functional form is that of a multivariate Gaussian, which is often a reasonable approximation to the posterior and is analytically tractable.  It also has the advantage that it can be completely specified by a parameter covariance matrix evaluated at the maximum a posterior point.  The method gives a complete analysis, given a particular choice of model.  However, the question of selecting a good model still remains. This can be addressed by the Evidence, which provides a relative measure of the probability of different models being the best-fit, given the data \\citep[see e.g.][]{JaynesStats-book}.  Bayesian Evidence is typically time-consuming to calculate.  This makes analytic approximations desirable, in terms of practicality.  In particular, the Bayesian Information Criterion \\citep[BIC,][]{Schwarz-78} provides a easily calculated approximation to the log(evidence).  \\begin{equation} \\label{eqn:BIC} BIC = -2 \\ln ({\\cal{L}}_{max}) + \\nu \\ln (N_{data}) \\end{equation}  Where ${\\cal{L}}_{max}$ is the maximum likelihood value for a given hypothesis, $\\nu$ is the number of free parameters in the model and $N_{data}$ is the number of (approximately equally weighted) data used.  When comparing how likely different models are, lower BIC values indicate higher probability of the model being the correct one.  Using model-selection criteria allows us to address the question of which from a range of models is the best description of the data, and to do so in a statistically rigorous way.  This becomes vital when one's data contains millions of sources, some point-like, some extended (and with different morphologies), and some not real at all, but rather the product of contamination.   \\subsection{Implementation: Bayesian source detection filter} The first implementation that we present of the above formalism is a Bayesian source detection filter. Source detection is necessary if one has observations of a region of sky but has no explicit knowledge of the positions of sources in the image (the case with many astronomical surveys).  Our task is therefore to analyse the entire image, identifying the positions where it is likely that there is a source present.  One consideration which is often critical for such source detection is speed of analysis.  Modern photometric surveys, in particular, often produce many large images, necessitating source detection methods that are computationally quick to apply.  With this in mind, we derive an analytic Bayesian solution to determine the relative probability (at each pixel position in an image) of the data being best described by an empty sky or a point source (with an unknown, uniform background in each case).  The two models we therefore consider are the following.  \\begin{description} \\item{\\tt Empty sky, uniform background.} This model consists solely of a flat, uniform background, described by a single parameter (the level of the background). \\item{\\tt Point source, uniform background.} This model builds on the empty sky model, adding a single point source, centred at the pixel currently being considered.  The point source is modeled as a circularly-symmetric 2D Gaussian profile of known FWHM.  This model has two parameters: the background level and the integrated flux of the source. \\end{description}  We will compare these models using BIC.  This means (see Equation \\ref{eqn:BIC} that we only need to calculate the maximum posterior value for each model.  By doing this at each (fixed) pixel position, we can therefore calculate a map of the relative evidence for point sources across the image.  Because we are considering (for each pixel) a fixed position, both models are comprised of a linear sum of fixed components.  This means that we can find analytic solutions in each case for the maximum likelihood values. Using the condition that the partial derivatives of the likelihood must be zero at the maximum likelihood solution, we can solve to find the following maximum likelihood solutions for each model.  For the point source model, we have the following description of the model. \\begin{equation} m_i = F {\\cal{P}}_i + B \\end{equation}  For the empty sky model, we have the following simple description of the model. \\begin{equation} m_i = B \\end{equation}  Where $m_i$ is the $i$th model pixel, $F$ is the source flux, ${\\cal{P}}_i$ is the (Gaussian) point spread function (normalised such that it integrates to unity) and $B$ is the uniform background.  >From this, we find the following analytic maximum likelihood solutions for $F$ and $B$.  \\begin{equation} F_{source} = {{\\gamma \\beta - \\delta \\epsilon} \\over {\\alpha \\beta - \\epsilon^2}} \\end{equation}  \\begin{equation} B_{source} = {{\\alpha \\delta - \\gamma \\epsilon} \\over {\\alpha \\beta - \\epsilon^2}} \\end{equation}  \\begin{equation} B_{empty} = {\\delta \\over \\beta} \\end{equation}  Where the calculated values used in the above equations are given by the following (with all sums being performed over the image pixels in a local region. See subsection \\ref{subsection:dataSubset} for a discussion of how to choose this region).  \\begin{equation} \\alpha = \\sum_{i=1}^{N_{pixels}}\\left({{{\\cal{P}}_i^2}\\over\\sigma_i^2}\\right) \\end{equation} \\begin{equation} \\beta =  \\sum_{i=1}^{N_{pixels}}\\left({{1}\\over\\sigma_i^2}\\right); \\end{equation} \\begin{equation} \\gamma =  \\sum_{i=1}^{N_{pixels}}\\left({{d_i{\\cal{P}}_i}\\over\\sigma_i^2}\\right)\\\\ \\end{equation} \\begin{equation} \\delta =  \\sum_{i=1}^{N_{pixels}}\\left({{d_i}\\over\\sigma_i^2}\\right); \\end{equation} \\begin{equation} \\epsilon =  \\sum_{i=1}^{N_{pixels}}\\left({{{\\cal{P}}_i}\\over\\sigma_i^2}\\right) \\end{equation}  By then feeding the best-fit model back into Equations \\ref{eqn:likelihood} and \\ref{eqn:chiSquared}, we obtain the maximum likelihoods.  We can therefore calculate the relative BIC at each pixel position. (we note that we implicitly use flat, uninformative priors in the preceding steps) The resulting map is an estimate of the (log of the) relative probability of there being a point source, rather than empty sky, at any given pixel position.  The local extrema of this map therefore give us the locations where one model is (locally) most favoured over the other.  Constructing the map so that (by convention) high values correspond to the point source model being more likely, we can identify the most likely source positions in the input image by identifying the local maxima in the map, subject to some minimum threshold value.  This method is closely modeled in some respects on the traditional filtering methods such as matched, scale-adaptive and wavelet filters.  It does, however, have several key advantages.  \\begin{enumerate} \\item{\\tt Simultaneous background estimation, subtraction.} In real astronomical data, background subtraction is a highly non-trivial task.  In particular, more traditional methods such as median filtering are biased by the presence of sources.  By performing the subtraction simultaneously, we largely avoid this problem. \\item{\\tt Proper accounting for flagged data and locally-varying noise.} Real astronomical images will typically have gaps due to flagging and uneven scan strategies, as well as point-to-point variations in noise levels.  This approach allows us to properly account for these effects by including an individual statistical weight (i.e. $1\\over\\sigma^2$) for each image pixel.  Similarly, setting a given weight to zero effectively flags out the corresponding datum (see Equation \\ref{eqn:chiSquared}).  This is mathematically well-defined; the principle challenge in such cases is in fact to estimate accurately the statistical weight (via the standard deviation) for each image pixel, which will depend on exactly how the image was created (for example, if it is the sum of many repeated observations, the multiple samples contributing to each image pixel can be used to estimate the standard deviation). \\item{\\tt Extensible method.} As it is based on a very flexible and general formalism, this source detection filter can be straightforwardly modified to accommodate more complex and realistic data models.  For example, many data are subject to 'glitches' (caused by cosmic ray hits on detectors).  By including in this method a third model of a single very high pixel value, it would be possible to distinguish between a source and a glitch-spike.   \\end{enumerate}  \\subsection{Implementation: Bayesian source photometry} \\clearpage \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[angle=-90, width=4.5cm]{f1a.ps} \\includegraphics[angle=-90, width=4.5cm]{f1b.ps} \\includegraphics[angle=-90, width=4.5cm]{f1c.ps} \\caption{The input flux image (left), a map of the background (as estimated by the source detection filter) (middle) and a map of the residuals between input flux image and background map (right) (with matched grey-scales).  As can be seen, even the presence of very bright sources does not appreciably bias the background estimation of our algorithm.  We also note that by setting a threshold of 550 (unusually high, due to small-scale fluctuations in the diffuse background, which act as a source of correlated noise) we are able to successfully detect 26 of the 33 sources (with 3 false detections). Simulated flux image courtesy of Philippe Andre, Bruce Sibthorpe and Tim Waskett.} \\label{fig:sourceDetection} \\end{minipage} \\end{figure*} \\clearpage Once a source has been detected, we wish to more completely measure and characterise it.  Considering only regions of the sky in which there is likely to be a source means that we can afford to devote substantially more computational effort to each candidate position. This is the principal advantage of performing source extraction in two distinct stages.  In this method, we again adopt the approach of fitting multiple models to the local data (again using the fact that we are interested in compact sources to minimise the data we must consider).  However in this case we will use a more in-depth approach, allowing more parameter to vary and mapping out the posterior probability distribution in each case.  The result will be more precise results (in particular, sub-pixel positional accuracy) and the determination of the errors on each parameter (without assumptions as to the form of the error distributions).  We proceed again by defining a number of models which we will fit to the data.  \\begin{description} \\item{\\tt Empty sky, uniform background.} This model consists solely of a flat, uniform background, described by a single parameter (the level of the background). \\item{\\tt Point source, uniform background.} This model builds on the empty sky model, adding a single point source at a given (parameterised) X, Y position.  The point source is modeled as a circularly-symmetric 2D Gaussian profile of known FWHM.  This model has four parameters: the background level, X and Y position and the integrated flux of the source. \\item{\\tt Extended source, uniform background.} This model is the logical extension of the point source model and is identical, with the exception that the FWHM is now allowed to vary as a model parameter (giving five in total).  This allows us to account either for circularly symmetric extended sources, or alternatively to measure the FWHM of  the point spread function, if this is not known. \\end{description}  We emphasise that there are many other models that can be usefully applied.  Examples would be non-circular extended sources, models where the noise is unknown or models where there are two or more adjacent (blended) sources.  For simplicity, we will again reply on the BIC for model comparison, although a full Bayesian Evidence calculation could be used (computing resources permitting).  As before, the likelihood functions are thus defined for any given set of parameter values of the relevant model.  Multiplying by the prior distribution for each model, we have the posterior for each model, which is mapped using MCMC sampling (except for the empty sky model, for which we only require the analytic best-fit solution, unless a prior is imposed).  The MCMC sampling returns the best-fit value for each model.  We use this to calculate BIC values and hence determine which model is mostly likely to be the best representation of the data.  This characterises the nature of the source in question.  Returning to the MCMC samples for the most likely model, we have also mapped the posterior probability distribution for that model.  From this we can straightforwardly determine the confidence intervals and best-fit values for all fitted parameters.  The power of this method lies in its ability to ask precise, statistical questions as to the nature of a source and to recover the theoretically optimal amount of pertinent information, given the data.  The flexibility of the Bayesian framework means that we are able to adapt this method, depending on the type of sources (and data) that we are expecting.  It is entirely realistic to deploy a whole battery of models, fitting each one in turn and determining which is the most likely representation.  \\subsection{Determining optimal data subset size} \\label{subsection:dataSubset} Because we are concerned with the extraction of compact sources, the above analyses need only consider a small subset of \\emph{local} data, for each source position.  This region should be large enough that we get good constraints on the source flux and local background, but small enough that our assumption of a flat background does not break down.  Our definition of the size and shape of this region will therefore have a direct impact on our source extraction.  We choose to determine an optimal region size in terms of a minimised BIC value (and hence maximised Bayesian Evidence).  This will give us a data model that best describes our data (for the types of model we are considering here).  If we define the region as circular, then we reduce this problem to an optimisation (in BIC) with respect to the radius of the region.  One complication is that for BIC comparisons to be valid, we require the same data set to be considered in each case.  This would plainly not be true if we simply used the data inside the circular \\emph{region-of-interest}.  To avoid this problem, we also define a larger super-region (also circular), and label the super-region image pixels that lie outside the region-of-interest as \\emph{external pixels}.  We then redefine our model as fitting the source/background to the \\emph{region-of-interest}, plus allowing additional free parameters for each of the \\emph{external pixels}, so that they are fitted exactly and do not contribute to the chi-squared of the model fit. Therefore, the \\emph{external pixels} will contribute to the BIC solely as extra parameters, the number of which will vary depending on the radius of the \\emph{region-of-interest}.  These nuisance parameters are not trivial to deal with and we highlight that the above procedure makes the simplifying assumption that the nuisance parameters can be fitted to the data with no uncertainty, so that marginalisation over them is not necessary.  In practice this is not true and would alter the BIC calculation (via the maximum likelihood value).  This could be accounted for in the photometry method because MCMC methods can straightforwardly include large numbers of nuisance parameters, which can be marginalised over without extra effort. Doing so will incur the need for longer sampling chains to be generated, to ensure adequate convergence.  One peculiarity of this procedure is that the BIC has a weak dependence (going as the logarithm) on the radius of the super-region (and hence the maximum possible region radius). While this is clearly undesirable, the effect will be small for reasonable radius ranges.  The minimum sensible region radius will typically be dictated by the FWHM of the point spread function, with perhaps a radius equal to the FWHM being a reasonable starting point.  The maximum region radius is less well-defined, but a value of four or five times the FWHM would seem intuitively reasonable, and our experiences in this paper suggest that this is not unreasonable.  For the case of fitting a single source (i.e. when we are applying source photometry), this process is unambiguous. In the case of the source detection filter, where we may have many detected sources (and that number may change as we optimise with radius), we need to choose what metric we will optimise.  In general, this choice will depend on the exact nature of both the data and the science in question..  One simple approach (and the one that we have adopted here) is to use an intermediate radius case (say, twice the FWHM of the point spread function) as a starting point, identifying the sources detected in this case.  We then use as our metric the sum of BIC values for the fits to these sources.  This procedure gives us a way of selecting an optimised definition for a local region of the data.  This both optimises the performance of the source extraction algorithm, but also means that the user need not waste any time optimising by trial and error.   \\subsection{Prior knowledge} A strength of the Bayesian formalism is its explicit inclusion of prior knowledge.  In the case of source extraction, one typically assumes that the noise characteristics and the point spread function are known (although this need not be the case).  One could also assume prior knowledge about any of the fitted parameters; for example, source position may have already been determined in another observing band.  It is also possible to select priors on the basis of more general knowledge.  For example, if one is attempting to detect a population of galaxies, it may be reasonable to assume a power law distribution for the source flux (e.g. from a model of the galaxy population).  Even in the absence of such knowledge, one could still choose the \\emph{Jeffreys prior} (a power law with index of $-1$), which is the indifference prior for a positive-only scaling parameter.  The source flux is also of particular note because it will typically (although not always; e.g. Sunyaev-Zel'Dovich effect for galaxy clusters in CMB observations) be subject to the constraint of being non-negative.  In this case, it is important to properly apply this as a prior constraint.  Because the above source photometry method uses MCMC sampling, it is straightforward to quantify the assumptions on the prior.  In section \\ref{subsection:sourcePhotometry}, we show examples of this.  The source detection filter relies on analytic solutions in order to give plausible speed of analysis.  This makes the application of non-top-hat priors more difficult, if convenient analytic solutions are to be possible.  The potential size of this topic takes it beyond the scope of this paper, but we note that the exploration of different priors represents a largely untapped area where source detection methods could be improved.   \\clearpage \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[angle=0, width=8cm]{f2a.ps} \\includegraphics[angle=0, width=8cm]{f2b.ps} \\caption{The simulated images on which photometry was performed.  Both images contain the same underlying signal, consisting of a uniform background (of level 0.5 units), plus a Gaussian point source with position (relative to the image centre) of (0.3, 0.4) pixels, a FWHM of 5.2 pixels and in integrated flux of 10 units (corresponding to a peak height of 0.231 units).  The left-hand image has Gaussian noise with rms of 0.075, the right-hand has rms noise of 0.3 (i.e. higher than the peak of the source).} \\label{fig:simulatedData} \\end{minipage} \\end{figure*} \\clearpage  \\section[]{Results} \\label{section:results} In this section, we present example results from the two methods detailed in the previous section.  We highlight the speed of analysis of these methods.  Running on a desktop machine (using two 2.4~GHz AMD Opteron 250 CPUs) and implemented in IDL, the source detection filter processed $9 \\times 10^4$ pixels per second (a $784 \\times 912$ pixel image in eight seconds), and the photometry method was able to analyse one source every nine seconds (producing $10^5$ MCMC samples per model, per source).  At this rate, for example, the whole Akari all-sky survey could be source detection filtered in four days and four hours (assuming $40,000$ square degrees of coverage, with $8 \\times 8 arcsecond$ image pixels and four observing bands, using a single desktop machine.  \\subsection{Bayesian source detection filter} Figure \\ref{fig:sourceDetection} shows images from the analysis of a simulated Herschel-SPIRE \\citep[see e.g.][]{HerschelOverview-2004} observation of a number of point sources, along with a diffuse galactic foreground (data courtesy of Philippe Andre, Bruce Sibthorpe and Tim Waskett).  Shown are the input flux image, a map of the background (as estimated by the source detection filter), and a map of the residuals between input flux image and background map.  The background map is created using the maximum a posterior estimate of the model background at each pixel position.  In each case, the model used is that which is most likely, on the basis of BIC score.  The residuals map is created by subtracting the background map from the original input flux image.  The residuals will therefore contain the point sources, plus any imperfections in the background estimation.   \\subsection{Bayesian source photometry} \\label{subsection:sourcePhotometry} \\clearpage \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[angle=0, width=8cm]{f3a.ps} \\includegraphics[angle=0, width=8cm]{f3b.ps} \\includegraphics[angle=0, width=8cm]{f3c.ps} \\includegraphics[angle=0, width=8cm]{f3d.ps} \\includegraphics[angle=0, width=8cm]{f3e.ps} \\caption{1D marginalised posterior probability distributions for the five parameters of the extended source photometry model (solid line).  Also shown are the case where the FWHM prior is known perfectly (dashed line) and where both the FWHM is known and there is a Gaussian prior on the background (dot-dash line).  The input values are marked by vertical, dotted lines.} \\label{fig:compactSourcePosterior} \\end{minipage} \\end{figure*} \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[angle=0, width=8cm]{f4a.ps} \\includegraphics[angle=0, width=8cm]{f4b.ps} \\caption{1D marginalised posterior probability distributions for two parameters of the point source photometry model. Prior positional knowledge has been included, in the form of a Gaussian prior on both X and Y (FWHM of 0.1 pixels) In this case, the rms noise of the observation has been increased four-fold, so that in the absence of the prior, the BIC value would favour an empty sky.  This shows the case where a source has been detected to high precision in another band, but is very faint in this band.  The Bayesian formalism allows us to fully and properly account for this.} \\label{fig:positionPriorPosterior} \\end{minipage} \\end{figure*} \\begin{figure*} \\includegraphics[angle=0,width=16cm]{f5.ps} \\caption{Contour plots of the 2D marginalised posterior probability distributions for the parameters of the 'compact' source photometry model (shown are the 68\\% and 95\\% confidence regions, plus the maximum a posterior point).  The contours are found using smoothed 2D histograms of the MCMC samples.  This gives estimates of the marginalised 2D posterior probability distributions for these parameter combinations.  Note that we exclude the 'Y' parameter, as its behavior simply mimics that of 'X'.  The FLUX and FWHM of the point spread function are correlated.  The FLUX and BACKGROUND are negatively correlated.  The FWHM and BACKGROUND also show a slight negative correlation.  As expected, the X position is uncorrelated with these three other parameters.} \\label{fig:2Dposterior} \\end{figure*} \\clearpage The analyses in this sub-section are carried on on a simple, simulated test image (shown in Figure \\ref{fig:simulatedData}). The image contains a single point source on a uniform background, with uncorrelated Gaussian random noise added to each pixel.  While this is a benign data set, it is instructive to consider such an idealised case in order to better understand features of the algorithm.   Figure \\ref{fig:compactSourcePosterior} shows the 1D marginalised posterior probability distributions for a variety of cases.  The solid lines show a five parameter 'compact' source model fitted to the data.  The five parameters are a flat background, the FWHM of a Gaussian point spread function, the flux of the source, and its X and Y co-ordinates within the image. The dashed lines show the 1D marginalised posteriors for the case where the FWHM of the point spread function is known(for example, it has been measured independently of this 'observation').  The dotted lines show the 1D marginalised posteriors for the case where the FWHM of the point spread function is know \\emph{and} we have prior knowledge of the level of the background. Figure \\ref{fig:positionPriorPosterior} shows the 1D marginalised posteriors arising from analysing the same source with four times the rms Gaussian noise.  The FWHM is taken as known, as is prior knowledge of the the source position.  This simulates the case where a source has been strongly detected at another band and we now wish to find an estimate of that source's flux in this band. Figure \\ref{fig:2Dposterior} shows examples of 2D marginalised posteriors for the five parameter 'compact' source model.  These illustrate the different correlations that exist between the fitted parameters. ",
        "conclusions": "\\label{section:conclusions} In this paper, we have described a Bayesian formalism for the extraction of sources from astronomical data and have used it to derive two new source extraction methods.  We then demonstrated the methods on simulated data.  The source detection filter is a deliberately uncomplicated implementation of this formalism; it is designed to analyse images quickly, something that is often crucial given the size of many modern astronomical surveys.  Estimation of the image background is an often-overlooked (and highly non-trivial) aspect of source extraction and the simultaneous estimation performed by our filter makes unbiased background subtraction much more tractable.  An additional point not to be under-estimated is that by combining background subtraction and source detection, we have created a method that has essentially only one user-defined parameter (threshold), substantially simplifying its use.  We applied this filter to a deliberately challenging simulated image. The presence of a strong diffuse astronomical background introduces fluctuations on similar angular scales to the point spread function, presenting a particular challenge for source extraction. In spite of this, we are still able to detect the majority of sources, with only a few spurious detections.  If computationally fast ways can be found to better model this background (work beyond the scope of this paper), even more impressive results may be possible in the future.  Once a candidate source position has been identified, we wish to characterise the source as precisely as possible. The advanced photometry method allows us to do just that.  It can determine the flux, position (to sub-pixel accuracy), local background and (if required) point spread function FWHM, along with the uncertainties on those estimates.  Furthermore, it allows the meaningful comparison of different models, allowing us to determine (in an automated way) whether any given source is point-like, extended or even just a patch of empty sky.  We can also include any additional prior knowledge we may have about the source. For example, if the FWHM is known then the precision of our flux estimate is improved.  With prior positional knowledge (from a strong detection in another band), we can obtain a flux estimate even when there is insufficient evidence from the data alone to identify a source.  This formalism allows us to ask precise, statistical questions of our data.  We are able to include all pertinent information, giving us the best possible measurement and characterisation of the sources.  We can also determine a number of figures-of-merit, such as Bayesian Evidence, BIC and reduced chi-squared, all of which give measures of the quality of the extraction.  Parameter space searching techniques such as MCMC sampling allow us to recover the statistical uncertainties on our measurements while making minimal assumptions.  And model selection techniques allow us to ask which of a range of models best characterise any given source.  In conclusion, in this paper we present the following. \\begin{enumerate} \\item{\\tt A Bayesian formalism for the detection and extraction of compact sources from astronomical data} \\item{\\tt The derivation of an analytic source detection filter that simultaneously detects point sources and estimates the image background} \\item{\\tt The detailing of an advanced photometry method, which determines source parameters such as flux and position (to sub-pixel accuracy), as well as their uncertainties. It also allows us to determine the nature of the source (point-like, extended) and to include any prior knowledge we may have, thus enhancing the precision of our results} \\item{\\tt A method for optimising the local region from which data should be used to make the source fits} \\end{enumerate}  Bayesian source extraction is a highly powerful and (perhaps just as importantly) immensely flexible methodology.  The ability to adapt our methods to the peculiarities of the data we are considering is a key degree of freedom when dealing with real astronomical data.  Bayesian methods have historically been limited by lack of computing power; this is demonstrably no longer the case, giving us an array of new statistical tool with which to improve astronomical source extraction and hence the astrophysical science that depends upon it.  The methods described in this paper have been implemented as a beta-version, publicly available software tool (written in IDL).  The code plus associated documentation and test data can be obtained from the following URL:  http://astronomy.sussex.ac.uk/$\\sim$rss23/sourceMiner\\_v0.1.2.0.tar.gz"
    },
    "0512/astro-ph0512583_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} Space Telescope observations of the $z$=2.38 \\lya-emitter over-density associated with galaxy cluster J2143-4423, the largest known structure (110 Mpc) above $z=2$. We imaged 22 of the 37 known \\lya -emitters within the filament-like structure, using the MIPS 24\\m\\ band. We detected 6 of the \\lya -emitters, including 3 of the 4 clouds of extended ($>$50kpc) \\Lya emission, also known as \\Lya Blobs. Conversion from rest-wavelength 7\\m\\ to total far-infrared luminosity using locally derived correlations suggests all the detected sources are in the class of ULIRGs, with some reaching Hyper-LIRG energies. \\Lya blobs frequently show evidence for interaction, either in {\\it HST} imaging, or the proximity of multiple MIPS sources within the \\Lya cloud. This connection suggests that interaction or even mergers may be related to the production of \\Lya blobs. A connection to mergers does not in itself help explain the origin of the \\Lya blobs, as most of the suggested mechanisms for creating \\Lya blobs (starbursts, AGN, cooling flows) could also be associated with galaxy interactions. ",
        "introduction": "The 110 Mpc filament mapped out by 37 Ly$\\alpha$-emitting objects around the $z$ = 2.38 galaxy cluster J2143-4423 is the largest known structure above $z=2$ \\citep{pal04}, comparable in size to some of the largest structures seen in the local Universe \\citep[i.e. the Great Wall,][]{gel89}. Initially identified from narrow-band imaging tuned to \\Lya \\ at $z$=2.38, it has since been spectroscopically confirmed \\citep{fra04,frb04}. In addition to its compact \\lya -emitters, this high-redshift ``Filament'' is also home to four extended \\lya -emitting clouds, known more commonly as Lyman $\\alpha$ blobs.  The Ly$\\alpha$ blob is a relatively new class of objects found among high-redshift galaxy over-densities \\citep{ste00,kee99,pal04}. While similar in extent ($\\sim$100 kpc) and \\Lya flux (\\ab 10$^{44}$ ergs s$^{-1}$) to high-redshift radio galaxies, blobs are radio quiet and are therefore unlikely to arise from interaction with jets. Current surveys have reported the discovery of roughly 10 of these giant \\Lya blobs, but they are not isolated high redshift oddities. \\cite{mat04} have demonstrated that the blobs are part of a continuous size distribution of resolved ($>$16 arcsec$^2$) \\Lya emitters, with more than 40 presently known.  One of the standing mysteries of the blobs is the source of their energy, as the measured ultraviolet flux from nearby galaxies is insufficient to produce the observed \\Lya fluxes. One possibility is that the \\Lya blobs are powered by supernova-driven superwinds \\citep{ohy04}, driving great plumes of gas into the surrounding ambient medium and producing shocks. An obscured AGN is another model, with the exciting ultraviolet illumination escaping along different lines of sight \\citep[i.e.,][]{bas04}. Cooling flows have also been suggested \\citep{fard01,fra01} as a possible power source.  There is growing evidence that \\Lya blobs mark regions of extreme infrared luminosity. Submillimeter flux has been detected in two of the giant \\Lya blobs, SMM J221726+0013 \\citep{chp01} and SMM J17142+5016 \\citep{sma03}. The submm source SMM J02399-0136 is also likely surrounded by a \\Lya blob halo, as its \\Lya emission covered most of a 15$\\arcsec$ slit \\citep{ivi98}. Most recently, \\cite{gea05} has detected submm flux from four of the smaller ($<$ 55 square arcsec) and less luminous ($<$ 2$\\times$10$^{43}$ L$_{\\odot }$) \\Lya blobs from \\cite{mat04}: LAB5, LAB10, LAB14 \\& LAB18. Also, \\cite{dey05} has discovered a single \\Lya blob (SST24 J1434110+331733) in the NOAO Deep Wide-Field Survey with strong 24\\m\\ flux (0.86 mJy).  In this paper we discuss the {\\it Spitzer} 24\\m\\ observations of these $z$=2.38 \\lya -emitters, both compact sources and blobs. We estimate the total far-infrared luminosity for all detections and discuss the possibility of a connection between mergers and \\Lya blobs. We assume an $\\Omega _{M}$=0.3, $\\Omega _{\\Lambda }$=0.7 universe with H$_o$=70 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0512/astro-ph0512060_arXiv.txt": {
        "abstract": "{We have conducted low-frequency radio observations with the Giant Metrewave Radio Telescope (GMRT) of 40 new hard X-ray sources discovered by the INTEGRAL satellite. This survey was conducted in order, to study radio emissions from these sources, to provide precise position and to identify new microquasar candidates. From our observations we find that 24 of the X-ray sources have radio candidates within the INTEGRAL error circle. Based on the radio morphology, variability and information available from different wavelengths, we categorize them as seventeen Galactic sources (4 unresolved, 7 extended, 6 extended sources in diffuse region) and seven extragalactic sources (2 unresolved, 5 extended). Detailed account for seventeen of these sources was presented in earlier paper. Based on the radio data for the remaining sources at 0.61 GHz, and the available information from NVSS, DSS, 2MASS and NED, we have identified possible radio counterparts for the hard X-ray sources. The three unresolved sources, viz IGR J17303$-$0601, IGR J17464$-$3213, and IGR J18406$-$0539 are discussed in detail. These sources have been identified as X-ray binaries with compact central engine and variable in X-ray and in the radio, and are most likely microquasar candidates. The remaining fourteen sources have extended radio morphology and are either diffuse Galactic regions or extragalactic in origin.  \\keywords {stars : X-ray binaries --  X-ray : galaxies -- X-ray : sources : INTEGRAL sources }-- Techniques : interferometry} ",
        "introduction": "Many new hard X-ray emitting sources have been discovered during the deep Galactic plane survey by {\\emph{INTEGRAL}} satellite mission. The IBIS instrument has a point source location accuracy (PSLA) of typically 1 -- 3$'$ within a large field of view 29$^\\circ$ $\\times$ 29$^\\circ$ (Ubertini et al. 2003). 55 new hard X-ray sources have been reported in literature. A majority of these sources are believed to be Galactic X-ray binaries  with a compact object orbiting a companion star (Bird et al. 2004). Some of these sources are identified as AGNs, radio galaxies, pulsars, CVs and dwarf nova. A detailed study of X-ray sources in the multi-wavelength band is essential to understand the emission mechanism and the  accretion process on to the compact companion  neutron stars (NS) or black holes (BH). The radio imaging of these sources can establish whether some of these are radio  emitting X-ray binaries (REXBs)  and show any  microquasar like features. Due to their similarity with quasars, the jet feature in microquasars provide important information about the underlying physical phenomenon and the possible disk-jet connection which may power the observed emission in different wave bands. Their X-ray, infrared and radio properties can lead to classification schemes.  \\par The detailed radio observations for these sources were made immediately after their discovery using the GMRT, to find the possible radio counterparts within the location error box of the X-ray source, to measure precise position if detected and to study the low frequency radio nature of the hard X-ray sources. The radio morphology of a source also provides its identification, viz Galactic, which are mainly compact (Becker et al. 1990) or extragalactic mostly extended (Jackson 1999).  At meter wavelengths REXBs show point source morphology (Pandey {et al.} 2005a). We have also analyzed the available NVSS images at 1.4 GHz and other  radio data in order to understand the radio spectrum of these sources. The archival data, from DSS, 2MASS and NED images is also used in our analysis to facilitate a complete study of these sources. \\begin{table*}[ht] \\begin{center} \\caption{List of target INTEGRAL sources observed with GMRT} \\begin{tabular}{p{25.7mm}p{15mm}p{11mm}p{13mm}p{15mm}p{41.3mm}p{46.5mm}ll} \\hline \\hline Source           &Type  &Integral    &Variable                      &X-ray      &X-ray/optical/&No.of sources \\\\ &      &Pos. Unc.   &100s--1ks                     &Flux       &UV/IR/Radio   &in the X-     \\\\ &      &1.6$\\sigma$ &$\\frac{\\rm IBIS}{\\rm ISGRI}$  &15--40~keV &sources in X- &ray error circ.       \\\\ &      &            &                              &(mCrab)    &ray error circ.   &       \\\\ \\hline IGR J00370$+$6122  &HMXB$^{1,2}$   &2$'$        &Yes         &-         &BD$+$6073                     &\\\\ IGR J01363$+$6610  &HMXB$^{2}$       &2$'$        &Yes         &17        &HD9603                        &\\\\ IGR J16167$-$4957  &-      &6$'$$'$     &-           &2         &-                             &\\\\ IGR J16195$-$4945  &HMXB(?)$^{*,10}$   &16$'$$'$    &-           &-         &HD146628                      &\\\\ IGR J16207$-$5129  &-      &2$'$        &-           &-         &HD146803                      &\\\\ IGR J16358$-$4726  &LMXB(?)$^{3,4}$   &0.6$'$$'$   &Yes         &20-50     &2MASS J163553$-$472539        &\\\\ &Pulsar$^{5}$ &            &            &4.63      &                              &\\\\ IGR J16393$-$4643  &HMXB(?)&2$'$        &Yes         &3         &                           &20$^{**}$\\\\ &Pulsar$^{5,6,7,11}$ &            &            &          &                              &\\\\ IGR J16558$-$5203  &-      &8$'$$'$     &-           &-         &1RXS J165605$-$520345         &\\\\ &       &            &            &          &USNO-B1.0 0379$-$00008129     &\\\\ IGR J17195$-$4100  &-      &8$'$$'$     &Yes         &-         &1RXS J171935$-$410054         &\\\\ &       &            &            &          &USNO-B1.0 0489$-$00511283     &\\\\ IGR J17200$-$3116  &-      &9$'$$'$     &Yes         &-         &1RXS J172006$-$311702         &\\\\ IGR J17252$-$3616  &HMXB$^{8}$&2$'$        &-           &-         &IRAS 17220$-$3615             &\\\\ &Pulsar &            &            &          &NVSS J172510$-$361614         &\\\\ &       &            &            &          &HD319824                      &\\\\ IGR J17254$-$3257  &-      &14$'$$'$        &-           &-         &1RXS J172525$-$325717         &\\\\ &       &            &            &          &USNO-B1.0 0570$-$00727635     &\\\\ IGR J17285$-$2922  &XB$^{9}$&2$'$       &Yes         &-         &IRAS 17252$-$2922             &\\\\ &       &            &            &          &[T66b]320                     &\\\\ IGR J17303$-$0601  &LMXB   &7$'$$'$        &Yes         &-      &H1726$-$058                   &\\\\ &       &            &            &          &USNO$-$A2.0 0825$-$10606993   &\\\\ &       &            &            &          &1RXS J173021.5$-$055933       &\\\\ IGR J17456$-$2901  &-      &1$'$        &-           &-         &1LC G359.923$-$00.013   &90$^{**}$ \\\\ IGR J17460$-$3047  &-      &2$'$        &-           &-         &                        &80 IR sources$^{**}$\\\\ IGR J17464$-$3213  &BHC$^{*}$&0.5$'$    &Yes         &60        &H1743$-$322                   &\\\\ &LMXB$^{3}$&         &            &          &2MASS 17461525-3213542 &\\\\ &          &         &            &          &USNO-A2.0 0525-294112269         &\\\\ IGR J17475$-$2822  &Sgr B2$^{5}$&2--3$'$    &-           &-         &                     &200 $^{**}$\\\\ IGR J17488$-$3253  &-      &12$'$$'$    &Yes         &-         &1RXS J174854.7$-$325444       &\\\\ IGR J18027$-$2016  &Pulsar$^{6}$ &1$'$        &-           &4.06      &HD312525                      &\\\\ &       &            &            &          &1LC G000.683$-$0.035          &\\\\ &       &            &            &          &IRAS 17594$-$2021             &\\\\ IGR J18406$-$0539  &-      &2--3$'$    &-           &-          &IRAS 18379$-$0546              &\\\\ &       &            &            &          &AX J1840.4$-$0537         &\\\\ &       &            &            &          &NVSS J184037$-$054317         &\\\\ &       &            &            &          &GSC2.2&&\\\\ IGR J18450$-$0435  &-      &2--3$'$    &-           &-          &IRAS 18422$-$0437              &\\\\ &       &            &            &          &PMN J1845$-$0433              &\\\\ IGR J18490$-$0000  &-      &2--3$'$    &-           &-         &-                              &\\\\  \\hline \\end{tabular} \\footnotemark{High mass X-ray binary$^{1}$, Reig {et al.} 2005$^{2}$, Bird {et al.} 2004$^{3}$, Low mass X-ray binary$^{4}$, Revnivtsev {et al.} 2004$^{5}$, Lutovinov {et al.} 2005$^{6}$,Boudaghee {et al.} 2005$^{7}$, Zurita {et al.} 2005$^{8}$, X-ray binary$^{9}$, Sidoli {et al.} 2005$^{10}$, Soldi {et al.} 2005$^{11}$ http://isdc.unige.ch/~rodrigue/html/igrsources.html$^{*}$, $^{**}$ field sources identified at different wavelengths} \\end{center} \\end{table*} ",
        "conclusions": "We have presented the radio analysis of 23 of the newly discovered {\\emph{INTEGRAL}} hard X-ray sources. Most of these sources are X-ray binaries; however, the identification of AGNs, radio galaxies, X-ray novas, CVs and pulsars are the other important byproducts. Among the twenty three sources observed, seventeen have a possible radio counterpart detected at radio wavelengths. The position offset of the possible radio counterparts with respect to the {\\emph{INTEGRAL}} position is of the order of few arc minutes. The consistent position provided by the GMRT will allow the search for infrared/optical counterparts for these sources to be detected. Based on the radio morphology of these source we have further grouped them into:\\\\ (a) Galactic point source,\\\\ (b) extended Galactic sources and sources in diffuse Galactic emission,\\\\ (c) extragalactic sources.\\\\ Three sources belong to group (a) and are REXBs. We carried out a detailed study of these three sources and their possible counterparts at other wavebands. Based on the variability in the radio and X-ray windows along with the information available about the counterparts in other wavebands we infer that IGR J17303$-$0601, IGR J17464$-$3213 and IGR J18406$-$0539 are possible microquasar candidates. However, devoted radio observations are necessary to confirm the jet emission from these sources. We have also detected a variable compact radio source within the field of IGR J17464$-$3213. Eleven sources can be associated with group (b), the diffused Galactic region, and three sources satisfy the radio morphology of extragalactic sources, group (c). No radio counterpart was detected for the remaining sources.  To conclude, we highlight that our observations were very important in pinpointing the possible microquasar candidates from the list of 40 {\\emph{INTEGRAL}} sources observed at radio wavelength. We have performed repeated observations on these sources of our interest in Cycle 7 to look for the radio variability and the data has to be analyzed.   \\begin{acknowledgement} I would also like to thank Dr. M. Ribo, Dr. D. Green and Dr. D. Ojha for the useful discussions and Prof. V. Kulkarni, Dept. of Physics, Mumbai University for his constant support and suggestions.  \\end{acknowledgement}"
    },
    "0512/astro-ph0512126_arXiv.txt": {
        "abstract": "We study polarization from scattering of light on a cloud in radial motion along the symmetry axis of an accretion disk. Radiation drag from the disk and gravitational attraction of the central black hole are taken into account, as well as the effect of the cloud cooling in the radiation field. This provides us with a self-consistent toy-model for predicted lightcurves, including the linear polarization that arises from the scattering. Strong gravitational lensing creates indirect images; these are formed by photons that originate from the disk, get backscattered onto the photon circular orbit and eventually redirected towards an observer. Under suitable geometrical conditions the indirect photons may visibly influence the resulting magnitude of polarization and light-curve profiles. Relevant targets are black holes in active galactic nuclei and stellar-mass Galactic black-holes exhibiting episodic accretion/ejection events. ",
        "introduction": "\\label{intro} The present-day evidence for black holes relies almost entirely on information carried by electromagnetic waves. X-rays play a particular role (\\cite{sew95}): they are supposed to emerge from gas near a black hole horizon and bring us imprints of physical processes and conditions in the place of their origin and along the ray path. Various spectral and lightcurve patterns have been identified as likely signatures of supermassive black holes in galactic nuclei and stellar-mass black holes. These features presumably arise when matter is accreted from an immediate vicinity of the black hole. The ultimate goal of this effort is to prove the existence of event horizons in `real' nature and this way to discover black holes. Clearly, an affirmative proof is a great challenge that may still be far ahead of us. Here we study a related task, which appears to be somewhat easier on the technical side, though it also represents an unresolved issue as yet: searching for rays of photons that encircle an ultra-compact star or a black hole.  Accretion disks represent a common way of feeding black holes and generating photons, which then experience strong-gravity while travelling to a distant observer (\\cite{kat98}; \\cite{kro99}). Outside the disk plane density is much less, however, even regions near axis are not empty: winds and jets emerge, roughly along the symmetry axis. We address a question whether future observations of polarized time-dependent signal can help recognizing the effects of strong-gravity light bending, and what features are expected in light curves. To this aim we examine a model of warm clouds moving radially along the symmetry axis. Primary photons from the disk are Thomson-scattered and polarized. We assume that the process takes place near a Schwarzschild black hole, where the higher-order (indirect) light rays contribute to the observed  radiation flux. The resulting modulation of intensity and polarization magnitude can reach a non-negligible level under suitable geometrical alignment of the black hole, the cloud and the observer. We treat the interaction of the cloud with the radiation and gravitational fields in the relativistic framework (\\cite{abr90}; \\cite{vok91}; \\cite{kea01}). This approach provides us with a self-consistent description of the cloud motion and the resulting observed signal. Both ingredients are conveniently expressed in terms of the radiation stress tensor.  There is the evidence for jets being formed only a few tens gravitational radii from a supermassive central black hole (\\cite{jun99}). The emission mechanisms producing the observed high-energy photons (X- and $\\gamma$-rays) are likely non-thermal, but it is not clear whether the synchrotron emission or the inverse Compton emission dominates in each particular source (\\cite{har02}). Here we concentrate on the latter case (within the Thomson approximation), which seems to be relevant for radio-quiet sources with the ambient radiation acting on particles and fluids and determining their terminal (equilibrium) speed; see \\cite{noe74}; \\cite{ode81}; \\cite{sik81}; and \\cite{phi82} for original papers. This dynamical influence of the radiation field has been studied further by Sikora et al.\\ (\\yearcite{sik96}), Renaud \\& Henri (\\yearcite{ren98}), Tajima \\& Fukue (\\yearcite{taj98}), Fukue (\\yearcite{fuk99}) and Watarai \\& Fukue (\\yearcite{wat99}). Ghisellini et al.\\ (\\yearcite{ghi04}) proposed a model of aborted (failed) jets in which colliding clouds and shells occur very near a black hole and are embedded in strong radiation. According to their scheme, most of energy dissipation should take place on the symmetry axis of an accretion disc where the individual clouds either move away from the black hole, or they fall back. The process of gravitational and radiative acceleration of plasma was studied also by Fukue et al.\\ (\\yearcite{fuk01}) and Fukue (\\yearcite{fuk05}), who examined the efficiency of collimation towards the disk axis and applied this model to the case of microquasars.  The effects of strong gravity are significant in this region and the challenge for future techniques is to identify subtle, yet specific patterns in X-ray lightcurves. Polarization studies could help to achieve this goal. Distinct features should arise from multiple trajectories of light rays connecting the source with the observer along several different paths. The rays winding up around the photon circular orbit should experience a characteristic mutual time delay. We discuss the expected features in this paper.  Polarization of light from scattering in winds and jets was examined by various authors. Following the early papers (e.g.\\ \\cite{dol67}; \\cite{ang69}; \\cite{bon70}), Begelman \\& Sikora (\\yearcite{beg87}) studied the linear polarization of initially unpolarized soft radiation up-scattered by cold electrons in a jet. Beloborodov (\\yearcite{bel98}) examined the case of fast winds outflowing from an accretion disk slab. He found that polarization direction depends on the wind velocity; the terminal speed of the outflow plays a critical role. Poutanen (\\yearcite{pou94}) and Celotti \\& Matt (\\yearcite{cel94}) considered the synchrotron self-Compton mechanism, a likely process producing polarization wherever magnetic fields interact with relativistic particles. The effect of electron temperature was also discussed: it reduces the final magnitude of polarization. Recently, Lazzatti et al.\\ (\\yearcite{laz04}) further discussed the Compton drag as a conceivable mechanism for polarization in gamma-ray bursts. Hor\\'ak \\& Karas (\\yearcite{hor05}) studied the polarization of scattered light from a compact star, taking into account the light-bending effect. In fact, it was demonstrated that retro-lensing images, which clearly require strong gravity, can give rise to specific polarimetric signatures in predicted lightcurves. Here we develop this model further by considering non-negligible temperature of the scattering medium and by changing the geometry of the primary source. We assume a Keplerian disk as the source, so we are able to examine situations that are relevant for accreting black holes.  \\begin{figure}[tb!] \\begin{center} \\FigureFile(0.48\\textwidth,0.48\\textwidth){fig1.eps} \\end{center} \\caption{Geometry of the model. An accretion disk is the source of primary unpolarized light, which is then Thomson scattered on a cloud. The cloud moves radially along the axis, $z\\,\\equiv\\,z(t)$, in the radiation field of the disk. Gravity of the central black hole influences the motion of the cloud. The light rays of primary and scattered photons are also affected. Direct and indirect (retro-lensed) light rays have different degree of linear polarization and they experience different amplification and the Doppler boosting. The observed lightcurve is produced by all the rays reaching an observer at view angle $i$ far from the centre (along $Z$-direction). The indirect photons contribute  most significantly to the the total signal if the cloud moves toward the black hole and the observer inclination is small.} \\label{fig1} \\end{figure} ",
        "conclusions": "In X-rays, future polarimeters could be employed to probe jets and winds in strong gravitational fields of the central compact object. Polarimetry is a powerful tool that can provide additional information, which would be difficult to obtain by other techniques such as traditional photometry and spectroscopy. In this way  polarimetry helps to discriminate between different geometries and physical states of sources where accretion processes are accompanied by fast radial motion of the blobs of material. Naturally, this goal would require sufficient sensitivity in X-rays; the scattered signal is mixed with primary photons, which reduce the final polarization. We have seen that the predicted features are flashing for only a brief period of time and the maximum polarization degree is typically $\\simeq1$ percent or less.  In order to allow the analytical treatment we employed various simplifications: we considered the bolometric quantities and the Thomson cross-section for scattering of primary photons on electrons (rather than Compton scattering and the Klein-Nishina cross-section; see \\cite{mel89}; \\cite{ski94}; \\cite{mad00}). We also assumed that the flow is not magnetically dominated, although astrophysically realistic models require magnetohydrodynamic effects to be taken into account (\\cite{beg84}; \\cite{bes04}). Effects of general relativity were taken into account in the limit of the non-rotating black hole spacetime (we neglected the effects of frame-dragging for the sake of simplicity; cf.\\ \\cite{vok91}). Likewise, we adopted the simplest possible parameterization of the disk emissivity via the standard Shakura--Sunyaev model; this could be improved by including relativistic effects (\\cite{pag74}) and a more realistic description of the disc itself. These changes will be necessary to provide quantitative and astrophysically realistic results for the polarization degree, however, we do not expect any qualitative change regarding the signature of indirect photons. The predicted polarization is either parallel or perpendicular to the projection of cloud velocity onto the observing plane.  \\medskip We gratefully acknowledge fruitful discussions at the Institute  of Theoretical Physics in Prague and we thank for helpful comments that we have received from the referee. We also acknowledge the financial support from the Academy of Sciences (ref.\\ IAA\\,300030510) and from the Czech Science Foundation (ref.\\ 205/03/H144). The Astronomical Institute has been operated under the project AV0Z10030501."
    },
    "0512/astro-ph0512310_arXiv.txt": {
        "abstract": "{We present a wide grid of models for the structure and transmission properties of warm absorbers in active galactic nuclei (AGN). Contrary to commonly used constant density models, our absorbing cloud is assumed to be under {\\it constant total (gas plus radiation) pressure}. This assumption implies the coexistence of  material at different temperatures and ionization states, which is a natural consequence of pressure and thermal equilibrium. Our photoionization code allows us to compute the profiles of the density, the temperature, the gas pressure, the radiation pressure and the ionization state across the cloud, and to calculate  the radiative transfer of continuum and lines including Compton scattering. Therefore, equivalent widths of both saturated and unsaturated lines are properly modeled. For each pair of the incident spectrum slope and the ionization parameter at the cloud surface there is a natural upper limit to the total column densities of the cloud due to thermal instabilities. These maximum values are comparable to the observational constraints on the column density of warm absorbers which may give support to constant total pressure models. In all models we note considerable absorption around 6.4 keV which modifies the intrinsic relativistically broadened iron line profile originating in an accretion disk illuminated atmosphere. Our models can be applied to fitting the spectroscopic data from the {\\it XMM-Newton} and {\\it Chandra} satellites. ",
        "introduction": "The first X-ray absorption feature due to ionized heavy elements was recognized by Halpern (1984) in the X-ray spectrum of the Sy1 galaxy MR 2251-178 observed by the {\\it EINSTEIN} satellite. The author related the observed jump  in flux around 1 keV to the absorption edge of O{\\sc viii}. Therefore, as Halpern concluded, X-rays emitted from the central region of an active galaxy traveling toward an observer encounter a ``warm absorber'' --- material with  an electron temperature lower than the temperature of collisionally ionized gas with a similar level of ionization.  Results from {\\it EXOSAT}, {\\it ROSAT}, {\\it GINGA}, {\\it ASCA}, and {\\it Beppo-SAX} satellites showed that warm absorbers were common among Seyfert galaxies. Nandra \\& Pounds (1994) have suggested that more than 50\\% of Sy1s contain a warm absorber, and later results confirmed such occurrence rate (e.g. Reynolds, 1997; George at al. 1998). Those conclusions were based exclusively on the detection of absorption edges (for instance in MCG-6-30-15, Nandra \\& Pounds 1992). The detector areas and the spectral resolution of those satellites did not allow us to see any absorption lines from  highly ionized species.  The situation changed since 1999, when large X-ray telescopes {\\it Chandra} and {\\it XMM-Newton} started to operate, with their X-ray grating instruments working in the energy range up to almost 10 keV. Several tens of absorption or emission lines were observed and identified in NGC 3783 (Kaspi el al. 2002, Behar el al. 2003, Netzer et al. 2003, Krongold el al. 2003), NGC 5548 (Kaastra el al. 2002), NGC 1068 (Kinkhabwala et al. 2002), NGC 7469 (Blustin et al. 2003), MCG-6-30-15 (Turner et al. 2004). In other objects the number of fitted lines is lower and/or their detection is less firm, but the  results still strongly support the presence of a warm absorber in those Seyferts (NGC 4051, Collinge et al. 2001; Mrk 509, Yaqoob el al. 2003; TonS180, \\Agata el al. 2004). Warm absorber lines were even detected in a distant blazar at $z = 4.4$ (Worsley et al. 2004).  Detected lines are basically consistent with the unification scheme of AGN based on the presence of the dusty/molecular torus. According to that scheme, Sy1 are objects seen face-on and Sy2 are objects seen edge-on (Antonucci \\& Miller 1985; for application to X-ray band see Mushotzky, Done \\& Pounds 1993). The major characteristic of the X-ray spectra of Sy2 galaxies is that they are strongly absorbed at low energies, and that their emission lines, especially the iron K$\\alpha$ line, have large equivalent widths (Turner et al. 1997) since their intensity is measured with respect to the heavily obscured direct continuum (Weaver \\& Reynolds 1998) or only with respect to the scattered continuum (in Compton-thick objects), as discussed by Bassani et al. (1999). In some cases of Compton-thin Sy2 obscuration may be due to the host galaxy and unrelated to disk/torus orientation (Matt 2000, Guainazzi et al. 2001) The spectra of Sy1 galaxies are predominantly featureless, show little low energy absorption, and the equivalent width of the iron K$\\alpha$ line is much lower than in Sy2. Some Sy1 show relativistically broadened iron K$\\alpha$ line profiles which additionally supports the view that in Sy1 we directly observe the innermost part of the nucleus. The scattering medium may be identical to the warm absorber and we have side view of this medium in Sy2 galaxies while we see the nucleus through it in Sy1. Indeed, X-ray spectra of some Sy2 galaxies show many narrow emission lines, (e.g. NGC 1068, Kinkhabwala et al. 2002; NGC 4507, Matt et al. 2004; Mkn 3, Pounds \\& Page 2005; NGC 4151, Schurch et al. 2004) indicating that ionized material may extend beyond the shielding torus. In Sy1 spectra mostly absorption lines are observed (e.g. NGC 3783, Kaspi et al. 2002: NGC 5548, Kaastra et al. 2002; Mrk 509, Yaqoob et al. 2003) without strong absorption of continuum, suggesting that in those objects we see directly the nucleus through the ionized plasma.  Some contribution of emission lines in Sy1 X-ray spectra is also expected (e.g. Netzer 1993, Collin, Dumont \\& Godet 2004) and seen in the data (e.g. NGC 3783, Behar et al. 2003; NGC 7469, Scott et al. 2005). However, the constraints from Sy2 galaxies for the ionized medium cannot be used directly to Sy1 galaxies since in Sy2 galaxies we may observe only the outer part of the plasma distribution, due to the obscuration by the torus, while warm absorber features may come predominantly from the inner part.  A majority of these absorption lines shows a velocity shift of the order of a few hundreds km/s (Kaspi et al. 2001, Kaastra et al. 2002) suggesting that the warm absorber is outflowing. A strong high velocity outflow was recently reported in several objects, mostly radio-quiet quasars (e.g. 60000 km/s and 120000 km/s in APM 08279+5255, Chartas et al. 2002; 23000 km/s in PG1211+143, Pounds et al. 2003a; 63000 km/s in PG0844+349, Pounds et al. 2003b, 30000 km/s in IRAS 13197-1627, Dadina \\& Cappi 2004; 26,000 km/s in PG 1404+226, Dasgupta et al. 2005). However, the results are based on a few detected lines, so line identifications and, consequently, the determined high outflow velocities can be questioned (Kaspi 2004a). The origin of the outflow and its geometry is still under discussion (Crenshaw, Kraemer \\& George 2003a, Blustin et al. 2005).  It is more difficult to estimate the radial distance of the warm absorber from the nucleus. Absorption lines most probably form somewhere between the broad line region (BLR) and   the narrow line region (NLR) (i.e. from about 0.01-0.1 up to 10 pc from nucleus; see Crenshaw et al. 2003a, Blustin et al. 2005). Variability in the overall warm absorber properties was reported for a few sources (MR 2251-178, Halpern 1984, Kaspi et al. 2004b; MCG -6-30-15, Reynolds et al. 1995; H1419+480, Barcons et al. 2003; NGC 4395, Shih et al. 2003; NGC 3516, Netzer et al. 2002). Netzer et al. (2002) and Barcons et al. (2003) concluded that the changes observed are consistent with varying ionization of the gas. The lack of short-timescale (days) response of the warm absorber to the change of the continuum can be used to put lower limits to the warm absorber distance (e.g. 0.5 - 2.8 pc for NGC 3783, Behar et al. 2003; similar limits were given by Netzer et al. 2003). Krongold et al. (2005) detected a response of the warm absorber to the changes of the continuum on a timescale of 31 days in NGC 3783, deriving an upper limit of 6 pc for the distance of the warm absorber. The shortest variability timescale of $\\sim 10^4$ s has been detected in the warm absorber in MCG -6-30-15 (Otani et al. 1996; see also Turner et al. 2004) locating the plasma responsible for the O{\\sc viii} edge within the distance of $10^{17}$ cm from the nucleus. The spectral analysis of this source indicates that the highly ionized warm absorber is dust-free, with dust contribution in this source coming from a distant zone, hundreds of pc from the nucleus (Ballantyne et al. 2003).  The same medium is most probably responsible for narrow absorption lines seen in the UV spectra of many AGN (for a review, see Crenshaw et al. 2003a). Some kinematic components discovered in UV coincide with those discovered in soft X-rays but for other components no such correspondence is seen (e.g. Behar et al. 2003 for NGC 3783, Crenshaw et al. 2003b for NGC 5548, Kaspi et al. 2004b for MR 2251-178; Scott et al. 2005 for NGC 7469; Gabel et al. 2005 for NGC 3783). However, the resolving power of the UV observations is high ($R \\sim 20 000$) while X-ray data resolution is much lower ($R \\sim 1 000$), which makes the comparison difficult, as discussed by Crenshaw et al. (2003a). Analysis usually suggests that UV and X-ray absorption features are consistent with arising in the same gas, but with stratified ionization (e.g. Barcons et al. 2003, Kaspi et al. 2004b, Scott et al. 2005, Gabel et al. 2005). Since the absorption features appear in the profiles of the broad emission lines like C{\\sc iv}, and are occasionally deep, this serves as an argument that the absorbing region is located outside the BLR.  The column density of the warm absorber is generally estimated to be about $10^{21-23}$ cm$^{-2}$, and absorbing gas contains heavy elements mostly in the form of helium- and hydrogen-like ions. However, accurate measurements of the column density are quite complex. Most determinations are based on detection of absorption edges, but in some cases edges are undetectable while lines are clearly seen (Kaastra el al. 2002, \\Agata et al. 2004). Also the estimates of column densities from edges do not always confirm estimates derived from the absorption line analysis (Kaspi et al. 2002).  The ionization state of the gas required to reproduce the observed absorption or emission lines seems to be quite complex. In many objects predictions based on a single cloud at one specific ionization state cannot explain the data, so an absorbing material is modeled using at least two photoionization regions, which  are required to explain presence of lines from the matter in different ionization states (see in the case of Sy1:  NGC 4051, Collinge et al. 2001; NGC 5548,  Kaastra et al. 2002;  NGC 3783,  Kaspi et al. 2002, Netzer et al. 2003 and Krongold et al. 2005, H0557-385, Ashton et al. 2005). The same conclusion was drawn from fitting of the absorbed continuum (strongly absorbed Sy1: Mkn 304, Piconcelli et al. 2004; IC 4329A, Steenbrugge et al. 2005a,b; in the case of a dwarf galaxy with an active nucleus: NGC 4395, Shih, Iwasawa \\& Fabian 2003; and in case of Sy2: NGC 4507, Matt et al. 2004).   The idea of the warm absorber being under the constant pressure was developed even before {\\it Chandra} and {\\it XMM-Newton} satellites were lunched  (Netzer 1993; Krolik \\& Kriss 1995; Netzer 1996). It is well known that cold/warm material irradiated by hard  X-rays should be strongly stratified, and eventual thermal instabilities lead to its clumping (McKee Krolik \\& Tarter 1981). If thermal instabilities are strong, none of the existing photoionization codes  can describe unstable zone. This is because radiative transfer codes are based on unique density and temperature profiles, and do not accept situations where for one value of  optical depth the solution gives three different values of temperatures and densities.  Therefore many models of the warm absorber being under constant pressure do assume that those two discrete phases already exist. Usually two or three zones at different constant densities are assumed to have the same dynamical ionization parameter, $\\Xi$, which is the ratio of ionization pressure to the gas pressure (McKee Krolik \\& Tarter 1981). In such situation, calculations of transfer of X-ray radiation through two constant density zones are done separately and then spectrum is merged together depending on covering factor  (Netzer 1993).  In this paper we solve the situation where X-ray radiation is no so hard and strong, that separation on two phases takes place. Instead, our warm absorber is strongly stratified and radiation passes through different densities and ionization stages. Theoretically this situation was considered by Krolik \\& Kriss (2001), Krolik (2002). The stratification of our warm absorber is determined physically by radiative properties, and the only assumption which we make is constant total pressure within a cloud. We use the photoionization code {\\sc titan} developed by Dumont et al. (2000) (see Dumont et al. 2003 for implementation of Accelerated Lambda Iteration method) to compute the synthetic spectra for a systematic set of model parameters. The advantage of our calculations is that we compute the full radiative transfer of continuum and lines taking care also of lines which are  saturated. Recent observations suggest that saturated lines are often present in warm absorber (Kaspi et al. 2002, \\Agata et al.2004). The transfer is done in the stratified medium, and the density profile is determined self-consistently with the radiation transfer to fulfill the condition of the pressure equilibrium across the cloud. We aim understanding these AGN which show clear absorption lines in their {\\it Chandra} or {\\it XMM} spectra so the warm absorber is located in the line of sight to the observer. We concentrate on modeling the transmission spectra and absorption lines therefore our models can be used for Sy1 galaxies.  We describe our model in Sec.~\\ref{sect:model} and ~\\ref{sec:grid}, while results are presented in Sec.~\\ref{sect:struc}, ~\\ref{sect:local_spot}, and \\ref{sec:rel}. Sec.~\\ref{sec:comp} describes comparison of our models to observations and  Sec.~\\ref{sec:disc} contains conclusion remarks. ",
        "conclusions": "\\label{sec:disc}  We have calculated a grid of models of warm absorber with density stratification determined by the condition of the constant pressure. The assumption of constant pressure instead of constant density was motivated by recent observations of AGN. In most observed objects the ionization states implied by the observed lines span a large range, which cannot be accounted for a single photoionized region, or by collisionally ionized matter. At least two photoionized absorbing regions are required to fit the data and the properties of these two regions are consistent with pressure equilibrium  Our models represent an essential improvement in the description of a single cloud over the models presented in the literature. All previous models adopted constant density approximation for the medium (Netzer 1993, Kaastra et al. 2002, Kinkhabwala et al. 2002, Krongold et al. 2005), or even constant temperature  (Krolik \\& Kriss 1995).  On the other hand, we do not consider here a global picture including the dynamics of the warm absorber flow which form a separate and a major issue (see Chelouche \\& Netzer 2005 and the references therein).  We have calculated the full radiative transfer of the illuminating radiation in continuum and in lines through the plane-parallel density-stratified slabs of different total column densities. For a low total column density the cloud is mostly ionized and optically thin, and does not differ from the constant density model. But for a higher total column density the optically thick dense zone arises at the back of the illuminated cloud, and the temperature falls dramatically. Therefore, absorbing matter contains zones of different ionization states coexisting under constant pressure.  For each set of models with spectral index $\\Gamma$ and surface ionization parameter $\\xi$ there is a maximum total column density $N^{Max}_{H_{tot}}$ for which the cloud is thermally stable. For higher column densities thermal instabilities do not allow the calculations to converge. Interestingly, $N^{Max}_{H_{tot}}$  is of the order of the maximum column density derived from X-ray observations of different AGN. It gives additional support to the idea of the warm absorber clouds being in pressure equilibrium.  The modeled lines are usually optically thick and their equivalent widths are of the order of the observed values. We conclude that the observed warm absorbers at column densities of the order of $10^{22}$ cm$^{-2}$ and higher possess saturated absorption lines and that full radiative transfer is required to model their equivalent widths properly. For full evaluation  of this problem we will perform a curve of growth analysis in our future work.  Generally lines are easier to detect than ionization edges and this tendency is observed in several AGN (Kaastra et al. 2002, \\Agata et al. 2004). Our models predict that edges are not observable up to column densities about $10^{22}$ cm$^{-2}$ and even higher, depending on $\\Gamma$ and $\\xi$.  The most interesting result of our computations is the shape of the spectrum around an iron K$\\alpha$ line. In almost all models there are strong and narrow absorption features due to highly ionized iron ions, above Fe{\\sc xvii}. Such absorption affects the shape of the broad iron emission line possibly originating in illuminated disk atmosphere. For instance, relativistic broad iron line profile after passing through warm absorber becomes disrupted into three narrower profiles, which can be fitted by Gaussian shape. Such lines were reported in several AGN as presented by Yaqoob \\& Padmanabhan (2004) (see also NGC 3783  Reeves et al. 2004).  In this paper we consider only single cloud but observations suggest that absorbing material forms distribution of clouds with covering factor which may depend on velocity of clouds (de Kool, Korista \\& Arav 2002). However, theoretical modeling of cloud distribution with complex velocity field is complicated. Also, the study of the transmission spectra, including absorption lines, do not give us any diagnostic of the matter along other directions than the line of sight. The {\\sc titan} code can be used, however, to study the radiative transfer in all directions. We address both issues to the future work."
    },
    "0512/astro-ph0512098.txt": {
        "abstract": "{We present an in-depth near-IR analysis of a sample of H$_2$ outflows from young embedded sources to compare the physical properties and cooling mechanisms of the different flows. The sample comprises 23 outflows driven by Class 0 and I sources having low-intermediate luminosity. We have obtained narrow band images in H$_2$~2.12\\,$\\mu$m and [\\ion{Fe}{ii}]~1.64\\,$\\mu$m and spectroscopic observations in the range 1-2.5\\,$\\mu$m. From [\\ion{Fe}{ii}] images we detected spots of ionized gas in $\\sim$74\\% of the outflows which in some cases indicate the presence of embedded HH-like objects. H$_2$ line ratios have been used to estimate the visual extinction and average temperature of the molecular gas. $A_{\\rm v}$ values range from $\\sim$2 to $\\sim$15 mag; average temperatures range between $\\sim$2000 and $\\sim$4000\\,K. In several knots, however, a stratification of temperatures is found with maximum values up to 5000\\,K. Such a stratification is more commonly observed in those knots which also show [\\ion{Fe}{ii}] emission, while a thermalized gas at a single temperature is generally found in knots emitting only in molecular lines. Combining narrow band imaging (H$_2$, 2.12\\,$\\mu$m and [\\ion{Fe}{ii}], 1.64\\,$\\mu$m) with the parameters derived from the spectroscopic analysis, we are able to measure the total luminosity of the H$_2$ and [\\ion{Fe}{ii}] shocked regions ($L_{H_2}$ and $L_{[\\ion{Fe}{ii}]}$) in each flow. H$_2$ is the major NIR coolant with an average $L_{H_2}$/$L_{[\\ion{Fe}{ii}]}$ ratio of $\\sim$10$^{2}$. We find that $\\sim$83\\% of the sources have a $L_{H_2}$/$L_{bol}$ ratio $\\sim$0.04, irrespective of the Class of the driving source, while a smaller group of sources (mostly Class I) have $L_{H_2}$/$L_{bol}$ an order of magnitude smaller. Such a separation reveals the non-homogeneous behaviour of Class I, where sources with very different outflow activity can be found. This is consistent with other studies showing that among Class I one can find objects with different accretion properties, and it demonstrates that the H$_2$ power in the jet can be a powerful tool to identify the most active sources among the objects of this class. ",
        "introduction": "Matter flows emitted from young stellar objects (YSOs) are manifestations commonly observed during all the phases of Pre-Main Sequence evolution: from the early accretion phase (Class 0), which lasts a relatively short time ($\\sim10^4$~yrs, for a low mass YSO), to the final contraction toward the Main Sequence (Class II and III, $\\sim10^7$~yrs). According to low mass star formation models, mass accretion and ejection rates ($\\dot{M_{acc}}$, $\\dot{M_{out}}$) are expected to be strictly related, because a significant fraction of the infalling material is ejected by the accretion disk. The efficiency of such a coupling, however,  has to necessarily decrease while the evolution proceeds. Therefore the outflow/jet properties are expected to remarkably change with time and their study offers both direct and indirect clues to understand the processes related to protostellar evolution. In addition, several issues can be addressed by systematic observations of jets since: ({\\it i}) they signal the presence of strongly embedded driving sources, often invisible up to far IR wavelengths; ({\\it ii}) their shape may help to reconstruct the star/disk rotation; ({\\it iii}) they have dynamical and chemical effects on the closeby interstellar material; ({\\it iv}) at the same time, they are also influenced by the properties of such material.  Since bipolar and collimated jets are easily observed from the ground over a wide range of frequencies, large imaging and spectroscopical data-bases have been accumulated by various groups. Our group, in particular, has gathered, over the last five years, a large spectroscopical data set in the near IR (1--2.5 $\\mu$m), complemented with a considerable amount of high-resolution imaging data ([\\ion{Fe}{ii}], H$_2$), on a sample of active H$_2$ outflows from both Class 0 and Class I sources. The material collected so far has allowed us to clarify specific aspects of the protostellar jets physics, such as the identification of the crucial spectral range for the study of the H$_2$ excitation conditions \\citep{paper3}; the role of [\\ion{Fe}{ii}] as a diagnostic tool of embedded atomic jets \\citep{paper4}; the observational tests of state-of-art shock models \\citep{paper5,mcoey,paper8}. Moreover, near IR spectroscopy allowed us to investigate the properties of H$_2$ jets in individual star-forming regions \\citep[IC1396;][]{nisini1}, \\citep[Vela Molecular Ridge;][]{paper6,paper1,paper7,loren}. Here we use our data-base to derive more general properties on the IR activity of jets from young stars. This has been done complementing our published data with new observations focused mainly on the youngest Class 0 sources.  The main aim of this study is to reveal any systematic difference in the derived physical parameters which can be attributed to an effect of the evolution in both the intrinsic jet properties and in the way the jet interacts with the ambient medium. In addition, we want to define if any relationship exists between the total IR cooling of the outflows and the evolutionary status of the driving source. Indeed, several authors \\citep[see e.\\,g.][]{cabrit0,bontemps,saraceno,andre,paper2} have observationally shown that during the first stages of protostellar evolution  a correlation exists between different tracers of the outflow activity and the source bolometric luminosity, which is believed to be largely dominated by the accretion luminosity ($L_{bol} \\approx L_{acc} = GM_{*} \\dot{M_{in}} / R_{*}$).  In the near IR the jet shocked excited regions mainly cool through H$_2$ quadrupole transitions, therefore the bright ro-vibrational lines in that range represent a suitable shock tracer which can be used to evaluate the molecular hydrogen luminosity ($L_{H_2}$) of the outflow. Due to its very rapid cooling, H$_2$ is more suitable to probe shock excitation and gas cooling than the CO lines which can only give a time-integrated response \\citep[see e.\\,g.][]{smith1}. The empirical classification of early protostellar evolution by means of the H$_2$ luminosity of the emitted jets has recently attracted much attention in different works \\citep{stanke,smith1,froebrich1,oco}, where the simplified assumption was adopted that the total H$_2$ luminosity is derivable from the 2.122\\,$\\mu$m (1-0S(1)) line luminosity. In the following, the validity of such an assumption as well as the most critical aspects in deriving a reliable value of $L_{H_2}$ will be also reviewed.  The structure of the paper is the following: in Sect.2 we define the investigated sample of jets and associated YSOs with their parameters; in Sect.3 our NIR observations are presented; in Sect.4 we derive, for each jet, the physical quantities from the molecular and ionic components; in Sect.5 a discussion is presented in terms of jet properties vs. luminosity of the central object. Our conclusions are summarized in Sect.6. Finally, in Sect.7 (Appendix) the detailed results on any individual star forming region are reported.    %%%%TABLE 1 THE SAMPLE%%%%%%%%%%%%%%%%%%%% \\begin{table*} \\caption[]{ The investigated sample \\label{sample:tab}} \\begin{center} \\begin{tabular}{ccccccccccccc} \\hline \\hline\\\\[-5pt] Id & Source &  Associated &      \\multicolumn{3}{c}{$\\alpha$(2000.0)} & \\multicolumn{3}{c}{$\\delta$(2000.0)} &  Class & $L_{bol}$ & Ref.  & D   \\\\ &        &   HH        &   ($^{h}$ & $^{m}$ & $^{s}$) & ($\\degr$ & $\\arcmin$ & $\\arcsec$)        &   & ($L_{\\sun}$) & & (pc) \\\\ \\hline\\\\[-5pt]  1 & L1448-MM       \t     & -\t    \t &03&25&38.8&30&44&05.0  & 0   & 8.3--9 & 1,2 & 300  \\\\ 2 & NGC1333-I4A    \t     & -\t    \t &03&29&10.5&31&13&30.5  & 0   & 14--18 & 2,1 & 350  \\\\ 3 & HH211-MM       \t     & HH211 *\t    \t &03&43&56.8&32&00&50.0  & 0   & 3.6--5 & 1,2 & 300   \\\\ 4 & IRAS05173-0555 \t     & HH240/1 *  \t     &05&19&48.9&-05&52&05.0 & 0/I & 17--26.6& 3,4 & 460    \\\\ 5 & HH43-MMS       \t     & HH43, HH38, HH64      &05&37&57.5&-07&07&00.0 & 0   & 3.6--5  & 3,4 & 450\t     \\\\ 6 & HH212-MM       \t     & HH212 *\t\t     &05&43&51.5&-01&02&52.0 & 0   & 7.7--14 & 1,2  & 400  \\\\ 7 & HH26IR         \t     & HH26 *\t\t     &05&46&03.9&-00&14&52.0 & I   & 28.8   & 5 & 450\t \\\\ 8 & HH25-MMS       \t     & HH25 *\t\t     &05&46&07.8&-00&13&41.0 & 0   & 6--7.2 & 6,1 & 450  \\\\ 9 & HH24-MMS       \t     & HH24 *\t\t     &05&46&08.3&-00&10&42.0 & 0 &  5 & 7 & 450 \t  \\\\ 10 & IRAS05491+0247 (VLA2)   & HH111, HH311, HH113 *   &05&51&46.1&02&48&30.6 & I & 24--42 & 1,4,8 & 450  \t\t \\\\ 11 & NGC2264G-VLA2  \t     & - *\t\t     &06&41&11.0&09&55&59.2  & 0 & 12--13 & 9,1  & 800    \\\\ 12 & IRAS07180-2356 \t     & HH72 *\t\t     &07&20&10.3&-24&02&24.0 & I & 170 & 3 & 1500\t \\\\ 13 & IRAS08076-3556 \t     & HH120\t\t     &08&09&32.8&-36&05&00.0 & I &  13--19 & 3,15 & 450    \\\\ 14 & IRAS08211-4158 (IRS8-2) & HH219 *\t\t     &08&22&52.1&-42&07&55.0 & I & 642& 10  & 400\t  \\\\ 15 & \\#40-3 (IRS17) \t     & - *\t\t     &08&46&32.6&-43&54&38.9  & I   & 11--245 & 11 & 700  \\\\ 16 & BHR71-MM (IRS1)\t     & HH321\t\t     &12&01&44.0&-65&09&00.1 & 0   & 7.9--10& 2,18 & 200    \\\\ 17 & BHR71 (IRS2)   \t     & HH320\t\t     &12&01&34.0&-65&08&44.0 & I  & 1--3 & 2,18 & 200\t  \\\\ 18 & IRAS12515-7641 \t     & HH54\t\t     &12&55&00.2&-76&57&00.0 & I & 0.22--0.44 & 16 & 180  \\\\ 19 & VLA1623-243    \t     & HH313 *\t\t     &16&26&26.5&-24&24&31.0 & 0 & 1 & 2   & 160\t   \\\\ 20 & IRAS18273+0113  \t     & HH460\t\t     &18&29&49.8&01&15&20.8  & 0/I & 45--72 & 1,12  & 310  \\\\ 21 & R CrA-IRS7     \t     & HH99\t\t     &19&01&55.3&-36&56&21.9  & I & 3.4 & 17,19 & 130\t  \\\\ 22 & L1157-MM \t\t     & - *\t\t     &20&39&05.7&68&02&16.0  & 0   & 8.4--11& 1,2 & 440   \\\\ 23 & IRAS21391+5802 \t     & HH593 *\t\t     &21&40&42.4&58&16&09.7  & 0   &150--350 & 13,14 & 750  \\\\ \\hline \\hline \\end{tabular} \\end{center} {\\bf References}: {\\bf (1)} \\citet{froeb}, {\\bf (2)} \\citet{andre}, {\\bf (3)} \\citet{reipurth}, {\\bf (4)} \\citet{molinari}, {\\bf (5)} \\citet{davis1}, {\\bf (6)} \\citet{lis}, {\\bf (7)} \\citet{chini1}, {\\bf (8)} \\citet{chini2}, {\\bf (9)} \\citet{ward}, {\\bf (10)} \\citet{paper6}, {\\bf (11)} \\citet{paper7}, {\\bf (12)} \\citet{larsson}, {\\bf (13)} \\citet{beltran}, {\\bf (14)} \\citet{nisini1}, {\\bf (15)} \\citet{persi}, {\\bf (16)} \\citet{Cohen}, {\\bf (17)} \\citet{marraco}, {\\bf (18)} \\citet{bourke}, {\\bf (19)} \\citet{wilking0} \\\\ Note: asterisks in column 3 indicate outlows with a defined morphology, where the knot assignation taken from literature or derived from this work (see Appendix) is certain.\\\\ \\end{table*}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We have measured the cooling and the physical properties at NIR wavelengths of a sample of protostellar jets (23 objects), originated by Class 0 and I low-intermediate solar mass YSOs, presenting new spectroscopic and imaging observations of 15 objects. We have investigated how the derived properties are correlated with the evolution of the jets and their exciting sources. The main results of this work can be summarized as follows:  \\begin{itemize} % \\item[-] [\\ion{Fe}{ii}] emission in active H$_2$ jets has been systematically investigated to define the occurrence of embedded ionized gas in young outflows. [\\ion{Fe}{ii}] emission spots are observed in $\\sim$74\\% of the investigated sample, irrespective of the driving source class. This indicates that dissociative-shocks are common even in high-density embedded regions which do not show optical HH objects.  \\item[-] H$_2$ line ratios have been used to estimate the visual extinction ($A_{\\rm v}$) and average temperature of the molecular gas. When observed, the [\\ion{Fe}{ii}] 1.53/1.64\\,$\\mu$m ratio was used to determine the electron density ($n_{\\rm e}$) of the atomic gas component. $A_{\\rm v}$ values range from $\\sim$2 to $\\sim$15 without any evidence of higher extinction associated with the Class 0 flows. The H$_2$ average gas temperatures range between $\\sim$2000 and 4000\\,K. In several knots, however, a stratification of temperatures is found with maximum values up to 5000\\,K. Generally, gas components at different temperatures are associated with the knots also showing [\\ion{Fe}{ii}] emission while thermalized gas at a single temperature is most commonly found in knots emitting only in molecular lines.  % \\item[-] We have computed the total cooling due to H$_2$ and [\\ion{Fe}{ii}] ($L_{H_2}$, $L_{FeII}$) adopting the parameters derived from the line ratio analysis in the single knots and the H$_2$ 2.12\\,$\\mu$m, [\\ion{Fe}{ii}] 1.64\\,$\\mu$m luminosities derived from the narrow band imaging. The determination of $L_{H_2}$ strongly depends on the local gas conditions and in particular on the $A_{\\rm v}$ value, therefore the often-used approximation $L_{2.12} \\sim 0.1 \\times L_{H_2}$ can be wrong by up to an order of magnitude if a proper reddening is not applyed. % \\item[-] By comparing the measured outflow H$_2$ luminosity with the source bolometric luminosity (assumed representative of the accretion luminosity), we find that for $\\sim$83\\% of the sources there is a correlation between these two quantities, with $L_{H_2}$/$L_{bol}$$\\sim$0.04. A small sample of four sources, however, display an efficiency $L_{H_2}$/$L_{bol}$ lower by about an order of magnitude. We interpret this behaviour in terms of evolution, with the sources which are less efficient H$_2$ emitters are more evolved than the others.  \\item[-] We also find that there is not a clear separation in terms of $L_{H_2}$/$L_{bol}$ efficiency between Class 0 and Class I sources (although the four objects with the lower $L_{H_2}$/$L_{bol}$ value are all from Class I or intermediate Class 0/I). This partially reflects the large heterogeneity between the evolutionary properties among Class I, which include sources with very different accretion properties. In addition, the efficiency of the H$_2$ cooling in a dense ambient medium, likely characterizing the Class 0 environment, can be limited by the relatively small critical density of H$_2$ IR lines ($\\sim$10$^3$-10$^4$\\,cm$^{-3}$). Indeed, the total cooling of flows from Class 0 sources prevalently occurs at far IR wavelengths through CO and H$_2$O emission lines and thus can be underestimated considering only the H$_2$ contribution. % \\item[-] On the basis of the observational experience, accumulated over the last decade, it maybe empirically adequate to refine the concept of Class I objects. According to the present analysis of the H$_2$ jets, we suggest to define a new Class 0.5, which is composed (as a starting point) of objects defined as Class I YSOs \\citep[according to][]{lada}, but presenting a $L_{H_2}$/$L_{bol}$ ratio similar to that of Class 0 sources (i.\\,e. $\\sim$0.04 or higher). We would classify as Class 0.5 the following objects: $IRAS05173-0555$, HH26IR, $IRAS07180-2356$, $IRAS08076-3556$, \\#40-3 (IRS17), BHR71(IRS2), $IRAS12515-7641$. Obviously, this tentative classification should be confirmed with further observations.  \\end{itemize}"
    },
    "0512/astro-ph0512476_arXiv.txt": {
        "abstract": "We present variability and multi-wavelength photometric information for the 933 known quasars in the QUEST Variability Survey. These quasars are grouped into variable and non-variable populations based on measured variability confidence levels. In a time-limited synoptic survey, we detect an anti-correlation between redshift and the likelihood of variability. Our comparison of variability likelihood to radio, IR, and X-ray data is consistent with earlier quasar studies. Using already-known quasars as a template, we introduce a light curve morphology algorithm that provides an efficient method for discriminating variable quasars from periodic variable objects in the absence of spectroscopic information. The establishment of statistically robust trends and efficient, non-spectroscopic selection algorithms will aid in quasar identification and categorization in upcoming massive synoptic surveys. Finally, we report on three interesting variable quasars, including variability confirmation of the BL Lac candidate PKS 1222+037. ",
        "introduction": "A considerable amount of analysis has been carried out on various aspects of quasar variability \\citep[see][for a review of past surveys]{Helfand01}. For example, \\citet{Cristiani96} merged variability results from quasars in three separate photographic plate fields: SA 57 \\citep{Trevese94}, SA 94 \\citep{Cristiani96}, and the South Galactic Pole \\citep{Hook94} to study ensemble and individual quasar variability for several hundred quasars. Recent work \\citep{Vanden04,Devries05} has utilized the large number of quasars discovered by the Sloan Digital Sky Survey (SDSS) to make statements on quasar variability on an ensemble basis for $\\sim\\!\\!10^4$ quasars with two or three data points per quasar light curve. In summary, most previous variability studies focus either on a large number of epochs for a relatively small sample of quasars, or a small number of epochs for a large sample of quasars. %  In this paper, we take a somewhat different approach. The QUEST Variability Survey \\citep[QVS:][hereafter, R04b]{Rengstorf04b} contains light curves in up to four bandpasses for nearly 200,000 objects. The QVS contains 69 scans of a thin strip of high Galactic latitude equatorial sky ($-1\\fdg2 \\leq \\delta \\leq 0\\fdg2$; $10^h \\leq \\alpha \\leq 15^h30^m$) collected between 1999 February and 2001 April. Four broadband filters (camera filter order: $RBRV$) were used throughout the variability scans. Typical seeing at the site was about $2\\farcs8$ and a limiting magnitude of $r = 19.06$ was reached. The filter set and observing cadence were chosen to optimize between multiple variability-driven projects, including a SNe 1a search and a recently published RR Lyrae catalog \\citep{Vivas04}. The synoptic study of $\\sim\\!\\!10^5$ light curves with several observations per lunation for several lunations per year for several years will serve as a testbed for algorithms and selection techniques for upcoming massive synoptic surveys (e.g., LSST, JDEM, Pan-STARRS). We see a need to develop robust techniques that bridge the gap between working with well-sampled individual quasars (e.g., 3C 273) and working with $10^5 - 10^6$ objects.  We combine the synoptic data from the QVS with published spectroscopic quasar catalogs to study the variability of confirmed quasars. We use the Global Confidence Level (GCL) parameter (see R04b for complete details) to construct populations of variable and non-variable quasars. This technique identifies variable objects based on the statistical likelihood that the object is reliably varying above the photometric noise of the entire ensemble. An object's GCL value is the weighted average of confidence levels for every bandpass in which we have data for that object. In using the likelihood of variability rather than the amplitude of variability to identify variable objects, we account for increased photometric noise near the magnitude limit of the QVS. Using a critical GCL value for the determination of variability allows easy fine-tuning for completeness and efficiency based on the particular needs of the study.  Extending our analysis to other wavelengths, we match the QVS synoptic data of the known quasars to several large-area, non-optical catalogs. First, we identify differences in global variability statistics based on quasar rest-frame time baseline and known quasar detection/non-detection in non-optical (Radio, IR, and X-ray) catalogs. Second, we look at the redshift and multi-wavelength luminosities of variable and non-variable populations of quasars.  Finally, we develop a light curve morphology test that cleanly delineates highly variable objects into periodic and aperiodic sources. We apply this novel technique to the data from the QVS, using spectroscopically confirmed variable objects from the SDSS Third Public Data Release \\citep[DR3;][]{Abazajian05} as our training sample. Such a morphology test can improve quasar-star separation among variable objects, aiding, for example, in identifying false positives in serendipitous quasar lens searches \\citep[see, e.g.,][]{Pindor05}.  Throughout this paper, we assume the  standard WMAP cosmology \\citep{Spergel03},  $\\Omega = 1$, $\\Omega_{\\Lambda} = 0.73$, and H$_o = 71$  km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have used a unique, synoptic dataset, federated with large-scale optical and non-optical public survey data, to explore quasar variability in a realm of phase space not previously examined. We have larger numbers of light curve points than recent work using SDSS data and we are spanning more area than earlier pointed observations. Using data from the QVS and known quasars from the SDSS, 2QZ, and VC03 catalogs, we have assembled 933 quasar light curves in multiple bandpasses. Using the GCL parameter from R04b, we defined populations of variable and non-variable quasars. Converting the light curves to the quasars' rest frames, we find a bias towards nearby objects being more likely to vary in a time-limited synoptic survey, likely due to luminosity and time lag effects. This trend will need to be accounted for in the upcoming synoptic surveys.  The anti-correlation between variability and redshift is evidence that we are sampling time lags in which quasar variability is still rapidly increasing with time lag. Previous work using ensemble structure function analysis has shown that quasar variability is expected to increase with time lag over at least several years \\citep[e.g.,][]{Hook94,Cristiani97,Vanden04,Devries05}. Recent work has indicated a positive correlation between redshift and variability, decoupled from the correlation between variability and photon wavelength \\citep{Vanden04}. We did not see the redshift-variability correlation in our data, which is not surprising given our sampled time baselines and the reported weakness of the redshift-variabilty correlation.  A near-UV (2500 \\AA) luminosity density was calculated for every known quasar. The median luminosity densities for the variable and non-variable quasar populations showed an anti-correlation with likelihood of variability, consistent with earlier work \\citep[e.g.,][]{Hook94,Trevese94,Cristiani96,Vanden04,Devries05}. The known quasars were also matched to the FIRST, 2MASS, and ROSAT surveys. Correlations between 2MASS and ROSAT detection and optical variability were found. Those quasars which were detected in either the IR or the X-ray were significantly more likely to be seen as variable than those quasars not detected. In each of the non-optical surveys, the detected sub-sample was then divided into variable and non-variable populations. Both the IR- and X-ray-detected populations showed that the variable population had a lower median luminosity than the non-variable population.  Unlike the IR and X-ray surveys, the FIRST survey showed no difference in detection between the variable and non-variable populations. The median integrated radio flux density, however, was larger for the variable population than the non-variable population, also in contrast with the other non-optical surveys. \\citet{Vanden04} reported the same behavior and suggest blazars as a possible explanation.  A simple light curve morphology analysis shows that the unique, aperiodic time signature inherent to quasars can be utilized to further refine variability selection techniques. The $Q_i$ parameter efficiently separates variable point source objects in our data into groups of periodic (i.e., stellar) and aperiodic (i.e., quasars etc.) objects. While somewhat fine-tuned to the particulars of the QVS, this is an appealing test, as it is easy to implement, takes advantage of the light curve details, and is fairly robust to unevenly sampled data. These techniques used on the QVS are a useful testbed for techniques to identify and characterize quasars using synoptic photometric data without the need for follow-up spectroscopy."
    },
    "0512/astro-ph0512640_arXiv.txt": {
        "abstract": "We analyzed a sample of high and low surface brightness (HSB and LSB) disc galaxies and elliptical galaxies to investigate the correlation between the circular velocity (\\Vc ) and the central velocity dispersion (\\sigmac ). We better defined the previous \\vsigma\\ correlation for HSB and elliptical galaxies, especially at the lower end of the \\sigmac\\ values.  Elliptical galaxies with \\Vc\\ based on dynamical models or directly derived from the \\hi\\ rotation curves follow the same relation as the HSB galaxies in the \\vsigma\\ plane. On the contrary, the LSB galaxies follow a different relation, since most of them show either higher \\Vc\\ (or lower \\sigmac ) with respect to the HSB galaxies. This argues against the relevance of baryon collapse in the radial density profile of the dark matter haloes of LSB galaxies. Moreover, if the \\vsigma\\ relation is equivalent to one between the mass of the dark matter halo and that of the supermassive black hole, these results suggest that the LSB galaxies host a supermassive black hole with a smaller mass compared to HSB galaxies of equal dark matter halo. On the other hand, if the fundamental correlation of SMBH mass is with the halo \\Vc , then LSBs should have larger black hole masses for given bulge \\sigmac . ",
        "introduction": "\\label{sec:introduction}  A possible relation between the central velocity dispersion of the spheroidal component (\\sigmac ) and the galaxy circular velocity (\\Vc ) measured in the flat region of the rotation curve (RC) was suggested by Whitmore et al. (1979). By measuring stellar velocity dispersions and \\hi\\ line widths for a sample of 19 spiral galaxies they found a a significant decrease in \\Vc$/$\\sigmac\\ with increasing bulge-to-disk ratio. Since \\sigmac\\ and \\Vc\\ probe the potential of the spheroid and dark matter (DM) halo, a mean value \\Vc$/$\\sigmac\\ $\\simeq 1.7$ implies these components are dynamically separate with the bulge substantially cooler than halo.  Gerhard et al. (2001) derived the \\vsigma\\ relation for the sample of giant ellipticals studied by Kronawitter et al. (2000). This was explained as an indication of near dynamical homology of these objects which were selected to be nearly round and almost non-rotating elliptical galaxies. These galaxies form a unique dynamical family which scales with luminosity and effective radius. As a consequence the maximum \\Vc\\ of the galaxy is correlated to its \\sigmac . Whether the same is true for more flattened and fainter ellipticals is still to be investigated. On the contrary, both shape and amplitude of the RC of a spiral galaxy depend on the galaxy luminosity and morphological type (e.g., Burstein \\& Rubin 1985). As a consequence for spiral galaxies the \\vsigma\\ relation is not expected a priori.  Nevertheless, Ferrarese (2002) and Baes (2003) found that elliptical and spiral galaxies define a common \\vsigma\\ relation. In particular, it results that for a given \\sigmac\\ the value of \\Vc\\ is independent of the morphological type. But \\sigmac\\ and \\Vc\\ are related to the mass of the supermassive black hole (hereafter SMBH, see Ferrarese \\& Ford 2005 for a review) and DM halo (e.g., Seljak 2002), respectively. Therefore Ferrarese (2002) argued that the \\vsigma\\ relation is suggestive of a correlation between the mass of SMBH and DM halo.  Previous works concentrated on high surface brightness (HSB) galaxies. It is interesting to investigate whether the \\vsigma\\ relation holds also for less dense objects characterized by a less steep potential well. This is the case of low surface brightness (LSB) galaxies, which are disc galaxies with a central face-on surface brightness $\\mu_B\\geq22.6$ mag arcsec$^{-2}$ (e.g., Impey et al. 1996). In Pizzella et al. (2005) we studied the behavior of LSB galaxies in the \\vsigma\\ relation. Here we present our results. ",
        "conclusions": ""
    },
    "0512/astro-ph0512530_arXiv.txt": {
        "abstract": "{ We continue the study of the effects of a strong magnetic field on the temperature distribution in the crust of a magnetized neutron star (NS) and its impact on the observable surface temperature. Extending the approach initiated in Geppert et al. (\\cite{GKP04}), we consider more complex and, hence, more realistic, magnetic field structures but still restrict ourselves to axisymmetric configurations. We put special emphasis on the heat blanketing effect of a toroidal field component. We show that asymmetric temperature distributions can occur and a crustal field consisting of dipolar poloidal and toroidal components will cause one polar spot to be larger than the opposing one. These two warm regions can be separated by an extended cold equatorial belt. As an example we present an internal magnetic field structure which can explain, assuming local blackbody emission, both the X-ray and optical spectra of the isolated NS RXJ 1856-3754, the hot polar regions dominating the X-ray flux and the equatorial belt contributing predominantly to the optical emission. We investigate the effects of the resulting surface temperature profiles on the observable lightcurve which an isolated thermally emitting NS would produce for different field geometries. The lightcurves will be both qualitatively (deviations from sinusoidal shape) and quantitatively (larger pulsed fraction for the same observational geometry) different from those of a NS with an isothermal crust. This opens the possibility to determine the {\\em internal} magnetic field strengths and structures in NSs by modeling their X--ray lightcurves and spectra. The striking similarities of our model calculations with the observed spectra and pulse profiles of isolated thermally emitting NSs is an indication for the existence of strong magnetic field components maintained by crustal currents. ",
        "introduction": "It has been understood since Greenstein \\& Hartke (\\cite{GH83}) that in presence of a sufficiently strong magnetic field, $\\ge 10^{10}$ G, the surface temperature of a neutron star (NS) will not be uniform as is expected in the unmagnetized case. Anisotropy of heat transport, caused by both classical and quantum magnetic field effects, in the thin low density ($\\rho \\le 10^{10}$ g cm$^{-3}$) upper layer, the so called {\\em envelope}, results in a strongly reduced conductivity in the direction perpendicular to the field and an enhanced one along the field. As a result, the regions around the magnetic poles are expected to be significantly warmer than the regions around the magnetic equator. Page  (\\cite{P95}) and Page \\& Sarmiento  (\\cite{PS96}) applied the Greenstein \\& Hartke formula to explain the lightcurves of the isolated thermally emitting NSs PSR 0833-45 (Vela), PSR 0656+14, PSR 0630+178 (Geminga) and PSR 1055-52, considering dipolar field configurations, and also the addition of a quadrupolar component, and including the General Relativistic curvature effects on photon trajectories. Much work has been dedicated to study the effects of the magnetic field on the properties of the NS envelope and crust (for reviews, see Yakovlev \\& Kaminker \\cite{YK94}, Ventura \\& Potekhin \\cite{VP01}, Lai \\cite{L01}). Due to the shallowness of the envelope, $\\sim$ 100 m, heat transport can be treated in the plane parallel approximation as a one dimensional problem in which the heat flux is purely radial with a, locally uniform, magnetic field having some arbitrary orientation with respect to the radial direction. The most recent studies (Potekhin \\& Yakovlev \\cite{PY01}, henceforth `PY01, and Potekhin et al. \\cite{PYCG03}) have included the best up to date transport coefficients and equation of state: they showed some deviations from the simple Greenstein \\& Hartke formula and finally give us a reasonably accurate description of a magnetized NS envelope.  In presence of a sufficiently strong magnetic field, $\\simgreater 10^{12} - 10^{13}$ G, the anisotropy of heat transport, simply due to the classical effect of Larmor rotation of the electrons, extends to much higher densities and can even be present within the whole crust. Recently, we have shown (Geppert et al. \\cite{GKP04}, subsequently Paper I) that in cases where the field geometry in the crust is such that the meridional component of the field dominates over its radial component in a large part of the crust, as is the extreme case of a magnetic field entirely confined to the crust, the non uniformity of the temperature, previously considered to be restricted to the envelope, may actually extend to the whole crust. This modified crustal temperature results in a different surface temperature distribution than what the simple Greenstein \\& Hartke model predicted. In contradistinction, in case the dipolar poloidal field has its currents localized only in the core of the star, i.e., the field in the crust is considered as the one of a vacuum dipole, the crust is practically isothermal and the non uniformity of the surface temperature is then only due to field effects in the envelope. This result, that the geometry of the magnetic field {\\em in the interior} of the NS leaves an observable imprint at the surface, potentially allows us to study the internal structure of the magnetic field through modelling of the spectra and pulse profile of thermally emitting NSs.   There exist growing observational evidence that the anisotropy of the heat transport in the envelope alone, assuming an otherwise isothermal crust, can not explain the surface temperature distributions of some observed NSs. The seven X-ray dim isolated NSs (XDINSs), dubbed \"The Magnificent Seven\" (for reviews, see, e.g., Haberl \\cite{H04,H04b,H05}) are nearby isolated NSs, all discovered in the X-ray band where they show a thermal spectrum. As a set they share several properties: 1) low X-ray absorptions imply small distances, $<$ 200 -500 pcs, 2) high blackbody  temperatures deduced from the fit of their X-ray spectra indicate ages of the order of $10^6$ yrs, possibly confirmed by measurements of proper motions which allow for tentative identification of their birth sites in clouds of the Gould's belt, 3) no detection of radio emission, 4) no detection of a hard X-ray tail, typical of magnetospheric emission in active pulsars, 5) no association with a supernova remnant, 6) detected spin periods (in 5 cases) between 3 to 11 s, 7) estimates of their surface dipolar field strength $B_0$ between $1 - 10 \\times 10^{13}$ G, either from a measurement of the period derivative $\\dot{P}$ (1 case) or the interpretation of broad absorption lines (in 5 cases) in the X-ray spectrum as being proton cyclotron lines, 8) optical broad band photometric detections (5 cases) can be interpreted as being the Rayleigh-Jeans tail of a blackbody. However, these optical data are well above the Rayleigh-Jeans tail of the blackbody detected in the X-ray (``optical excess'') and indicate the presence of an extended cold component of much larger area than the warm component observed in X-ray, the latter having an emitting radius ($\\sim 3-5$ km) much smaller that the usually assumed radius of a NS($\\sim 10 - 15$ km). Schwope et al. (\\cite{SHHM05}) tried to fit the lightcurve of  RBS 1223 and concluded that only a surface temperature profile with relatively small, about 4-5 kms across, hot polar regions may explain the observations. Pons et al. (\\cite{Petal02}) and Tr\\\"umper et al. (\\cite{TBHZ04}) arrived qualitatively at the same conclusion when they fitted the combined X-ray and optical spectrum of  RX J1856.5-3754. In both cases, the smallness of the hot region is much below what can be reached by considering anisotropic heat transport limited to only a thin envelope. There are, however, not many possibilities to produce non uniform surface temperature distributions. Principially, rotation leads to an oblateness which enhances the thickness of the insolating outer layer at the equatorial belt making that region cooler than the poles (e.g. Geppert \\& Wiebicke \\cite{GW86}). This effect is, however, even for the fastest rotating millisecond pulsars, too small and a fortiori neglibigle for those isolated thermally emitting NSs whose rotational period is in the order of seconds. Another way to heat the polar regions is the bombardment by ultrarelativistic charged particles as can be expected in radiopulsars (see Gil et al. \\cite{GMG03}). Such heating is mostly confined to the polar caps whose angular radius can be roughly estimated as $\\theta_\\mathrm{pc} \\sim (2\\pi R/cP)^{1/2}$, $R$ being the star's radius, $c$ the speed of light and $P$ the pulsar's rotational period, giving a polar cap radius $r_\\mathrm{pc} \\sim 0.46 \\; P^{-1/2}$ km (for $P$ measured in sec.), which is much smaller than the estimated size, a few kms, of the warm regions in the X-ray dim isolated NSs whose periods are in the range of 3 to 11 s.   The previous Paper I  was devoted to demonstrate that a strong magnetic field can have significant effects on the crustal temperature distribution. We solved the stationary heat transport equation in Newtonian approximation and considered very simple, purely dipolar poloidal crustal field structures. Our goal here is to consider more realistic models in which the currents maintaining the poloidal magnetic field are distributed between the core and the crust and we also consider the possible presence of a strong toroidal component in the crust. We will still consider only axisymmetric, dipolar, magnetic field configuration. Moreover, we now perform a wholly general relativistic formulation of the heat transport and energy balance equations.  There exist also theoretical reasons that the presence of a strong toroidal field is very likely. In a series of papers Markay \\& Tayler \\cite{MT73} (and references therein) showed that a stable magnetic field configuration needs a coexistence of poloidal and toroidal field components, having approximately the same strength. Recently, Braithwaite \\& Spruit (\\cite{BS04}) considered the stability of pure poloidal field geometries numerically in Ap stars and in magnetic white dwarfs and arrived at the same conclusion.  Little is known about the magnetic field structure in NSs which is very likely determined by processes during the proto-NS phase and/or in a relatively short period after that epoch. A proto-NS dynamo (Thompson \\& Duncan \\cite{TD93}) is unlikely to generate purely poloidal fields while differential  rotation will easily wrap any poloidal field and generate strong toroidal components (Klu\\'zniak \\& Ruderman \\cite{KR98}; Wheeler, Meier,  \\& Wilson \\cite{WMW02}). The magneto-rotational instability (Balbus \\& Hawley \\cite{BH91}) also most certainly  acts in proto-NSs (Akiyama {\\em et al.}  \\cite{AWML03}) and results in toroidal fields from differential rotation (Balbus \\& Hawley \\cite{BH98}). It is not yet explored whether the onset of supercondctivity in the core and the relatively  fast crystallization of the largest part of the crust allows the field to relax into a stable (force free) state. However, it seems realistic to consider the effect of magnetic field configurations which consist of poloidal AND toroidal crustal components as well as of a star centered poloidal one. The relative strengths as well as the positions of the crustal components will depend on the - certainly varying - creation processes of those fields and will therefore be varied in reasonable limits.  The paper is organized as follows: In the next section, \\S~\\ref{sec:magn-param}, we define the axisymmetric field structures and consider the resulting 2D heat transport and energy balance equations. The next section, \\S~\\ref{sec:results1}, presents our results for the crustal temperature distributions resulting from a large set of different field geometries and \\S~\\ref{sec:results2} studies the observable effects of the resulting surface temperature distributions, assuming simple isotropic blackbody emission. As an example we present a crustal magnetic field structure which allows us to reproduce the two temperature spectral fits of RX J1856.5-3754. A discussion of our results and conclusions are the objects of \\S~\\ref{sec:discon}. ",
        "conclusions": "\\label{sec:discon}  We have studied the impact of different crustal magnetic field configurations on the transport of heat within the crust of a NS and the resulting surface temperature distributions with their observable consequences. In a previous paper (Paper I) we had first considered the simplest cases of a poloidal dipolar field either entirely confined to the crust (``crustal field'') or exclusively produced by currents in the core and hence having in the crust the same structure as a vacuum dipole (``core field''). In the present work we have considered the effect of an additional (dipolar) toroidal component, located within the crust, with various combinations of core and crustal fields, and all three components having the same symmetry axis.  Since we consider here only poloidal and toroidal dipolar components, due to the axial symmetry the toroidal field component is confined to the interior of the NS and the external field geometry is always a purely poloidal dipolar one uniquely characterized by its strength at the magnetic pole, $B_0$. Therefore, the effects of various magnetic field structures with the same $B_0$ in the magnetosphere, as e.g. pulsar properties and rotational evolution of the NS, are indistinguishable. However, these different inner geometries of the magnetic field result in very different surface temperature distributions which are potentially distinguishable through observations of thermal emission from isolated cooling NSs.  We had shown in Paper I that a strong magnetic field practically channels the heat flow along its field lines. We had found that a core field, with field lines essentially radial in most of the crust, results in an almost isothermal crust and the surface temperature is hence only controlled by the anisotropic transport in the uppermost layers of the envelope (Greenstein \\& Hartke \\cite{GH83}, Page \\cite{P95}) exhibiting two symmetric extended warm regions, one in each hemisphere, and a cold belt along the magnetic equator. In contradistinction, a crustal field, having field lines with large meridional components, can produce strongly non-isothermal crusts and the heat flow is much more concentrated around the magnetic poles, resulting in a much wider cold equatorial belt. Our main result here is that the addition of a toroidal component, inside the crust, is able to add a very efficient heat blanket which may force the heat to flow within a narrow region along the polar axis and result in a cold region covering most of the stellar surface with two warm spots arround the two magnetic poles. However, we find that for this effect to occur it is essential to have a significant part of the poloidal field produced in the crust: the toroidal field in itself is not sufficient in case the field lines of the poloidal component are almost radial.  The presence of two small warm regions separated by an extended cold belt has two inmediate observational consequences. The first is that the observable pulsed fraction in the X-ray band can be very large, above 30\\% assuming isotropic blackbody emission, and, second, the cold region's emission contributes little to the X-ray flux but dominates the detectable flux in the optical range, appearing as an ``optical excess''. Notice, however, that addition to the dipolar poloidal field of a quadrupolar component also allows, potentially, to reach high pulsed fractions (Page \\& Sarmiento \\cite{PS96}) within the same isotropic blackbody emission scheme and without any todoidal component. In this latter case, high $P_f$'s are reached when the quadrupole sufficiently deforms the field to push the two magnetic poles close to each other but without significantly altering the star's effective and average temperature (see, e.g., Figure 1 of Page \\& Sarmiento \\cite{PS96}), and hence producing no significant ``optical excess''. The ``Magnificent Seven'' briefly presented in the Introduction seem to be very good candidate NSs for which our results may be relevent. Their estimated dipolar surface field strengths $B_0$ are above $10^{13}$ G, opening the possibility of strong anisotropic heat transport in their crust, and the ``optical excess'' observed in 5 cases is then a natural consequence of the large $B_0$'s in case their crusts harbor very strong poloidal fields. As we has seen, channeling of the heat flux toward two small polar regions requieres moreover that part of the poloidal flux be confined to the crust: this can be achieved either by the action of some dynamo process which produced such a poloidal field component or/and by the expulsion of flux from the core through the action of the outward migration of the neutron superfluid fluxoid resulting from the NS spin-down and the consecutive expulsion of the magnetic fluxoid resulting from the core proton superconductor. This latter process is likely to have occured in the ``Magnificent Seven'' given their unusually long period, for their relatively young ages, which are very likely the results of fast spin-down due to their strong dipolar poloidal fields. However, our axisymmetric field configurations produce symmetric, but not sinusoidal, light-curves and the superposition of a quadrupolar component may be necessary to produce precise fits of the observed pulse profiles of the seven XDINS (Zane \\& Turolla \\cite{ZT05})  Results very similar to the ones presented here have been recently obtained by P\\'erez-Azor\\'in, Miralles \\& Pons (\\cite{PAMP05}) who also conclude that a toroidal component localized within the crust lead to surface temperature distributions with very localized warm regions and an extended cold equatorial belt. These authors also took into account the transport of heat by phonons, which is unaffected by the magnetic field and thus reduces the anisotropicity, and the small quantum corrections present at high densities for very strong fields, two effects we did not consider. Moreover, they performed their heat transport calculations to lower densities than we have considered here and assumed that the upper layer is cut at a finite density due to solidification of the material while we simply assumed that the envelope extends into an atmosphere and ``glued'' our models with pre-existent envelope models. Nevertheless, despite of these differences, the similarity of the results prove that the effect of a toroidal field is independent of the details and make our common results very robusts.  Claiming that the observed optical-X-ray properties of the ``Magnificent Seven'' prove the existence of strong toroidal fields in their crust may still be premature. Given our poor understanding of the emissive properties of the NS surface, being it either an atmosphere, a liquid or a solid depending on its chemical composition, temperature and magnetic field strength, one cannot exclude that the ``optical excess'' is achieved locally, i.e., is an intrinsic feature of the spectrum produced at each point of the surface, instead of being a global effect due to strong surface temperature differences as we have obtained. In contradistinction, three very active pulsars exhibit very different properties (Vela: Romani, Kargaltsev \\& Pavlov 2005, \\cite{RKP05}, PSR B0656+14: Shibanov et al. \\cite{SSLGL05}, and Geminga: Kargaltsev et al. \\cite{KPZR05}): a thermal spectrum is clearly seen in the X-ray but near \\& far UV observations give upper limits on its Rayleigh-Jeans tail which are in agreement or even {\\em below} the extrapolations of the X-ray spectrum. The reasons for this different behavior is presently unknown and may be due to their weaker magnetic field, different surface chemical composition, and/or reprocessing of the thermal photons in the magnetosphere.  Finally, an issue raised by our models is the time evolution and stability of the field configurations we have considered. We intend to perform coupled time-dependent cooling and field evolution calculations to try to eliminate at least the most unlikely configurations, i.e., either highly unstable configurations and/or configurations based on current distribution which will evolve fast and reajust or decay. Nevertheless we can speculate that a very strong crustal toroidal field could produce such strong tension that the crust should readjust itself, as is seen in magnetars, and this may be the cause of the observed change in the spectrum and pulse profile of RX J0720.4-3125 (de Vries et al. \\cite{dVVMV04})."
    },
    "0512/hep-ph0512118_arXiv.txt": {
        "abstract": "We consider supersymmetric models with right-handed neutrinos where neutrino masses are {\\it purely Dirac-type}.  In this model, right-handed sneutrino can be the lightest supersymmetric particle and can be a viable candidate of cold dark matter of the universe. Right-handed sneutrinos are never thermalized in the early universe because of weakness of Yukawa interaction, but are effectively produced by decays of various superparticles.  We show that the present mass density of right-handed sneutrino can be consistent with the observed dark matter density. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512277_arXiv.txt": {
        "abstract": "The relation between X-ray luminosity and near-infrared luminosity for early-type galaxies has been examined.  Near-infrared (NIR) luminosities should provide a superior measure of stellar mass compared to optical luminosities used in previous studies, especially if there is significant star-formation or dust present in the galaxies.  However, we show that the X-ray-NIR relations are remarkably consistent with the X-ray-optical relations.  This indicates that the large scatter of the relations is dominated by scatter in the X-ray properties of early-type galaxies, and is consistent with early-types consisting of old, quiescent stellar populations.  We have investigated scatter in terms of environment, surface brightness profile, Mg$_{2}$, H$\\beta$, H$\\gamma$ line strength indices, spectroscopic age, and nuclear H$\\alpha$ emission.  We found that galaxies with high Mg$_{2}$ index, low H$\\beta$ and H$\\gamma$ indices or a `core' profile have a large scatter in \\lx, whereas galaxies with low Mg$_{2}$, high H$\\beta$ and H$\\gamma$ indices or `power-law' profiles, generally have \\lx$<10^{41}$ erg s$^{-1}$.  There is no clear trend in the scatter with environment or nuclear H$\\alpha$ emission. ",
        "introduction": "Early-type galaxies are known to emit X-rays via discrete populations of low-mass X-ray binaries and via a hot thermal plasma (\\citealt{fab84}; \\citealt{for85}; \\citealt{fab89}).  The correlation of X-ray luminosity (\\lx) and \\emph{B}-band optical luminosity (\\lb) has been studied to examine the variation of gas properties as a function of galaxy mass (see e.g.\\ \\citealt{osul01}; \\citealt{mat03} for a review).  It is found that there is a transition in the X-ray emission, and hence the hot gas content, of galaxies at \\lb$\\approx3 \\times 10^{10}$\\lbsun.  Galaxies less luminous than this threshold have X-ray emission dominated by discrete sources, and display a relation \\lx$\\propto$\\lb.  Above this threshold, there is an excess of X-ray emission attributed to hot gas emitting via thermal Bremsstrahlung, and hence a steeper relation of \\lx$\\propto$\\lb$^{\\sim 2}$ (\\citealt{osul01}).  The scatter of the \\lx:\\lb\\ relation is very large, varying by $\\sim2$ orders of magnitude for galaxies of similar \\lb\\ (\\citealt{mat03}).  Several physical phenomena have been examined to account for this scatter.  \\citet{mat03} list: environmental and intrinsic galaxy properties (\\citealt{whit91};  \\citealt{esk95a}; \\citealt{esk95b}; \\citealt{esk95c}; \\citealt{hen99}; \\citealt{bro00}), differences between power-law ellipticals and core ellipticals due to their different location within groups and the additional contribution of intragroup gas to centrally located ellipticals (\\citealt{pel99}; \\citealt{hel01}; \\citealt{mat01}), the influence of rotation and flattening (\\citealt{nul84}; \\citealt{kle95}; \\citealt{brig96}; \\citealt{pel97}; \\citealt{der98}), the influence of type Ia supernov\\ae\\ and cooling flows (\\citealt{cio91}), the extent of the diffuse gas (\\citealt{mat98}) and ram-pressure stripping (\\citealt{ton01}).  Recently \\citet{pip05} have shown that a late secondary accretion of gas can account for the \\lx:\\lb\\ relation of massive galaxies, and the large observed scatter.  These phenomena play a role in contributing to the large scatter in the \\lx:\\lb\\ relation, but there is often conflicting evidence and/or too little scatter explained.  All the phenomena above relate to an increase in scatter of the X-ray luminosity.  The assumption that the scatter is due to variance in \\lx\\ rather than \\lb\\ is reasonable, since the stellar populations of early-type galaxies are generally thought to be old, passively evolving systems, especially for galaxies in clusters (e.g.\\ \\citealt{sta98}; \\citealt{dep99}; \\citealt{bel04}; \\citealt{hol04}; \\citealt{ell04}).  However, some authors find that there is some evolution of early-type galaxies not attributable to purely passive evolution (e.g.\\ \\citealt{but84}; \\citealt{dre97}; \\citealt{dro03}; \\citealt{vdv03}; \\citealt{red04}) in agreement with models of structure formation which predict a merger origin for elliptical galaxies (e.g.\\ \\citealt{bau96}).  If elliptical galaxies are not all old and quiescent then there may be significant inconsistencies between the optical \\emph{B}-band luminosity and the mass of a galaxy.  The presence of young stars can significantly enhance the blue luminosity of a galaxy, and hence the \\emph{B}-band would yield an inaccurate estimate of galaxy mass in the presence of young stellar populations.  Thus, the assumption that the \\emph{B}-band luminosity contributes an insignificant scatter to the \\lx:\\lb\\ relation needs to be tested.  Near-infrared (NIR) luminosities are much more closely correlated with dynamical and total galaxy mass than optical luminosities (\\citealt{gav96}), since NIR emission is dominated by longer lived stellar populations, and has little contribution from young, bright stars.  Figure~\\ref{fig:examplespecs} shows the model spectral energy distribution of a simple stellar population at ages, 10Myr, 100Myr, 500Myr, 1Gyr and 9Gyr, from the libraries of \\citet{bru03}.  Also shown are the transmission profiles of the \\emph{B}, \\emph{H} and \\emph{Ks} bands.  The difference between the two stellar populations due to the massive, short lived stars will clearly have a much stronger effect on the \\emph{B}-band luminosity than the \\emph{H} or \\emph{Ks} luminosities, e.g.\\ between the ages of 100Myr and 9Gyr, the \\emph{B} band luminosity of a simple stellar population decreases by $\\approx 14$ times more than the decrease in the  \\emph{K} band luminosity, or $\\approx 5$ times between 1Gyr and 9Gyr.  \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.4]{fig1.eps} \\caption{Model spectral energy distributions from the libraries of \\protect\\citealt{bru03}.  The spectra are 10Myr, 100Myr, 500Myr, 1Gyr and 9Gyr old, with the younger spectra being brighter.   The units of flux are arbitrary.  Also illustrated are the locations of the transmission profiles of the \\emph{B}, \\emph{H} and \\emph{Ks} bands. } \\label{fig:examplespecs} \\end{figure}  The NIR is also much less affected by the presence of dust.  Galactic extinction in the \\emph{K} band is typically a factor of $\\sim0.08$ times the \\emph{B}-band extinction (\\citealt{schl98}).  Thus if some elliptical galaxies contain significant amounts of dust the NIR will be a better indicator of the underlying stellar mass.  About 80\\% of early-type galaxies contain dust, though this is mainly confined to the central kpc (\\citealt{mat03}).  Theoretical work suggests that more widely distributed dust will be rapidly sputtered (e.g.\\ \\citealt{tsa95}), therefore the presence of dust is unlikely to affect global NIR$-$optical colours.  However, more extended filaments of dust have been observed in some galaxies (\\citealt{tra01}).  Any effect due to dust, though expected to be small, will be much reduced in the NIR.  The improved measure of galaxy mass attainable through NIR photometry provides an opportunity to recast the \\lx:\\lb\\ relation as an \\lx:\\lk\\ relation (or \\lx:\\lh).  Doing so will test the importance of variance in \\lb\\ in contributing to the scatter of the \\lx:\\lb\\ relation. In turn the comparison of the \\lx:\\lb\\ and \\lx:\\lk\\ relations may yield information on the stellar populations of elliptical galaxies, in particular the presence of young stars.  Accordingly we have reanalysed the X-ray catalogue of early-type galaxies of \\citet{osul01}, supplemented with \\emph{H} and \\emph{Ks} band luminosities from the Two Micron All Sky Survey (2MASS).  The paper is organised as follows.  Section~\\ref{sec:samp} introduces the X-ray catalogue of \\citet{osul01}, and section~\\ref{sec:nir} describes the NIR data.  Section~\\ref{sec:method} describes the statistical analyses used to fit the relations.  Section~\\ref{sec:results} describes the correlations, and the new measurements of the scatter as a function of environment, surface brightness profile, Mg$_{2}$, H$\\beta$ and H$\\gamma$ emission line strength and nuclear H$\\alpha$ emission.  Finally section~\\ref{sec:discussion} discusses the results in terms of the origin of the scatter and the stellar populations of early-type galaxies. ",
        "conclusions": "\\label{sec:discussion}  The relationship between X-ray luminosity and NIR luminosity for early-type galaxies has been examined.  The results are consistent with previous work studying the relationship between X-ray luminosity and optical luminosity (\\citealt{osul01}).   The \\lx:\\lb, \\lx:\\lh\\ and \\lx:\\lk\\ relations bear the same overall trend, becoming steeper for galaxies that are bright in the optical or NIR.  This is consistent with previous interpretations in which more massive galaxies contain larger amounts of hot gas (see \\citealt{mat03} for a review).  The scatter of the relations is the same in the \\emph{B}, \\emph{H} and \\emph{Ks} bands.  This is probably largely because early-type galaxies generally contain very little dust and are composed of old, quiescent stellar populations.  Any scatter introduced  by residual star-formation or dust present in early-types,  must be insignificant compared to the large scatter of the X-ray luminosities.  The scatter and slope of the relations has been investigated as a function of environment.  There is no clear trend with environment, either in terms of cluster, group or field membership or in terms of local galaxy density. Brightest group galaxies display a steeper relation and significantly higher \\lx/\\lb, \\lx/\\lh\\ and \\lx/\\lk, than other types of galaxies, perhaps indicative of a different formation mechanism or their location in the densest regions of the IGM, in agreement with \\citet{osul01} and \\citet{hel01}.  There may also be a slight increase in the \\lx/\\lb\\ etc.\\ values for group galaxies in general again suggestive that the IGM is influential on the X-ray emission from galaxies through processes such as accretion or stifling of galactic winds.  However cluster galaxies, which would be subject to similar processes are consistent with \\lx/\\lb\\ values of field galaxies, making any firm conclusions difficult to draw.  It is possible that differences between group and cluster galaxies, such as merger history, ram-pressure stripping, etc.\\ may be responsible for the lower \\lx/\\lb\\ of cluster galaxies.  For example, it is known that group galaxies are involved in  significantly more galaxy-galaxy interactions than cluster galaxies, since the higher velocity dispersion of clusters make the chances of galaxy-galaxy interaction much smaller (\\citealt{zep93}).  Galaxies with a shallow inner core profile cover a wider range in \\lx\\ than power-law galaxies with a steep inner profile, which are generally restricted to log\\lx$<41$ (erg s$^{-1}$).  This is in agreement with previous work by \\citet{pel99}.  Furthermore the core profile galaxies are observed mainly to occupy the steeper part of the \\lx:\\lb\\ relation, while the power-law galaxies are generally on the shallower part of the relation.  This is suggestive that whatever is causing the break in the \\lx:\\lb\\ relation is also responsible for the observed differences between core and power-law galaxies, such as rotation velocity, luminosity etc.    However, this it is difficult to test this as the large scatter in both relations makes a comparison of the \\emph{B} band luminosities of the breaks hard.  An obvious difference is the mass of the systems, but different formation mechanisms are also often speculated as the cause of the difference between power-law and core galaxies (e.g.\\ \\citealt{fab97}; \\citealt{ryd01}; \\citealt{burkert03}; \\citealt{kho03}).  The effects of age and metallicity were investigated via absorption line strengths.  Mg$_{2}$ index, which is degenerate in age and metallicity, shows a clear demarcation with high index galaxies having a large spread in \\lx\\ while lower index galaxies are confined to lower values of \\lx.  Similar trends are also seen for H$\\beta$ and H$\\gamma$ index which are less degenerate, and more sensitive to age (\\citealt{ter02}).  Although this provides support for a relation between age and \\lx, as reported by \\citealt{osul01b}, it seems likely that the large spread in \\lx\\ for high Mg$_{2}$ index galaxies is also partly due to a higher mass and subsequent higher metallicity.  The age-metallicty degeneracy can be broken using combinations of emission lines.  Using the spectroscopic ages of \\citet{ter02} reduces the size of the sample, but reveals a strong correlation between \\lx/\\lk\\ and age.  The contribution to this correlation due to the dimming of \\lk\\ over time is negligible, supporting the results and conclusions of \\citet{osul01b}.  Known AGN were removed from the sample before any statistical analysis was performed.  However it is possible that some galaxies with a low level of nuclear activity remain in the sample.  Galactic nuclear activity has been quantified using H$\\alpha$ emission line strength, as reported by \\citet{ho97a} after subtraction of the contribution from starlight.  Broad H$\\alpha$ emission lines are a classic indication of a Seyfert or LINER (low-ionisation nuclear emission-line region) galaxy .  Seyfert and LINER galaxies are generally spirals, but occasionally are found in ellipticals.  The lack of correlation of either \\lx\\ or log\\lx/\\lb\\ with H$\\alpha$ emission is suggestive that galactic nuclear activity is not very influential on the observed scatter of the \\lx:\\lb\\ relation.  This is corroborated by the similar range in values of \\lx\\ and log\\lx/\\lb\\ displayed by the H$\\alpha$ emitting galaxies and the whole sample.  In conclusion, the origin of the large scatter in \\lx/\\lb\\ remains elusive.  It is clear that mass, age, metallicity and core-profile are all linked with the range in \\lx\\ displayed by galaxies, with more massive and older galaxies having a wider range in \\lx\\ and cuspy profiles.   The underlying cause of these differences may well be linked with formation and merger history as well as the  environment of the galaxies, but disentangling the effects is very problematic. Conversely it appears that star-formation and recent post-merger objects do not significantly contribute to the scatter, and thus studies of the optical-X-ray luminosity relations are a valid choice.   Similarly the local density and environment of early-type galaxies do not appear to affect the scatter."
    },
    "0512/astro-ph0512041_arXiv.txt": {
        "abstract": "The $r$-modes of accreting neutron stars could be a detectable source of persistent gravitational waves if the bulk viscosity of the stellar matter can prevent a thermal runaway. This is possible if exotic particles such as hyperons are present in the core of the star. We compute bulk viscous damping rates and critical frequencies for $r$-modes of neutron stars containing hyperons in the framework of relativistic mean field theory. We combine the results of several previous calculations of the microphysics, include for the first time the effect of rotation, and explore the effects of various parameters on the viability of persistent gravitational wave emission. We find that persistent emission is quite robust, although it is disfavored in stars below 1.3--1.5~$M_\\odot$ depending on the equation of state. In some cases persistent emission is compatible with temperatures as low as $10^7$~K, observed in some accreting neutron stars in quiescence. ",
        "introduction": "The $r$-modes (fluid oscillations governed by the Coriolis force) of rapidly rotating neutron stars have attracted much interest as possible sources of gravitational waves and mechanisms for regulating the spins of neutron stars. See Ref.~\\cite{Stergioulas:2003yp} for a recent review of the many physical and astrophysical issues related to the $r$-modes; here we focus on gravitational wave emission. Gravitational radiation drives the $r$-modes unstable and could lead to detectable gravitational wave emission in two scenarios. In one scenario, a newborn neutron star could radiate a substantial fraction of its rotational energy and angular momentum as gravitational waves changing in frequency on a timescale of a year or more~\\cite{Owen:1998xg}. In the other scenario, $r$-modes in rapidly accreting neutron stars in low-mass x-ray binaries (LMXBs) could be persistent sources of periodic gravitational waves~\\cite{Bildsten:1998ey, Andersson:1998qs}. (Here the $r$-modes provide a specific mechanism for a more general torque balance argument~\\cite{Papaloizou:1978, Wagoner:1984pv}).  Currently the latter scenario for gravitational wave emission (from LMXBs) looks like a brighter prospect for detection. There is now evidence from several approaches that the amplitude of an $r$-mode growing due to the instability is limited by nonlinear fluid dynamics to a relatively small value~\\cite{Schenk:2001zm, Morsink:2002ut, Arras:2002dw, Brink:2004bg, Brink:2004qf, Brink:2004kt, Sa:2004gn}. While low-amplitude long-lived $r$-modes in newborn neutron stars still can lead to astrophysically interesting effects such as the regulation of spins, a low mode amplitude renders the gravitational wave signal undetectable unless there is a very nearby supernova~\\cite{Owen:1998xg}. Also, if neutron stars contain particles more exotic than neutrons and protons---such as hyperons, where an up or down quark in a nucleon is replaced by a strange quark---there are additional viscous damping mechanisms which may eliminate the instability altogether in very young, hot neutron stars~\\cite{Langer:1969, Jones:1970, Jones:1971, Jones:2001ie, Jones:2001ya, Lindblom:2001hd, Haensel:2001em, vanDalen:2003uy} or strange quark stars~\\cite{Madsen:1998qb}. However, for the LMXB scenario, low amplitude is not a problem~\\cite{Bildsten:1998ey, Andersson:1998qs} and the additional viscosity actually renders stars with exotic particles better candidates for gravitational wave detection. The reason for the latter is subtle, and bears explanation.  Real neutron stars are not perfect fluids, and thus viscous (and other) damping mechanisms compete with gravitational wave driving of the $r$-modes. The strengths of the driving and damping mechanisms can be expressed as timescales which depend on the rotation frequency and temperature of the star (usually assumed to be nearly isothermal). Therefore it is useful to consider the location of a star in the temperature-frequency plane, as shown in Fig.~\\ref{runaway}. The strength of viscosity can be graphically represented by a curve in that plane that is the locus of all points where the driving and damping timescales are equal---this defines a critical frequency as a function of temperature. Since the driving timescale decreases with frequency, stars above the critical frequency curve have one or more unstable $r$-modes, while stars below it are stable and stars on it are marginally stable.  Figure~\\ref{runaway} plots examples of critical frequency curves in the temperature range $10^8$~K$ \\ll T \\ll 10^{10}$~K appropriate for the LMXB emission scenario. (The needed range is higher than most observed temperatures of LMXBs in quiescence because of a thermal runaway; see below.) At low temperatures the damping in Fig.~\\ref{runaway} is taken to be dominated by shear viscosity in a boundary layer between the solid crust and fluid core (from Ref.~\\cite{Lindblom:2000gu}, augmented by a constant relative crust-core velocity in the range discussed in Ref.~\\cite{Levin:2000vq}). There are many other possible curves for the low temperature part of the plot, corresponding to more complicated damping mechanisms such as turbulence~\\cite{Wu:2000qy} or superfluid magnetoviscous effects~\\cite{Kinney:2002mq}, but they generally share the qualitative property of decreasing with temperature. (Some low-temperature curves lie above the observed range of spins entirely, but the observed spins of low-mass x-ray binaries are easier to explain if the curve lies in the region indicated~\\cite{Bildsten:1998ey, Andersson:1998qs, Chakrabarty:2003kt}.) At high temperatures the damping is probably dominated by bulk viscosity, either from the Urca process (perturbation of $\\beta$-equilibrium) or from nonleptonic processes involving strange particles such as hyperons. The Urca process, which requires no exotic particles, does not affect the critical frequency curve for $T < 10^{10}$~K and thus does not show up in Fig.~\\ref{runaway}. Thus the top plot in Fig.~\\ref{runaway} shows a critical frequency curve for a neutron star, and the bottom plot shows such a curve for a star with hyperons (with the high-temperature part of the curve derived from Ref.~\\cite{Lindblom:2001hd}). This type of plot looks encouraging for gravitational wave detection because there is plenty of room above the curve for stars to be unstable and thus emitting gravitational waves.  However, the gravitational wave emission duty cycle could be much smaller than 100$\\%$ due to a thermal runaway~\\cite{Levin:1998wa, Spruit:1998mk}. This happens generically when the critical frequency decreases with temperature. In that case, plotted at the top of Fig.~\\ref{runaway}, a star will execute a loop as shown and radiate only during the time it spends above the curve. A stable star which begins at the bottom left of the loop is spun up by accretion until it moves above the critical frequency and the instability is triggered. The shear from the growing $r$-modes then causes the star to heat up, moving it rightwards on the loop. As the temperature rises, so does the rate of neutrino cooling, causing the star to drop down toward the critical frequency again at a temperature of a few times $10^9$~K. After falling below the critical frequency, the $r$-mode heating is removed and the star drifts leftward along the bottom of the loop until returning to its initial position. Although the time it takes for a star to complete the loop is dominated by the accretion rate, the timescale for gravitational wave emission depends mainly on the saturation amplitude of $r$-mode oscillations. For a saturation amplitude ($\\alpha$ in the notation of Ref.~\\cite{Lindblom:1998wf}) of order unity, the duty cycle for gravitational wave emission is of the order $10^{-6}$~\\cite{Levin:1998wa}. Whereas the duty cycle can be as high as about $30\\%$ for the lowest predicted values of the saturation amplitude ($\\alpha\\simeq10^{-5}$), for typical estimates of the saturation amplitude ($\\alpha\\simeq10^{-4}$) the duty cycle is only of order $10^{-1}$~\\cite{Heyl:2002pe}. Advanced LIGO will be able to detect at most one LMXB (Sco X-1) without narrowbanding (and hurting its ability to see other sources), or 6--7 LMXBs with narrowbanding around a series of different frequencies~\\cite{Cutler:2002me}. The small number of detectable systems and the fact that the timescale for a star to complete a loop is much longer than a human lifetime mean that a duty cycle of order $10^{-1}$ or less is pessimistic for gravitational wave searches for the $r$-modes.  \\begin{figure} \\centerline{\\includegraphics[width=2.8in]{fig1a.eps}} \\centerline{\\includegraphics[width=2.8in]{fig1b.eps}} \\caption{ \\label{runaway} Critical frequency curves are given (qualitatively) by the dashed lines. The top plot includes no bulk viscosity due to hyperons or other strange particles. In this case an accreting neutron star traversing the loop indicated undergoes a thermal runaway and has a low gravitational radiation duty cycle. The bottom plot includes hyperon bulk viscosity. In this case the thermal runaway is blocked, and an accreting star is a source of persistent gravitational waves as it remains in equilibrium at the last arrowhead. } \\end{figure}  If the critical frequency increases with temperature as in the bottom plot of Fig.~\\ref{runaway}, the thermal runaway can be blocked. A rapidly accreting star in an LMXB can have a duty cycle of order unity for emission of gravitational radiation as it sits on the curve or makes small peregrinations about it~\\cite{Wagoner:2001kx}. (It could also keep emitting, although less detectably, for some time after the rapid accretion shuts off~\\cite{Reisenegger:2003cq}.) This could happen for stars which exhibit a rise in the critical frequency curve at a low enough temperature. The rise is typical for stars where high bulk viscosity processes involving hyperons~\\cite{Wagoner:2002vr, Wagoner:2003vi} or strange quarks~\\cite{Andersson:2001ev} are at work. Thus the case for the gravitational wave emission scenario is in fact strengthened by high bulk viscosity from processes which are fundamentally quark-quark interactions. One might expect similar critical frequency curves to arise for stars containing other forms of strange matter such as a kaon condensate or mixed quark-baryon phase, but this has yet to be investigated. The thermal runaway is blocked if the increase in temperature (width of the loop in Fig.~\\ref{runaway}) is enough to take the star from the negatively sloped part of the instability curve to the positively sloped part. The larger the saturation amplitude of $r$-modes the greater is the increase in the temperature. Whether the star makes the jump to the positively sloped curve also depends on which cooling mechanisms are operative, and on the shape of the low-temperature instability curve, which in turn depends on which damping mechanism dominates at low temperatures. Notwithstanding the wide range of estimates for $\\alpha$, the shape of the negatively sloped curve, and the cooling mechanisms, the question of the $r$-modes in LMXBs as a persistent source of gravitational radiation comes down to asking whether there is a rise in the instability curve around $10^9$~K or lower.   Our purpose in this paper is to revisit the question of whether there is a rise in the critical frequency curve around $10^9$~K or lower, as in the bottom of Fig.~\\ref{runaway}, leading to persistent gravitational wave emission. We focus (for now) on neutron stars containing hyperons (hyperon stars), because hyperons are in some sense the most conservative and robust of the many proposals for exotic matter in the cores of neutron stars: Some properties of hyperons can be measured in the laboratory, both in vacuum and in the environment of a light nucleus, which combined with astronomical observations allows one to constrain some of the many uncertainties in building an equation of state~\\cite{Glendenning:1997wn, Lackey:2005tk}. Equations of state which allow for hyperons generally produce them at densities relevant for neutron stars, about twice nuclear density and up.  We synthesize and extend results of previous work on this topic. In arguing the case for persistent gravitational wave emission, Wagoner~\\cite{Wagoner:2002vr, Wagoner:2003vi} and Reisenegger and Bonacic~\\cite{Reisenegger:2003cq} base their critical frequency curves on bulk viscosity coefficients obtained by combining the results of Lindblom and Owen~\\cite{Lindblom:2001hd} (hereafter LO) and Haensel, Levenfish, and Yakovlev~\\cite{Haensel:2001em} (hereafter HLY). The LO and HLY viscosities were obtained by different calculations and produced somewhat different results. LO used a detailed self-consistent model of a multi-component fluid described by relativistic mean field theory~\\cite{Glendenning:1997wn}, but since they were primarily concerned with the high-temperature regime appropriate to newborn neutron stars they treated the effect of superfluidity with a rough approximation. LO also made some errors which resulted in a bulk viscosity coefficient a factor of 20 or more too high. HLY were more careful with superfluidity, using consistent damping factors in the collision integrals, but instead of evaluating reaction rates and thermodynamic derivatives within relativistic mean field theory they used ``order of magnitude estimates'' of some quantities which resulted in more than an order of magnitude disagreement with the microscopic results of LO (and with Jones~\\cite{Jones:2001ya}). We correct some mistakes in LO and combine their self-consistent microphysical model with a more careful treatment of superfluidity similar to that of HLY. We also treat the macroscopic physics more carefully than LO, including the effect of rotation on stellar structure which in some cases can significantly affect damping timescales and critical frequency curves. (HLY made only order of magnitude estimates of mode damping timescales and did not plot critical frequency curves.)  We also address the question ``Does the viability of persistent gravitational wave emission require fine tuning of parameters?'' Many of the parameters that go into building the equation of state have significant uncertainties, and those that go into computing reaction rates and bulk viscosities are even more uncertain. We investigate the viability of persistent gravitational wave emission with respect to variation of several microphysical numbers such as hyperon coupling constants and the superfluid bandgap, and even which nonleptonic reaction is most important (taking into account the results of van Dalen and Dieperink~\\cite{vanDalen:2003uy}). We also investigate the dependence on the mass of the star, since cooling observations~\\cite{Yakovlev:2004iq} and timing of radio pulsars in binaries~\\cite{Nice:2005fi, Ransom:2005ae} indicate a wider mass range than the traditionally assumed clustering around 1.4~$M_\\odot$. (Cooling observations also might be interpreted to favor the existence of strange particles such as hyperons, although the data still can be fit by exotic cooling from purely nucleonic matter.)  The organization of the rest of the paper is as follows. In Sec.~II we describe the microphysical model which leads to the equation of state and ultimately the macroscopic coefficient of bulk viscosity. In Sec.~III we plot the critical frequency curves for a range of neutron star masses and microphysical parameters. In Sec.~IV we summarize our findings and discuss possible improvements. ",
        "conclusions": "We have extended previous investigations of $r$-modes in accreting hyperon stars in LMXBs as persistent sources of gravitational waves, focusing on the bulk viscosity which is needed to prevent thermal runaway. We have used improved microphysics compared to previous treatments, and have accounted for the most important macroscopic correction due to rotation of the star.  We find that persistent gravitational wave emission is quite robust. Even the stiffest hyperonic equations of state in relativistic mean field theory produce enough damping to stop the runaway and persistently radiate, although some of them require neutron stars somewhat more massive than 1.4~$M_\\odot$. Stars below about 1.3~$M_\\odot$ are not likely to be persistent sources regardless of the equation of state. The mass thresholds are somewhat more favorable for lower superfluid bandgaps than for higher bandgaps. Other details of the microphysics are found to be considerably less important. Our results seem robust for typical values of the crust-core viscosity and of saturation amplitude of $r$-modes, though this issue requires further investigation. Stars with high masses and stiff equations of state could exist in thermal and torque equilibrium at temperatures down to $10^7$~K.  One possible avenue for substantial improvement is the hydrodynamics. We used a fluid expansion in the dissipation integrals which is a reasonable approximation for a Newtonian normal fluid, but not for a superfluid which can be much more complicated due to multiple components and entrainment. Work is underway to deal with the superfluid problem in general (\\cite{Andersson:2005pf} and references therein)."
    },
    "0512/astro-ph0512088_arXiv.txt": {
        "abstract": "We study the equation of state, polarization and radiation properties for nonideal, strongly magnetized plasmas which compose outer envelopes of magnetic neutron stars. Detailed calculations are performed for partially ionized hydrogen atmospheres and for condensed hydrogen or iron surfaces of these stars. ",
        "introduction": "Neutron stars can be considered as natural laboratories for studying the properties of matter under extreme physical conditions (see, e.g., \\cite{ST83}). In their cores the density $\\rho$ may exceed $10^{15}$ g cm$^{-3}$ and the temperature $T$ lies typically between $10^7$ and $10^9$~K, but their thermal emission depends on the properties of the outer envelopes with lower $\\rho$ and $T$. The temperature decreases in the heat-blanketing envelopes (at $\\rho \\lesssim 10^8$ g cm$^{-3}$) to $T\\sim T_\\mathrm{eff}$ near the radiative surface (at $\\rho \\sim10^{-3}$--$10^6$ g cm$^{-3}$, depending on the stellar parameters), where $T_\\mathrm{eff}$ is the \\emph{effective surface temperature} related to the thermal flux through the Stefan's law. In these envelopes, the magnetic field strength $B$ may reach $10^{15}$ G; most often $B\\sim 10^{11} - 10^{13}$ G. To calculate spectra of neutron-star emission that can be compared with the observations, one should take into account nonideality of the plasma in the envelopes and the effects of strong magnetic fields on the properties of the plasma and electromagnetic radiation.  In the cases where $T_\\mathrm{eff}\\sim\\mbox{a few}\\times10^5$--$10^6$~K (characteristic of middle-aged neutron stars, i.e., of age $\\sim10^4$--$10^6$ yr) and $B\\lesssim10^{13}$~G, the spectrum of thermal emission from a neutron star forms in an stellar \\emph{atmosphere}. In a strong magnetic field, the state of matter can change significantly (for review, see, e.g., \\cite{Lai01}). For instance, at $T_\\mathrm{eff}\\sim10^6$~K, an atmosphere composed of hydrogen might be treated as fully ionized if the magnetic field were zero, but there can be a significant fraction of bound H atoms, if $B\\gtrsim10^{11}$~G. At lower $T$ (characteristic of older stars) and sufficiently high $B$, \\emph{magnetic condensation} may occur, resulting in formation of a solid or liquid metallic surface, composed of a strongly coupled Coulomb plasma and covered by a thin gaseous atmosphere.  Although the possible importance of these effects has been realized long ago, they were not included in neutron-star atmosphere models until recently (e.g., \\cite{Pavlov95} and references therein). Early considerations of partial ionization in the magnetized neutron-star atmospheres (e.g., \\cite{Miller92,RRM}; also reviewed briefly in \\cite{ZP02}) were not very reliable because of oversimplified treatments of atomic physics and nonideal plasma effects in strong magnetic fields. However, in the last decade the studies of dense nonideal plasmas in the strong magnetic fields made a spectacular progress, that now allows one to calculate the spectrum and polarization of electromagnetic radiation formed in an extended partially ionized hydrogen atmosphere or emitted from a condensed metallic surface of a magnetic neutron star. In this paper, we briefly review these achievements. ",
        "conclusions": "We have briefly reviewed the main effects of strong magnetic fields on the EOS, opacities and properties of electromagnetic radiation in the surface layers of neutron stars. In order to calculate realistic spectra of thermal radiation from neutron stars, one must carefully take these effects into account. We expect that comparison of the calculated models with observations will help to improve the constraints on neutron star parameters (for example, their effective temperatures, see \\cite{CSP}) and thus provide a powerful tool for testing the theories of superdense matter in neutron star interiors.  \\ack The work of GC was partially supported by the CNRS French-Russian grant PICS 3202. The work of AYP was partially supported by the RLSS grant 1115.2003.2 and the RFBR grants 05-02-16245, 03-07-90200, and 05-02-22003. WCGH is supported by NASA through Hubble Fellowship grant HF-01161.01-A awarded by the STScI, which is operated by the AURA, Inc, for NASA, under contract NAS~5-26555."
    },
    "0512/astro-ph0512331_arXiv.txt": {
        "abstract": "We use line-of-sight velocity information on the filamentary emission-line nebula of NGC\\,1275 to infer a dynamical model of the nebula's flow through the surrounding intracluster gas.  We detect outflowing gas and flow patterns that match simulations of buoyantly rising bubbles from which we deduce that some of the nebula filaments have been drawn out of NGC\\,1275. We find a radial gradient of the ratio [N{\\sc ii}]$\\lambda$6584/\\ha\\ which may be due to a variation in metallicity, interactions with the surrounding intracluster medium or a hardening of the excitation mechanism. We find no preferred spatial correlation of stellar clusters within the filaments and there is a notable lack of [O{\\sc iii}]$\\lambda$5007 emission, therefore it is unlikely that the filaments are ionized by stellar UV. ",
        "introduction": "NGC\\,1275 is the central galaxy of the X-ray luminous Perseus cluster (A426), which has bright centrally peaked X-ray emission, and a cool core with a central temperature of a third the virial temperature \\citep{Schmidt, Sanders}. Cavities in the X-ray emission are observed in a number of locations surrounding the central galaxy \\citep{Fabian2003b}. Where these coincide with GHz radio emission they are interpreted as bubbles filled with relativistic plasma, which has been injected into the intracluster medium (ICM) by the central engine.  Cavities with no observed radio emission have been described as `ghost bubbles', and are thought to have originated from an earlier epoch of activity in the central engine.  \\citet{Minkowski} discovered that the nebula of NGC\\,1275 comprises of two distinct emission-line systems: a high-velocity system (8200\\,km\\ps), identified as a disrupted foreground galaxy \\citep{Boroson} at least 60\\,kpc in front of NGC\\,1275 \\citep{Gillmon}, and a low-velocity system (5265\\,km\\ps). Although the high-velocity system lies directly in front of NGC\\,1275, the emission lines are easily distinguished in wavelength from those of the low-velocity system, and clearly indicate photoionization by hot, young stars \\citep{Kent}.  The low-velocity system associated with the central galaxy NGC\\,1275 is the focus of this work. It is known to extend over 100\\,kpc in a large array of thin filaments \\citep{Lynds,conselice}.  Whilst the nebula is extremely luminous--4.1$\\times10^{42}$\\erg\\ps\\ in \\ha\\ and [N{\\sc ii}] \\citep{Heckman}, with a total line luminosity probably 20 times that in \\ha-- the power source remains unknown.  Ionization by the central active galactic nucleus (AGN) residing in NGC\\,1275 can be ruled out as the dominant source of power for the extended nebula on the grounds that the \\ha\\ luminosity does not decrease with distance from the nucleus \\citep{RodAndy}, although it may be important in the luminous inner regions.  Ionization by hot young stars is an attractive option as it is a local mechanism, but the line ratios are drastically different to those seen in H{\\sc ii} regions \\citep{Kent}. Models of heating by X-rays from the ICM have been put forward \\citep{Donahue91}, as well as conduction from the ICM \\citep{Donahue}, shocks \\citep{SabraShieldsFlip} and turbulent mixing layers \\citep{Crawfordmixinglayers}.  Soft X-ray emission is associated with some of the optical filaments of NGC\\,1275 \\citep{Fabian2003}. The filaments are less luminous in the X-ray than the optical/UV by up to two orders of magnitude implying they are not excited by X-radiation.  The soft X-ray emission indicates an interaction between the warm filaments and the hot ICM, possibly via heat conduction.  Large deposits of molecular hydrogen have been discovered in the central regions of NGC\\,1275 \\citep{Krabbe,Donahue} similar to other central cluster galaxies with emission-line nebulae \\citep{Edge}. Recently, molecular hydrogen was observed in the outer filaments of NGC\\,1275 \\citep{Hatch}, indicating gas at 2000\\,K exists within the hot ICM at radii of over 25\\,kpc.  The origin of the filaments remain a mystery. Current theories include condensing gas from the ICM in the form of a cooling flow \\citep{FNC, Heckman, Donahue91}, gas accreting from previous mergers \\citep{Braine}, the explosive expulsion of gas from NGC\\,1275 \\citep{Burbidge} or gas drawn out \\citep{Fabian2003}. The filaments are very thin, long and the majority are radial. Submerged within the ICM, they enable us to constrain the level of turbulence in the ICM and argue for a laminar flow. As the ICM moves it may drag the warm optically emitting gas, thus the filaments can act as streamlines tracing the flow direction\\citep{Fabian2003}.   None of these problems are exclusive to NGC\\,1275 as extended emission-line nebulae are commonly found surrounding other massive galaxies in the centre of X-ray bright `cool cores', where the X-ray emission is centrally peaked \\citep{Crawford1999}.  This work presents new spectroscopic data that explore the kinematic and line-emission properties of the nebula that surrounds NGC\\,1275. After analysing and interpreting the kinematics we put forward a dynamical model of the nebula and discuss the origin of the filaments. The redshift of NGC 1275 is 0.0176, which using H$_{0}$=70\\,kms$^{-1}$Mpc$^{-1}$, gives 1\\,kpc$\\simeq$2.7\\,arcsec. ",
        "conclusions": "The kinematic data presented in section \\ref{sec:kinematics} rules out any dynamical models of purely infalling filaments.  The low velocities and line-widths presented here argue strongly against inflow, as one would expect both these parameters to rise sharply toward the centre of the nebula \\citep{Heckman,Donahue91}. The most conclusive evidence lies in the velocity structure of the Northern and Northwest filaments.  The lower half of the Northern filament is redshifted,whilst the upper section is blueshifted, thus the upper section of the filament is moving in the opposite direction to the material below: part of the filament is flowing away from the galaxy, whilst the other part is flowing into the galaxy.  In order to explain the outflow we appeal to the models of \\citet{BohringerM87,Churazov} and \\citet{Fabian2003} in which the radio emitting plasma from the AGN forms bubbles in the ICM, which detach to become buoyant and rise, dragging cool material from the galaxy below. The Northwest filaments lie directly underneath a ghost bubble \\citep{Fabian2003}. In addition to the morphological resemblance noted by \\citet{Fabian2003}, the kinematic signature of these filaments matches simulations of gas flow under a buoyantly rising bubble \\citep{Reynolds}, including details such as gas above the bubble moving in the opposite direction to the filaments below.  The data suggests the filaments are outflowing and therefore their origins lie within the galaxy. NGC\\,1275 contains a large reservoir of cold molecular gas \\citep{Krabbe,Donahue} that can fuel these filaments. As we observe part of the Northern filament falling back into the galaxy it is possible that the fate of all filaments lie in an eventual return to the galaxy.  However, if the filament falls back in segments as observed in the Northern filament, this would stretch and possible narrow the filament, making them more susceptible to evaporation by the ICM.  The detection of [O{\\sc i}] and [N{\\sc i}] indicates the presence of warm atomic hydrogen, and warm molecular hydrogen has been found in the outer filaments \\citep{Hatch}. It is possible the filaments holds a significant amount of cooler, so-far undetected gas.  In the manner proposed above, the central galaxy can efficiently lose mass and pollute the ICM with metals.  We report a radial variation of the [N{\\sc ii}]$\\lambda$6584/\\ha\\ ratio, indicating either progressive hardening of the excitation mechanism close to NGC\\,1275, or a variation in the nitrogen/oxygen abundance.  Although NGC\\,1275 is surrounded by numerous stellar clusters, we have presented details of HST images which show there is no preferential association with the optical filaments (Figs. \\ref{hstloops} and \\ref{hstfiladetail}).  As [O{\\sc iii}] emission (commonly found in H{\\sc ii} regions) is also lacking in the filaments, it is unlikely to be powered by or the birthplace of hot young stars. The central region exhibits [O{\\sc iii}]$\\lambda$5007 line emission, in contrast to the outer nebula, therefore an additional hard excitation source may be influential in the central region."
    },
    "0512/astro-ph0512107_arXiv.txt": {
        "abstract": "{We have carried out a study of the orthogonal polarisation mode behaviour as a function of frequency of \\NPulsars pulsars, using average pulsar data from the European Pulsar Network (EPN). Assuming that the radiation consists of two 100\\% polarised completely orthogonal superposed modes we separated these modes, resulting in average pulse profiles of each mode at multiple frequencies for each pulsar. Furthermore, we studied the frequency dependence of the relative intensity of these modes. We found in many pulsars that the average pulse profiles of the two modes differ in their dependence on frequency. In particular, we found that pulse components that are dominated by one mode tend to increase in intensity with increasing frequency with respect to the rest of the profile. ",
        "introduction": "Single pulse studies of the position angle (PA) of linearly polarised radiation from a number of pulsars show that it is built up of two modes of polarisation, which are separated in angle by 90 degrees \\citep[e.g.,][]{Manchester75, Stinebring84, Rankin88a, Gil91, Gil92}. It is believed that these orthogonally polarised modes (OPM) reflect the eigenmodes of the magneto-active plasma in the open magnetic field-lines above the pulsar polar cap. Three modes of wave propagation are allowed in this region, one of which is the ordinary sub-luminous mode, which cannot escape the pulsar magnetosphere due to Landau damping and is thus of no interest to the present work. The remaining two modes are the ordinary super-luminous mode (O-mode) which is polarised in the local plane defined by the external magnetic field and the wave-vector, and the extraordinary mode (X-mode), which is polarised perpendicular to this plane \\citep{Arons86,Petrova01}. According to \\citet{Barnard86} refraction can separate the X- and O-modes beams, which have different indexes of refraction, by many beam-widths. The tendency for average profiles to have constant widths above a critical frequency is then explained by this separation occurring above a critical height \\citep{Sieber75}. Two conditions are required for this: (1) the radio emission mechanism has to be broadband in frequency over a narrow range of heights above the stellar surface and (2) the gradient in the plasma density has a transverse component to the radial direction \\citep{McKinnon97}. The independent propagation of the two modes in the open flux zone can produce the abrupt orthogonal transitions in polarisation position angle that are commonly observed in studies of individual pulse polarisation \\citep{Manchester75}. The merging of the beams might also account for depolarisation of pulse average profiles with increasing radio frequency \\citep{McKinnon97, Hoensbroech97}. Alternatively, the OPM transitions might not be due to refraction, but due to switching between significant and insignificant conversion of O-mode into X-mode \\citep{Petrova01}.  One of the most important questions about OPM is whether the polarisation modes are disjoint or superposed. For disjoint modes the polarisation of the observed radiation is given by either one of the two modes at each point in time.  Superposed modes occur simultaneously and the polarisation of the observed electromagnetic radiation is given by the vectorial addition of the Stokes parameters $Q$, $U$ and $V$ of both modes. \\citet{Cordes78}, who first posed this question based on the polarisation from \\object{PSR B2020+28}, assumed the modes to be disjoint and stated three arguments for this assumption: (1) the degree of polarisation is fairly steady from one pulse to the next, (2) if the modes are superposed, one might expect occasional sign changes in the complex value $Q+iU$ when the linear polarisation is low, due to random fluctuations in the relative strength of the two modes and (3) the instantaneous mode of polarisation seems correlated with the total intensity, which is consistent with disjoint occurrence of the modes.  However, \\citet{Stinebring84} found that the modes appear to be superposed by studying sensitive polarisation observations of a number of pulsars. It was found that average values of the fractional linear polarisation were small at longitudes where both polarisation modes occurred with nearly equal frequency. The same effect was later found in the instantaneous values of the fractional linear polarisation in single pulse observations of \\object{PSR B2020+28} \\citep{McKinnon98}. This phenomenon is present on timescales of hundreds of microseconds to hours, which is what one would expect for superposed OPM. For disjoint modes one would expect the fractional polarisation to be either high or low, depending on the fractional polarisation of the active mode.  Assuming that the polarisation modes are superposed, that they are completely polarised and completely orthogonal, the average intensity of the individual polarisation modes as a function of longitude can be determined from the Stokes parameters of the average pulsar signal \\citep{McKinnon00}. In this paper we make these assumptions and determine the individual polarisation profiles of \\NPulsars pulsars as a function of both longitude and frequency (using data from the EPN database). In section~\\ref{sec:Method} we give a description of our analysis, the results of which are shown in section~\\ref{sec:Results}. In section~\\ref{sec:Discussion} we discuss the results and finally we give our conclusions in section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} We have determined the average pulse profiles of two polarisation modes for \\NPulsars pulsars at multiple frequencies, assuming that the modes are completely orthogonal and 100\\% polarised. We find many cases where each polarisation mode can be wholly associated with one component in the average intensity and cases where a 90$\\deg$ phase jump in the PA is associated with the two polarisation modes changing dominance. This suggests that the two polarisation modes are successfully determined in these cases. In other cases where there is a (gradual) jump in the PA of less than 90$\\deg$ it would appear that one of the assumptions enabling the determination of the modes is not always valid. The spectra of the ratio of the integrated intensity of the polarisation modes show that in many cases there is a trend for one of the modes to become stronger with increasing frequency. This can be both the mode which we have defined as strongest, as well as the weakest mode. We also find that the average profiles of the modes often differ from each other at different frequencies. In particular, we find that when a component is dominated by one mode at low frequency it tends to increase in intensity with increasing frequency with respect to the rest of the profile. Furthermore, out of \\NPulsars pulsars there are 2 pulsars that clearly show one polarisation mode changing its longitudinal position with frequency. From our results we cannot determine whether the longitudinal transitions of one polarisation mode into the other at a fixed frequency are the result of the X- and O-modes beams being separated due to refraction \\citep{Barnard86}, or due to switching between significant and insignificant conversion of O-mode into X-mode \\citep{Petrova01}. Both explanations can also account for some of the complex changes that occur in the average profiles of the polarisation modes over frequency, since plasma waves with different frequency will traverse different paths through the pulsar magnetosphere. In the model of \\citet{Petrova01} complex changes in the average profile of the polarisation modes over frequency implies that there are differences between the plasma number densities in polarisation limiting regions corresponding to emission at different frequencies. Following the technique as described in \\citet{Petrova03} and assuming radius to frequency mapping, calculation of the plasma density distributions from polarisation profiles at different frequencies will then yield changes in the plasma density both as a function of longitude as well as altitude."
    },
    "0512/astro-ph0512382_arXiv.txt": {
        "abstract": "\\noindent  We present results from a survey of a 1300 arcmin$^2$ region of the Orion B South molecular cloud, including NGC 2024, NGC 2023, and the Horsehead Nebula (B33), obtained using the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope. Submillimeter continuum observations at 450\\,$\\mu$m and 850\\,$\\mu$m are discussed.  Using an automated algorithm, 57 discrete emission features (``clumps'') are identified in the 850\\,$\\mu$m map. The physical conditions within these clumps are investigated under the assumption that the objects are in quasi-hydrostatic equilibrium.  The best fit dust temperature for the clumps is found to be $T_d = 18 \\pm 4\\,$K, with the exception of those associated with the few known far infrared sources residing in NGC 2024. The latter internally heated sources are found to be much warmer. In the region surrounding NGC 2023, the clump dust temperatures agree with clump gas temperatures determined from molecular line excitation measurements of the CO molecule. The bounding pressure on the clumps lies in the range $\\log(k^{-1}\\,P\\ {\\rm cm}^{3}\\,{\\rm K}^{-1}) = 6.1 \\pm 0.3$. The cumulative mass distribution is steep at the high mass end, as is the stellar Initial Mass Function. The distribution flattens significantly at lower masses, with a turn-over around 3 -- 10\\,$M_\\odot$. ",
        "introduction": "This paper continues an effort to quantify the necessary pre-conditions for star formation via observations of the continuum emission from cold dust at submillimeter wavelengths. Such observations allow the determination of the mass spectrum and physical characteristics of the population of cold (typically 10--30\\,K), dense concentrations of dust, or ``clumps'', each of which is expected to be a site of active, or eventual, star formation. It is also well known that stars form in groups (see e.g. Clarke et al. 2000, Elmegreen et al. 2000) and that these in turn originate from a hierarchy of dust clumps for which the mass spectrum is closely similar to the stellar initial mass function (IMF; see e.g. Motte et al. 1998; Johnstone et al. 2000b, hereafter Paper II; Johnstone et al. 2001, hereafter Paper III; Motte et al. 2001). The recently launched Spitzer Space Telescope will observe many star-forming regions in the infrared and thus provide information on which dust and gas clumps contain deeply embedded protostars.  Together, the submillimeter continuum dust maps and the infrared images will yield statistics on the number of prestellar, Class 0, and Class I protostars and significantly enhance our understanding of the lifetime in each stage and the changes in natal environment with protostellar phase.  Previous papers in this series on large area mapping at 850\\,$\\mu$m have discussed image reconstruction techniques (Johnstone et al. 2000a, hereafter Paper I), and results for the Ophiuchus cloud (Paper II) and the northern part of the Orion B region (Paper III).  In the present paper we turn our attention to the southern part of the Orion B molecular cloud (Lynds 1630). This region includes a number of objects well known to both amateur and professional astronomers. In particular, the Horsehead Nebula, B33 (see, e.g., Pound et al. 2003), is seen in silhouette against the bright background of the HII region IC\\,434 near the three bright stars of Orion's Belt. At a distance of {approximately 400\\,pc [Anthony-Twarog (1982) measured a distance of 390\\,pc]} the Orion B South region also contains the young star clusters NGC\\,2023 and NGC\\,2024, and as part of the greater Orion Molecular Cloud is well-known as one of the nearest extensive regions of active star formation. An optical image showing most of the region discussed in this paper, and its relationship with the ionizing stellar system $\\sigma$~Orionis, is given by Abergel et al. (2002).  One of the first systematic searches for dense molecular cores in the Orion B region was carried out by Lada et al. (1991a; hereafter LBS) using the CS J = 2--1 transition, for which the effective critical density is about $10^4$~cm$^3$ (Evans 1999).  With a beamwidth of 1.8$'$ a total of 42 dense cores were found to be concentrated in two relatively small regions of the Orion B molecular cloud; 19 of these lie within the region surveyed in the present paper. Seven of the latter objects were subsequently mapped by Launhardt et al. (1996) at 1300\\,$\\mu$m wavelength with an angular resolution of 12$''$, comparable to that of the present observations with the JCMT. From these surveys it is evident that active star formation is confined to two rather limited areas within the Orion B region.  The region has been surveyed at many wavelengths. {Early 2.2\\,$\\mu$m observations of a large part of the molecular cloud made by Lada et al. (1991b) identified four embedded star clusters, and} subsequent widefield imaging of the stellar component is available via the 2MASS images (Carpenter 2000). The distribution of relatively warm dust was mapped by IRAS, although for our purposes with limited angular resolution, and in far infrared emission at 138 and 205\\,$\\mu$m (Mookerjea et al. 2000). The molecular component has been surveyed by a number of authors, e.g. Miesh and Bally (1994) in the $^{13}$CO J = 1--0 transition, and by Kramer et al. (1996) with lower angular resolution in the $^{13}$CO and $^{12}$CO J = 2--1 and 3--2 transitions.  The Horsehead Nebula in particular has been extensively investigated. Pound et al. (2003) have obtained an interferometric image in the CO J = 1--0 transition of the Horsehead Nebula with an angular resolution of 10$''$. Abergel et al. (2003) combine mid-infrared ISOCAM data in the LW2 and LW3 bands with CO, $^{13}$CO and C$^{18}$O 1--0 and 2--1 observations to model the transient heating of the outer edge of the dust cloud by photodissociating radiation from the relatively nearby O9.5V system $\\sigma$~Orionis. Teyssier et al. (2004) have obtained detailed observations of simple cyclic and linear hydrocarbons, which they show to be distributed similarly to CO within the western edge of the nebula.  Finally, early results from submillimeter-wavelength mapping of the Horsehead Nebula region have been reported by Sandell et al. (2001); these latter archival data have been included in the dataset for the present work, which covers a total area about 1300 arcmin$^2$ in extent. ",
        "conclusions": "\\label{s_sum}  We have presented images of the southern part of the Orion B region obtained at submillimeter wavelengths of 850 and 450\\,$\\mu$m, covering an area 1300 arcmin$^2$ in extent.  Utilizing the techniques developed in Papers I, II, and III, we have identified 57 independent dust concentrations, or ``clumps'', within this region. The majority of these clumps are concentrated in two star-forming regions, NGC\\,2024 and NGC\\,2023; a third minor grouping appears in the Horsehead Nebula cloud. Most of the dust mass is concentrated in NGC\\,2024.  We have been able to derive estimated masses and temperatures for the majority of the clumps by noting that they can be modelled as Bonnor-Ebert spheres; that is, they can be approximated by almost constant density models with low internal velocities. On this basis the typical clump temperature is found to be 18$\\pm$4~K; this value is in reasonable agreement with that (21$\\pm$9~K) derived from the spectral indices obtained from the 850 and 450\\,$\\mu$m data.  Thus for most of the clumps a temperature of 20~K is a good working approximation.  However, 9 clumps are centrally concentrated enough that they may be collapsing, and 4 of these objects are particularly bright at 850\\,$\\mu$m. All of the latter lie within NGC\\,2024 and have temperatures which are considerably higher than that for the majority of the clumps.  Using CO isotopomer spectral line data obtained for NGC2023 we have attempted to obtain independent estimates of temperatures and masses for the clumps in this particular region. With one exception the results obtained agree with those derived from the continuum data. It can be difficult, however, to clearly relate CO structures with equivalent continuum clumps. This problem becomes more acute at the lower mass end of the clump distribution, where continuum sensitivity is limited and individual clumps overlap.  The mass function of submillimeter clumps found in Orion B South is steep, in agreement with the results in other nearby regions (Motte et al. 1998; Paper II; Paper III; Motte et al. 2001). The fraction of mass in each CS-identified core which has fragmented in observable submillimeter clumps is $\\sim 20$ percent."
    },
    "0512/astro-ph0512457_arXiv.txt": {
        "abstract": "Using a large sample of MgII absorbers with $0.4<z<2.2$ detected by Nestor \\e\\ (2005) in the Early Data Release of the SDSS, we present new constraints on the physical properties of these systems based on two statistical analyses: (i) By computing the ratio between the composite spectra of quasars with and without absorbers, we measure the reddening effects induced by these intervening systems; and (ii) by stacking SDSS images centered on quasars with strong MgII absorption lines and isolating the excess light around the PSF, we measure the mean luminosity and colors of the absorbing galaxies.  This statistical approach does not require any spectroscopic follow up and allows us to constrain the photometric properties of absorber systems. ",
        "introduction": "Since it was realized that strong metal absorption lines detected in quasar spectra are due to intervening galaxies (e.g. Bahcall \\& Spitzer 1969, Boksenberg \\& Sargent 1978, Bergeron \\& Boiss\\'e 1991), significant progress has been made in our understanding of absorber systems.  A number of studies have constrained some of their physical properties such as temperature, abundances and kinematics (for a review see for example Churchill's contribution). However, the precise nature of absorber systems is still elusive and more observational constraints are needed in order to improve our understanding on how these systems trace various structures such as galactic disks, halos, winds, etc. This would allow us to include and use absorbers in our global picture of galaxy formation and evolution.  In this contribution we illustrate how imaging and spectroscopic data from the Sloan Digital Sky Survey (SDSS) can be used to constrain the nature of absorber systems. Having selected a population of quasars with strong absorbers as well as a population of reference quasars (\\S 2), we construct their composite spectra and isolate the reddening effects due to the presence of absorbers (\\S 3). In \\S 4, we show how the photometric properties of absorbing galaxies can be constrained by measuring the statistical excess of light present around quasars in SDSS images. We summarize the results in \\S 5. ",
        "conclusions": "Using a sample of over 800 MgII systems with $0.4<z<2.2$ we have presented two statistical analyses constraining the properties of these absorber systems. After having carefuly selected a population of reference quasars, we have measured the mean reddening curve induced by these intervening systems by computing the ratio between the composute spectra of quasars with and without absorbers.  We do not find any presence of the 2175 \\AA\\ bump present in the Milky way and LMC reddening curves and we find the SMC reddening curve to provide a good fit to the measured reddening. Assuming the mean MgII absorbers to have the same metallicity and dust-to-gas ratio as the SMC, we measured the mean hydrogen column density of these systems to be $\\sim10^{21}$ atoms cm$^{-2}$.  In addition, we have constrained the mean photometric properties of MgII absorbing galaxies by measuring the statistical light excess present around quasars with absorbers. This signal is detected at 5.5, 7.1 and 9.2$\\sigma$ in the $g$, $r$ and $i$-band, respectively and such an excess is not seen for the reference sample. This light contribution behaves as expected as a function of scale and redshift, and the measured colors indicate that the mean MgII absorbing galaxies is a late-type galaxy.  The results presented in this contribution are based on EDR data. Future analyses will provide much more accurate constraints and will allow us to investigate the absorber properties as a function of redshift and for different ion species."
    },
    "0512/astro-ph0512511_arXiv.txt": {
        "abstract": "{  Using recent spectroscopic observations, we show that the triple system AO\\,Vel with an eclipsing BpSi primary is in fact a remarkable quadruple system formed by two double-lined spectroscopic binaries with components close to the ZAMS. All available data have been reanalyzed to derive proper orbital parameters for both binary systems and to calculate absolute parameters of the eclipsing system. For the first time, direct determination of the radius and the mass have been obtained for a BpSi star. The study of the physical parameters of this unique system is especially important since it can be used as a test of evolutionary models for very young stars of intermediate mass. ",
        "introduction": "Upper main sequence stars with chemical peculiarities and large-scale organized magnetic fields constitute about 10\\% of A and B-type stars. To understand the physics of such stars, it is most important to identify the origin of their magnetic field. Two main hypotheses have been proposed: either these stars have retained the primordial field present at their formation, the ohmic decay being longer than the main sequence stellar lifetime (fossil field theory), or the field is generated and maintained by a contemporary dynamo. The first hypothesis implies that strongly magnetic stars must be found on the ZAMS (as well as over the whole width of the main sequence), while this is not necessarily true according to the second hypothesis. Hubrig et al. (\\cite{hu00}) suggested that magnetic stars tend to accumulate in the middle of the main sequence strip if their mass is below 3\\,M$_\\odot$ (which would challenge the fossil field theory), while more massive magnetic stars seem more evenly distributed over the whole width of the main sequence, including near the ZAMS (Hubrig et~al.\\ \\cite{hu05}). However, the chemical peculiarities of these stars distort their colour indices, so that their effective temperatures suffer from an uncertainty which also propagates to other stellar parameters like e.g. the surface gravity (North \\ \\cite{n98}). Therefore, it is most important to determine the stellar parameters of magnetic stars in the most fundamental way, especially as the extent to which their evolution within the main sequence is identical to that of their non-magnetic siblings is not known.  The most suitable systems for the accurate assessment of  fundamental stellar parameters are double-lined eclipsing binaries. While a few double-lined spectroscopic binary systems (SB2) containing an Ap star of mass  below 3\\,M$_\\odot$ are currently known, the rate of close binaries is much smaller among magnetic Bp stars (Gerbaldi et~al.\\ \\cite{g85}; Carrier et al.\\ \\cite{cnub02}; North \\& Debernardi \\cite{nd04}) and only two double-lined eclipsing binaries with a Bp component, namely AO\\,Vel (= \\object{HD\\,68826}) and HD\\,123335, are known to date (e.g., Clausen, Gim\\'enez \\& van Houten \\cite{c95}, from here on CGvH; Hensberge et~al.\\ \\cite{he04}). In this paper, we perform a spectroscopic analysis of the eclipsing SB2 system AO\\,Vel consisting of two late B-type stars with a BpSi primary component (Bidelman \\& MacConnell \\cite{bi73}), i.e.\\ of a type generally displaying rather large spectroscopic variations, so that one can infer that there are significant surface inhomogeneities. Numerous four-colour photometric observations have been obtained from 1974 to 1991 with the Danish 50\\,cm telescope at La Silla (CGvH). Times of mid-eclipse for a total of 16 primary and 10 secondary minima have been determined from all available photoelectric observations and photometric elements and apsidal motion parameters have been determined from complete $uvby$ light curves. This study revealed the presence of a third component in a wide orbit with an orbital period of 25.6\\,yr. The orbital period of the eclipsing system appeared to be 1.58~days and the orbit was found only slightly eccentric, with a quite fast apsidal motion (U=56.8\\,yr). Very recently, Wolf \\& Zejda (\\cite{wz05}) published seven additional times of minima obtained from the ASAS-3 database (Pojmanski \\cite{asas}) and new photoelectric $UBV$ observations. As a result, they have improved the apsidal motion parameters and the orbit of the third body, for which they found a period of 33\\,yr.  In spite of the considerable amount of observational data collected for decades, absolute dimensions of the components remained unknown since no accurate radial velocity observations had ever been performed for this system. ",
        "conclusions": "Using our recent spectroscopic observations we discovered that the triple system AO\\,Vel with an eclipsing BpSi primary is in fact a remarkable quadruple system close to the ZAMS. The analysis of our data as well as all available published data revealed completely different orbital parameters for this system compared to previous works. We found that both spectroscopic binaries move in a wide eccentric orbit ($e=0.29$) with a period of 41\\,yr and this orbit is fully consistent  with the observed difference between the center-of-mass radial velocities of the two spectroscopic binaries. The calculated orbit imposes a lower limit to the total mass of the binary C+D: M$_C$+M$_D$ $\\ge$ 4.9\\,M$_\\odot$. The analysis of the stellar parameters and of the evolutionary state of the components reveals their extreme youth. Since no magnetic field measurements have been performed for this system, we have applied for VLT time to obtain spectra in circular polarized light with FORS\\,1, which is equipped with polarization analyzing optics.  We also plan to obtain additional high resolution, high signal-to-noise spectra of this system at different orbital phases, in order to improve the determination of the stellar parameters and atmospheric elemental abundances of all companions. This unique system is especially important to test evolutionary models of very young stars of intermediate mass. Detailed observations during primary eclipses will provide strong constraints on the distribution of the abundance anomalies on the surface of the A component, complementing the standard Doppler imaging technique based on rotation."
    },
    "0512/astro-ph0512505_arXiv.txt": {
        "abstract": "We have performed an abundance analysis for F- and G- dwarfs of the Galactic thick disk component. A sample of 176 nearby ($d \\leq 150$ pc) thick disk candidate stars was chosen from the $Hipparcos$ catalogue and subjected to a high-resolution spectroscopic analysis. Using accurate radial velocities combined with  $Hipparcos$ astrometry, kinematics $(U, V,$ and $~W)$ and Galactic orbital parameters were computed. We estimate the probability for a star to belong to the thin disk,  the thick disk or the halo. With a probability $P \\geq 70$\\%  taken as certain membership, we assigned 95 stars to the thick disk, 13 to the thin disk, and 20 to the halo. The remaining 48 stars in the sample cannot be assigned with reasonable certainty to one of the three components.  Abundances of C, O, Na, Mg, Al, Si, Ca, Sc, Ti, V, Cr, Mn, Fe, Co, Ni, Cu, Zn, Y, Ba, Ce, Nd, and Eu have been obtained. The abundances for thick disk stars are compared with those for thin disk members from Reddy et al. (2003). The ratios of $\\alpha$-elements (O, Mg, Si, Ca and Ti) to iron for thick disk disk stars show a clear enhancement compared to thin disk members in the range $-0.3 <$ [Fe/H] $ < -1.2$. There are also other elements --  Al, Sc, V, Co, and possibly  Zn -- which show enhanced ratios to iron in the thick disk relative to the thin disk. The abundances of Na, Cr, Mn, Ni, and Cu (relative to Fe) are very similar for thin and thick disk stars. The dispersion in abundance ratios [X/Fe] at given [Fe/H] for thick disk stars is consistent with the expected scatter due to measurement errors, suggesting a lack of `cosmic' scatter.  A few stars classified as  members of the thick disk by our kinematic criteria show thin disk abundances. These stars, which appear older than most thin disk stars, are also, on average, younger than the thick disk population. They may have originated early in the thin disk history, and been subsequently scattered to hotter orbits by collisions. The thick disk may not include stars with [Fe/H] $> -0.3$. The observed compositions of the thin and thick disks seem to be consistent with models of galaxy formation by hierarchical clustering in a $\\Lambda$CDM universe. In particular, the distinct abundance  patterns observed in the thin and thick disks, and the chemical homogeneity of the thick disk at different galactocentric distances favor a scenario in which the majority of thick-disk stars were formed {\\it in situ}, from gas rich merging blocks. ",
        "introduction": "Stars of the solar neighbourhood are overwhelmingly members of the Galactic disk, with a small admixture of halo stars. The assignment of a star to the disk or the  halo is based on  differences in chemical composition and kinematics.  The local disk population is subdivided into stars of the thin disk and others belonging to the thick disk, with chemical composition and kinematics again playing a role in effecting this subdivision.  The modern division of the disk into  the thin and thick disk was proposed by Gilmore \\& Reid (1983). Star counts led them to divide the disk population in the solar neighbourhood into a thin disk with a scale height of 300 pc and a thick disk with the much greater scale height of 1450 pc. Thin disk stars outnumber thick disk stars by about twenty to one  in the Galactic plane. Many other estimates of scale heights and relative densities of thin and thick disk populations now exist (e.g., Buser et al. 1999; Ojha 2001; Cabrera-Lavers, Garz\\'on \\& Hammersley 2005; Juric et al. 2005). The thick disk stars are generally older than most thin disk stars. The metallicity distribution of the thick disk population is shifted to lower values relative to the distribution for the thin disk by about 0.5 dex. Although both distributions can be reasonably approximated by Gaussians with a FWHM of roughly 0.5 dex, the thick disk includes a tail at lower metallicities. In contrast to the thin disk stars, which orbit the Galactic centre on nearly circular orbits, the thick disk stars are on moderately elliptical orbits that typically reach higher distances from the plane. Thick disk stars also revolve around the Galactic center slower than those in the thin disk.  The origin of the thick disk has occasioned much debate. Keys to the origin lie within the kinematics and the composition of the thick disk stars. A number of recent spectroscopic studies have set out to compare the chemical compositions of thick and thin disk stars. This avenue was explored by Gratton et al. (1996) and Fuhrmann (1998). Gratton et al. showed that O/Fe ratios for thick disk stars are distinctly different from thin disk stars but similar to halo stars. Fuhrmann confirmed this based on Mg abundances and showed a clear cut difference in the Mg/Fe ratio for thick and thin F-G dwarf stars of the same [Fe/H]\\footnote{Here and throughout the paper we use the so-called {\\it bracket} notation to indicate chemical abundance ratios of two elements X and Y: [X/Y]  $\\equiv \\log$ N(X)/N(Y) $- \\log $ (N(X)/N(Y))$_{\\odot}$}. These studies stimulated several investigations of elemental abundances in samples of thick and thin disk stars -- see, for example, Prochaska et al. (2000), Feltzing, Bensby, \\& Lundstr\\\"{o}m (2003), Reddy et al. (2003, hereafter Paper I), Bensby, Feltzing, \\& Lundstr\\\"{o}m (2003, 2004), Bensby et al. (2005), and Mishenina et al. (2004). Although the pattern of abundance differences between thick and thin disk is emerging, many details remain obscure, largely, one suspects, because these investigations cover small numbers of thick disk stars: Prochaska et al. considered ten, Bensby and colleagues analyzed 36, and Mishenina et al. less than 30 stars. Considering that the thick disk may span a range of 1 dex in [Fe/H], these samples, even when combined, are probably too small to define in detail the differences between  compositions of  thick and thin disk stars over their full range in [Fe/H], even if the two disk components were themselves  chemically homogeneous as a function of metallicity. Additionally, different definitions of what constitutes a thick disk star have been adopted by different authors.  Tens of thousands of thick-disk stars at a few kpc from the plane have been spectroscopically observed as part of the Sloan Digital Sky Survey (York et al. 2000; Adelman-McCarthy et al. 2005). These spectra, however, have a much lower resolving power than the surveys mentioned above, and although they may yield abundance ratios for a number of metals that produce strong spectral lines, they have, to this date, been used to derive iron abundances only (Allende Prieto et al. 2005).  Exploration of the chemical compositions of local thin disk stars is now well advanced. In particular, several surveys have investigated many elements in F-G dwarfs whose spectra are amenable to quantitative analysis. Our recent study of 26 elements in 181 F-G dwarfs (Paper~I) was the precursor for the work presented in this paper. The vast majority of the 181 stars belong to the thin disk, as judged (see below) by their kinematics. Our Paper~I sample may be combined with other large samples to which thin disk stars are the major contributor: e.g., Edvardsson et al. (1993) for 189 stars, and Chen et al. (2000) for 90 stars.  A key result of our 2003 survey was the finding that `cosmic' scatter in an abundance ratio X/Fe at a given Fe/H for thin disk stars was less than the small measurement errors. Here, we  apply the same analytical techniques to a large sample of thick stars for which the cosmic scatter and, indeed, the form of the run of [X/Fe] with [Fe/H] was not known at the outset of this project. There were clear indications of the sign and magnitude of some abundance differences between thick and thin disk, as recognized by Fuhrmann (1998), and Prochaska et al. (2000) and further examined by Bensby and colleagues, and by Mishenina et al. (2004).  The present survey provides abundances for 23 elements from C to Eu for  176 stars in the solar neighbourhood, of which 95 are attributed to the thick disk. The full sample is introduced in the next section. The observations and abundance analysis are based closely on Paper~I's approach is described in Section 3 and 4. Full results, and comparisons with other studies are given in Section 5. Chemical evolution of the thick disk and evolution of the Galactic disk are discussed in Section 6. This section includes discussion on stars which have thick disk kinematics and thin disk abundances (TKTA), disk heating, and merger scenarios. Concluding remarks are  given in Section 7. ",
        "conclusions": "On the basis of the  available abundance data on the thin and thick disks and of the published simulations of disk formation through mergers in a $\\Lambda$CDM universe, scenario (C) appears to be a plausible leading explanation for the origin of the thick and thin disks. In terms of a continued exploration of the observational frontiers, there is, for example, a need for (i) a detailed abundance analysis of stars apparently attributable to an extension of the thick disk to metallicities below [Fe/H] of $-$1, i.e., the so-called metal weak thick disk, which comprises only one per cent of the thick disk (Martin \\& Morrison 1998); (ii) a more complete investigation of the four dimensional space (U,V,W,[Fe/H]), e.g., the pursuit of stars at low $W_{LSR}$ with [Fe/H] less than about $-0.7$, stars, which, if present, would be assigned to the thin disk; (iii) analysis of a larger  sample of stars with [Fe/H] greater than about $-$0.3 and well-determined kinematics is needed to confirm or deny the presence of the knee in the thick disk Mg-like abundances connecting to the thin disk abundances for the most metal-rich stars; (iv) a larger sample of thick disk stars is needed to determine radial and vertical gradients in compositions of thick disk stars. The vertical gradient, if any, appears to be very shallow (e.g; Bensby et al. 2005, Allende Prieto et al. 2005).  The realisation that thick and thin disk stars of the same [Fe/H] differ in composition and the strong suggestion that thick and thin disk stars span overlapping but distinctly different ranges in [Fe/H] (see, for example, Schuster et al. 2005) has consequences for the cottage industry providing models of chemical evolution of the Galactic disk, especially for models of the solar neighbourhood. The industry standard supposes that chemical evolution as represented by a plot of [X/Y] vs [Y/H], where Y is traditionally taken to Fe or sometimes O, is a continuous process from the halo to the disk, i.e., initially metal-free gas experiences star formation leading to the halo stars, collapse of gas to a disk with enrichment from stellar nucleosynthesis and possibly continuing infall of gas leads to a steady continous chemical evolution. No account is taken  of the fact that the disk has the two components - thin and thick - from different origins and most probably covering different metallicity ranges. The time has come to change the industry standard!"
    },
    "0512/astro-ph0512169_arXiv.txt": {
        "abstract": "We present 118 new optical redshifts for galaxies in 12 clusters in the Horologium-Reticulum supercluster (HRS) of galaxies.  For 76 galaxies, the data were obtained with the Dual Beam Spectrograph on the 2.3m telescope of the Australian National University at Siding Spring Observatory.  After combining 42 previously unpublished redshifts with our new sample, we determine mean redshifts and velocity dispersions for 13 clusters, in which previous observational data were sparse.  In six of the 13 clusters, the newly determined mean redshifts differ by more than 750 \\kms\\ from the published values.  In the case of three clusters, A3047, A3109, and A3120, the redshift data indicate the presence of multiple components along the line of sight. The new cluster redshifts, when combined with other reliable mean redshifts for clusters in the HRS, are found to be distinctly bi-modal. Furthermore, the two redshift components are consistent with the bi-modal redshift distribution found for the inter-cluster galaxies in the HRS by \\citet{fle05}. ",
        "introduction": "The Horologium-Reticulum supercluster (HRS) is an extended region of high galaxy density \\citep{sha35,luc83,ein03,fle05}, covering $\\sim$150 square degrees of sky at a mean redshift of $\\sim$20,000 \\kms. The HRS also contains more than 20 galaxy clusters \\citep{ein97,ein02}. As discussed in \\citet{hud99} and \\citet{ein01}, the HRS is the second largest mass concentration within $\\sim$300 Mpc, where it is only surpassed by the Shapley supercluster (SSC).  While the SSC has been extensively studied \\citep{qui95,qui00,dri99,dri04,bar98,bar00}, the HRS remains relatively unexplored.  Due to the potential importance of such a large-scale structure in the present-epoch universe, we have embarked on a redshift survey to provide a comprehensive mapping of the HRS. Our initial results, which contain 547 galaxy redshifts in the {\\it inter-cluster} regions of the HRS, are reported in \\citet[][hereafter Paper I]{fle05}. A key result from Paper I is that the distribution of inter-cluster galaxies is separated into two distinct redshift components.  On the other hand, the published mean redshifts for 21 galaxy clusters in the HRS do not exhibit such a bi-modal distribution. The differing results between the cluster and inter-cluster redshift distributions appear to contradict the view that galaxy clusters share the kinematics of the inter-cluster galaxy distribution as a result of their location at intersecting filaments of galaxies \\citep[e.g.,][]{van93,bon96,col99,col05}. However, the mean redshift for many of these clusters is based on fewer than four galaxy redshifts per cluster, i.e., sparse information. To clarify the distribution of cluster redshifts in the HRS, we have obtained new data for 12 clusters in which the previously published data were sparse.  The results of this program are reported below and, when combined with other previous redshift data, give an improved assessment of the distribution of cluster redshifts in the HRS. Throughout the paper, we adopt the following cosmological parameters: $\\Omega _m = 0.3$, $\\Omega _\\Lambda = 0.7$, and $H_o = 70 $ \\kms\\ Mpc$^{-1}$, which implies a spatial scale of 4.6 Mpc degree$^{-1}$ (77 kpc arcmin$^{-1}$) at the $\\sim$20,000 \\kms\\ mean redshift of the HRS. ",
        "conclusions": "We have obtained 76 new optical redshifts within 12 galaxy clusters of the Horologium-Reticulum supercluster (HRS). These observations, augmented by 42 previously unpublished redshifts, have led to the determination of more accurate cluster properties. Using the methods for calculating robust mean redshifts (location) and velocity dispersions (scale) described in BFG90, we have calculated mean redshifts and dispersions for 13 clusters, including A3109 for which no new observations are reported. The mean redshifts for several clusters have changed by at least 750 \\kms\\ (in 6/13 observed) from their previously reported values. In addition, three clusters are observed to consist of multiple components (A3047, A3109, and A3120). The new cluster redshift data have been compared to previously compiled redshift data for the inter-cluster galaxies in the HRS from \\citet{fle05}. Primarily, we now find consistency between the large-scale kinematic features of the clusters and the inter-cluster galaxies. Specifically, there is a principal kinematic axis in the HRS at a PA of $-$80\\degr\\ east from north, along which a systematic increase in redshift with position is observed for both clusters and inter-cluster galaxies. After this overall spatial-kinematic trend is removed, the distribution in redshift for both clusters and inter-cluster galaxies is distinctly bi-modal, with the two redshift peaks separated by $\\sim$3000 \\kms.  We thank the Australian National University and Mount Stromlo/Siding Spring Observatories for facilitating and supporting these observations. We also thank Clair Murrowood for her assistance with the observations, Ilana Klamer for supplying her unpublished 2dF data of A3104, and Bruce Peterson for the use of Mathams' thesis data. M. C. F. acknowledges the support of a NASA Space Grant Graduate Fellowship at the University of North Carolina-Chapel Hill. R. W. H. acknowledges grant support from the Australian Research Council. M. J. H. acknowledges support through IRGS Grant J0014369 administered by the University of Tasmania. A portion of this work was supported  by NSF grants AST-9900720 and AST-0406443 to the University of North Carolina-Chapel Hill. 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."
    },
    "0512/astro-ph0512443_arXiv.txt": {
        "abstract": "{We propose and test a new method based on Richardson-Lucy deconvolution to reconstruct three-dimensional gas density and temperature distributions in galaxy clusters from combined X-ray and thermal Sunyaev-Zel'dovich observations. Clusters are assumed to be axially symmetric and arbitrarily inclined with respect to the line-of-sight. No equilibrium assumption other than local thermal equilibrium is needed. We test the algorithm with synthetic observations of analytically modeled and numerically simulated galaxy clusters and discuss the quality of the density and temperature reconstructions in idealised situations and in presence of observational noise, deviations from axial symmetry and cluster substructure. We find that analytic and numerical gas density and temperature distributions can be accurately reconstructed in three dimensions, even if observational noise is present. We also discuss methods for determining the inclination angle from data and show that it can be constrained using X-ray temperature maps. For a realistic cluster and including observational noise the three-dimensional reconstructions reach a level of accuracy of about 15$\\%$. } ",
        "introduction": "In hierarchical models of structure formation, galaxy clusters are not only the most massive gravitationally bound objects in the Universe, but also the most recently forming. Numerous examples show that they are typically irregularly shaped and occasionally undergoing violent merger events. Cluster-sized dark-matter halos in simulations can often be well described as triaxial ellipsoids, but not as spheres \\citep{JI02.1}.  At the same time, observations of galaxy clusters are often interpreted based on spherically-symmetric models in hydrostatic equilibrium. The beta model \\citep{CA76.1} is still routinely being used for analyses of the X-ray emission and also of the amplitude of the thermal Sunyaev-Zel'dovich effect. Given the rapidly improving quality and diversity of cluster data, it appears timely to search for an algorithm which avoids the assumption of spherical symmetry and allows the joint analysis of different types of cluster data.  Several such algorithms have been proposed. \\cite{ZA98.1} suggested to base the reconstruction of axisymmetric, three-dimensional gravitational cluster potentials on the Fourier slice theorem, extrapolating Fourier modes into the ``cone of ignorance''. They applied their technique to simulated data and showed that it performs well \\citep{ZA01.1}. \\cite{DO01.3} followed a perturbative approach, and \\cite{LE04.1} proposed to adapt parameters of triaxial halo models, all by combining different data sets such as X-ray, (thermal) Sunyaev-Zel'dovich (SZ) and gravitational-lensing maps. A similar method was applied to data by \\cite{DE05.1}. An alternative approach based on the iterative Richardson-Lucy deconvolution was suggested by \\cite{RE00.1} and \\cite{RE01.1}. It aims at the gravitational potential, assumes only axial symmetry of the main cluster body, avoids extrapolations in Fourier space, and can easily be extended to include additional data sets.  In this paper, we develop the latter algorithm further. However, aiming at the potential would require us to assume a relation between the gas distribution and the gravitational field, which would be most conveniently given by hydrostatic equilibrium. But even ignoring this common equilibrium assumption, it should be possible to reconstruct the three-dimensional distributions of intra-cluster gas density and temperature by a joint analysis of X-ray and thermal SZ data.  We demonstrate here that this is indeed possible under the one simplifying assumption that the underlying three-dimensional distributions be axially (not spherically!) symmetric. The inclination of the symmetry axis can be arbitrary. We introduce the algorithm in Sect.~2 and apply it to the idealised case of an analytically modeled, ellipsoidal cluster without substructure in Sect.~3. Results obtained first without, then with observational noise are highly promising: both the three-dimensional density and temperature distributions are accurately reproduced. Noise suggests smoothing, and we describe a suitable smoothing algorithm.  We study the less-ideal case of a numerically-simulated galaxy cluster in Sect.~4. Here, axial symmetry is typically violated by the main cluster body, and substructures are present which further perturb the symmetry. Yet, faithful reconstructions are possible even in presence of realistic noise.  Section~5 finally describes how inclination angles can be constrained using temperature maps, and Sect.~6 summarises and discusses our results. ",
        "conclusions": "We propose a new method for deprojecting hot gas in galaxy clusters which combines X-ray and thermal Sunyaev-Zel'dovich effect observations and reconstructs three-dimensional density and temperature distributions. We start from the iterative deprojection algorithm suggested by \\cite{BI90.1} which employs Richardson-Lucy deconvolution and assumes axial symmetry of the physical quantity whose three-dimensional distribution shall be reconstructed from two-dimensional maps of its projection along the line-of-sight.  This approach does not restrict the orientation of the symmetry axis to be parallel to the line-of-sight, but the inclination angle between the symmetry axis and the line-of-sight is assumed to be known. This algorithm runs into problems when it is used to reconstruct strongly peaked distributions such as the X-ray emissivity of a galaxy cluster. There, one obtains spike-shaped artifacts through the centre of the reconstructed halo. We suppress the formation of such artifacts by introducing a regularisation scheme for the iterative corrections used in the deprojection. Then, we generalise this algorithm to simultaneously reconstruct three-dimensional distributions of several physical quantities by combining two-dimensional maps of projections which probe these three-dimensional distributions in different ways.  Here, we discuss how three-dimensional density and temperature distributions of the intra-cluster medium can be reconstructed from combined X-ray and thermal Sunyaev-Zel'dovich effect observations. We test the method using synthetic data of analytically modeled and of numerically simulated galaxy clusters and discuss the quality of the reconstructions and the impact of observational noise, cluster substructure and deviations from axial symmetry. For numerical clusters which are of course not strictly axisymmetric, we use one of the principal inertial axes as the ``symmetry'' axis for the deprojection.  Our main findings, if we neglect observational noise and assume that the inclination angle between the symmetry axis and the line-of-sight is known, are:  \\begin{itemize}  \\item Spike-shaped artifacts of the deprojection are efficiently suppressed by the regularisation of the iterative corrections.  \\item Densities and temperatures of the ICM of axisymmetric analytic clusters can be reconstructed very accurately from X-ray flux and Sunyaev-Zel'dovich effect maps. Errors are of the order of 1\\% unless the angle between the symmetry axis and the line-of-sight is small.  \\item The three-dimensional density and temperature distributions of hot gas in numerically simulated clusters, although not strictly axisymmetric, can still be reliably reconstructed. Relative errors reach roughly 10\\%. Smoothing of the X-ray flux and the Sunyaev-Zel'dovich effect maps can be used to suppress artifacts caused by subclumps.  \\item Accurate gas density and temperature profiles can be obtained from the reconstructions.  \\end{itemize}  We then add photon noise corresponding to a total number of $10^4$ observed photons to the X-ray and observational noise corresponding to a four-hour ALMA Band 3 observation to the Sunyaev-Zel'dovich effect maps, respectively. We smooth the maps before repeating the deprojections to suppress small-scale fluctuations which the algorithm would otherwise attempt to approximate and thereby reduce the quality of the reconstructions. From the repeated reconstructions of the analytic and numerical halos, this time including observational noise, we conclude:  \\begin{itemize}  \\item Gas densities and temperatures of axisymmetric analytic halos can also be efficiently reconstructed from maps that contain observational noise. The relative errors of the reconstructions are about 5\\% to 10\\%.  \\item The three-dimensional structure of the ICM of numerically simulated clusters can also be reliably reconstructed from X-ray flux and Sunyaev-Zel'dovich effect maps that contain observational noise. Relative errors reach roughly 15\\%.  \\item Accurate profiles can be obtained from the reconstructions.  \\item Five iterations are sufficient for the ICM deprojection. Using a larger number does not increase the quality of the reconstruction.  \\end{itemize}  For these deprojections, we assumed that the inclination angle between the symmetry axis and the line-of-sight is known beforehand. This will usually not be the case for observations of real clusters. In principle, one could try to find the inclination angle by deprojecting the cluster assuming different values of the inclination angle and comparing the original X-ray and Sunyaev-Zel'dovich effect maps to those expected from the reconstructed halo. They should match best if the correct inclination is used for the deprojection. Unfortunately, the minima of the deviations of the maps as a function of the assumed inclination angle are quite broad and not always centred on the correct value. However additional data which are independent of the X-ray flux and the Sunyaev-Zel'dovich effect maps, can provide information on the inclination angle. We show that high-quality emission-weighted temperature maps which become more and more routinely available, can constrain the inclination angle of a cluster's symmetry axis to values for which the quality of the reconstruction is close to its optimum."
    },
    "0512/astro-ph0512396_arXiv.txt": {
        "abstract": "{ We report on the discovery of a $z = 3.16$ Lyman-$\\alpha$ emitting blob in the Great Observatories Origins Deep Survey (GOODS) South field. The discovery was made with the VLT, through narrow-band imaging. The blob has a total Ly$\\alpha$ luminosity of $\\sim10^{43}$~erg~s$^{-1}$ and a diameter larger than 60~kpc. The available multi-wavelength data in the GOODS field consists of 13 bands from X-rays (Chandra) to infrared (Spitzer). Unlike other known Ly$\\alpha$ blobs, this blob shows no obvious continuum counter-parts in any of the broad-bands. In particular, no optical counter-parts are found in deep HST/ACS imaging. For previously published blobs, AGN (Active Galactic Nuclei) or ``superwind'' models have been found to provide the best match to the data. We here argue that the most probable origin of the extended Ly$\\alpha$ emission from this blob is cold accretion onto a dark matter halo. ",
        "introduction": "Narrow-band surveys for Lyman-$\\alpha$ (Ly$\\alpha$) emitting galaxies at high redshift have recently revealed a number of luminous (up to $5 \\cdot 10^{43}$ erg s$^{-1}$), very extended (from a few times ten kpc to more than 150 kpc) Ly$\\alpha$-emitting objects, so-called Ly$\\alpha$ ``blobs'' (Fynbo et al. 1999; Keel et al. 1999; Steidel et al. 2000; Francis et al.  2001; Matsuda et al. 2004; Palunas et al. 2004; Dey et al. 2005; Villar-Martin et al. 2005). At least three mechanisms have been suggested as energy sources for Ly$\\alpha$ blobs. These are: \\emph{i)} hidden QSOs (Haiman \\& Rees 2001; Weidinger et al. 2004, 2005), \\emph{ii)} star formation and superwinds from (possibly obscured) starburst galaxies (Taniguchi et al. 2001; Ohyama et al. 2003; Mori et al. 2004; Wilman et al. 2005), and \\emph{iii)} so-called cold accretion (Haiman, Spaans \\& Quataert 2000; Fardal et al. 2001; Keres et al. 2004; Maller \\& Bullock 2004; Birnboim \\& Dekel 2003; Sommer-Larsen 2005; Dijkstra et al. 2006(a,b); Dekel \\& Birnboim 2006). Cooling flows are phenomena observed in galaxy clusters for more than a decade (Fabian 1994). These are explained by gas which is cooling much faster than the Hubble time through X-ray emission in the centres of the clusters. However, cooling emission from a galaxy, or a group sized halo can be dominated by Ly$\\alpha$ emission (e.g. Haiman, Spaans \\& Quataert 2000; Dijkstra et al. 2006(a,b)). In this \\emph{Letter} we present the discovery of a Ly$\\alpha$ blob at redshift $z \\approx 3.16$ located in the GOODS South field, which we argue is the first piece of evidence for cold gas accretion onto a dark matter halo.  Throughout this paper, we assume a cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega _{\\rm m}=0.3$ and $\\Omega _\\Lambda=0.7$. All magnitudes are in the AB system. ",
        "conclusions": "We have here reported the results of an extensive multi-wavelength investigation of a redshift $z = 3.16$ Ly$\\alpha$~emitting blob discovered in the GOODS South field. The blob has a diameter larger than 60~kpc diameter and a total luminosity of $\\mathrm{L}_{\\mathrm{Ly}\\alpha} \\sim 10^{43}$~erg~s$^{-1}$. Deep HST imaging show no obvious optical counterpart, and the lack of X-ray or IR emission suggest there are no AGN or dusty starburst components associated with at least the centroid of the blob. Two galaxies within a $10''$ radius have photometric redshifts consistent with the redshift of the blob, but follow-up spectroscopy is needed to establish if there is a connection. We have run simulations of Ly$\\alpha$ surface brightness arising from cold accretion and found that such extended Ly$\\alpha$ emission may be explained by accretion along a filament onto a galaxy group sized dark matter halo. Another possibility is that such emission in very short lived, i.e. significantly shorter than the 2.5 Myr resolution of our simulation. We argue that other previously suggested origins of Ly$\\alpha$ blobs (hidden AGN and ``super-winds'') can be ruled out in this case due to the lack of detected continuum counter-parts. Hence, though our cold accretion simulation cannot perfectly match our data, it is the only explanation that is plausible. Our results combined with the fact that previously studied blobs appear to be caused by superwinds and/or AGN in turn implies that the energy sources for blob Ly$\\alpha$ emission are diverse.  \\vspace{-0.1cm}"
    },
    "0512/astro-ph0512113_arXiv.txt": {
        "abstract": "We report the discovery of a double-double radio galaxy (DDRG), J0041+3224, with the Giant Metrewave Radio Telescope (GMRT) and subsequent high-frequency observations with the Very Large Array (VLA). The inner and outer doubles are aligned within $\\sim$4$^\\circ$ and are reasonably collinear with the parent optical galaxy. The outer double has a steeper radio spectrum compared with the inner one. Using an estimated redshift of 0.45, the projected linear sizes of the outer and inner doubles are 969 and 171 kpc respectively. The time scale of interruption of jet activity has been estimated to be $\\sim$20 Myr, similar to other known DDRGs.  We have compiled a sample of known DDRGs, and have re-examined the inverse correlation between the ratio of the luminosities of the outer to the inner double and the size of the inner double, $l_{in}$. Unlike the other DDRGs with $l_{in}$$\\gapp$50 kpc, the inner double of J0041+3224 is marginally more luminous than the outer one. The two DDRGs with  $l_{in}$$\\lapp$few kpc have a more luminous inner double than the outer one, possibly due to a higher efficiency of conversion of beam energy as the jets propagate through the dense interstellar medium. We have examined the symmetry parameters and find that the inner doubles appear to be more asymmetric in both its armlength and flux density ratios compared with the outer doubles, although they appear marginally more collinear with the core than the outer double. We discuss briefly possible implications of these trends. ",
        "introduction": "One of the important issues concerning galaxies is the duration of their active galactic nuclei (AGN) phase and whether such periods of activity are episodic. In the currently widely accepted paradigm, activity is believed to be intimately related to the `feeding' of a supermassive black hole whose mass ranges from $\\sim$10$^6$ to 10$^{10}$ M$_\\odot$. Such an active phase may be recurrent with  an average total timescale of the active phases being $\\sim$10$^8$ to 10$^{9}$ yr (cf. Marconi et al. 2004, and references therein).  Of the galaxies harbouring an AGN, a small fraction appears to be luminous at radio wavelengths. For example, in the SDSS (Sloan Digital Sky Survey) quasars, $\\sim$8 per cent of the bright ones ($i<$18.5) are radio loud in the sense that the ratio of radio to X-ray flux exceeds unity (Ivezi\\'{c} et al. 2002, 2004). Although what physical conditions determine loudness still remains unclear, Nipoti, Blundell \\& Binney (2005) have suggested recently that this may simply be a function of the epoch at which the source is observed.  For the radio-loud objects, an interesting way of probing their history is via the structural and spectral information of the lobes of extended radio emission. Such studies have been used to probe sources which exhibit precession or changes in the ejection axis, effects of motion of the parent galaxy, backflows from hotspots as well as X-shaped sources and major interruptions of jet activity. For example, the radio galaxy 3C388 exhibits two distinct regions of emission separated by a jump in spectral index, which has been interpreted to be due to two different epochs of jet activity (Burns, Schwendeman \\& White 1983; Roettiger et al. 1994). A ridge of emission, reminescent of a jet but displaced towards the south of the nucleus in the radio galaxy 3C338 could also be due to intermittent jet activity (Burns, Schwendeman \\& White 1983). Other suggestions of distinct epochs of jet activity based on spectral index studies include Her A, where the bright inner regions have flatter spectra with a sharp boundary delineating it from the more extended lobe emission, especially in the western lobe (Gizani \\& Leahy 2003), and 3C310 where the inner components B and D have substanially flatter spectra than the surrounding lobes (van Breugel \\& Fomalont 1984; Leahy et al. 1986). An interesting example of different epochs of jet activity is the well-studied radio galaxy Cen A, where in addition to the diffuse outer lobes there are the more compact inner lobes and a northern middle lobe or NML (Burns, Feigelson \\& Schreier 1983; Clarke, Burns \\& Norman 1992; Junkes et al. 1993). Morganti et al. (1999) have detected a large-scale jet connecting the northern lobe of the inner double and the NML, and have suggested that the formation of the NML may be due to a `bursting bubble' in which plasma accumulated in the inner lobe bursts out through a nozzle. A lobe of emission on only one side of the nuclear region has also been seen in the Gigahertz Peaked Spectrum radio source B0108+388, and has been suggested to be a relic of a previous cycle of jet activity (Baum et al. 1990). However, the one-sidedness of the emission is puzzling (cf. Stanghellini et al. 2005), and it would be interesting to examine whether a `bursting bubble' model may also be applicable in such cases.  One of the more striking examples of episodic jet activity is when a new pair of radio lobes is seen closer to the nucleus before the `old' and more distant radio lobes have faded. Such sources have been christened as `double-double' radio galaxies (DDRGs) by Schoenmakers et al. (2000, hereinafter referred to as S2000). They proposed a relatively general definition of a DDRG as a double-double radio galaxy consisting of a pair of double radio sources with a common centre. S2000 also suggested that the two lobes of the inner double should have an edge-brightened radio morphology to distinguish it from knots in a jet. In such sources the newly-formed jets propagate outwards through the cocoon formed by the earlier cycle of activity rather than the general intergalactic or intracluster medium, after traversing through the interstellar medium of the host galaxy. Approximately a dozen or so of such DDRGs are known in the literature (S2000; Saripalli et al. 2002, 2003; Schilizzi et al. 2001; Marecki et al. 2003). We have included objects such as 3C236 (Schilizzi et al. 2001) and J1247+6723 (Marecki et al. 2003) in this category since the principal difference is only in the size of the inner double which is $\\lapp$few kpc. It is important to identify more DDRGs not only for understanding episodic jet activity and examining their time scales, but also for studying the propagation of jets in different media. For example, to explain the edge-brightened hotspots in the inner doubles of the DDRGs, Kaiser, Schoenmakers \\& R\\\"{o}ttgering (2000) have suggested that warm (T$\\sim$10$^4$K) clouds of gas in the intergalactic medium are dispersed over the cocoon volume by surface instabilities induced by the passage of the cocoon material.  In this paper we report the discovery of a new DDRG, J0041+3224, identified from observations made with the Giant Metrewave Radio Telescope (GMRT) of candidate DDRGs from the B2 sample (Padrielli, Kapahi \\& Katgert-Merkelijn 1981; Saikia et al. 2002). Our candidates were identified by comparing the large- and smaller-scale images made either by us or those available in the literature. J0041+3224 was reported by Padrielli et al. (1981) to be double-lobed with an angular size of 32 arcsec.  They identified the radio source to be associated with a galaxy with a visual magnitude of 20.0 and located at RA 00$^h$ 41$^m$ 46.$^s$11, Dec: +32$^\\circ$ 24$^\\prime$ 53.$^{\\prime\\prime}$8 in J2000 co-ordinates. There is no measured redshift of the galaxy. From the V magnitude$-$redshift diagram (Guiderdoni \\& Rocca-Volmerange 1987) we estimate the redshift to be $\\sim$0.45 and use this value for this paper. ",
        "conclusions": "We have reported the discovery of a new double-double radio galaxy, J0041+3224, with the GMRT and subsequent observations with the VLA. Using an estimated redshift of 0.45, the projected linear size of the outer double is 969 kpc. Large linear sizes are characteristic of most of the known DDRGs. The lobes of the outer double have steeper spectral indices compared with those of the inner double. The kinematic age of the outer double is $\\sim$3$\\times$10$^7$ yr, while the time scale of interruption of jet activity is $\\sim$20 Myr.  Unlike most DDRGs with $l_{in}$$\\gapp$50 kpc, the inner double of J0041+3224 is marginally more luminous than the outer one. S2000 reported an inverse correlation between the ratio of the luminosities of the outer and inner doubles, P$_{o:in}$, and $l_{in}$. For their sample of DDRGs the luminosity ratio, P$_{o:in}$, is as high as $\\sim$100 for $l_{in}$$\\sim$50 kpc and approaches unity when $l_{in}$ approaches values close to a Mpc, the ratio being always $>$1. Considering the more compact inner doubles with $l_{in}$ $\\lapp$ few kpc, namely J1006+3454 (3C236) and J1247+6723, the inner doubles are significantly more luminous than the outer ones. This suggests that in the early phase of the evolution of the inner double, where it is ploughing its way through the dense interstellar medium, conversion of beam energy into radio emission may be more efficient. In this case, the ratio of the luminosity of the outer to the inner double could increase with size before decreasing and approaching a value of about unity when $l_{in}$ approaches $\\sim$1 Mpc. We have re-examined the inverse correlation for sources with $l_{in}$$\\gapp$50 kpc with the addition of a few more sources and find the correlation to have a rank correlation co-efficient of $-$0.37, compared with a value of $-$0.57 for the objects in the sample of S2000.  We have compared the symmetry parameters of the inner and outer doubles and find that the inner doubles appear to be more asymmetric in both its armlength and flux density ratios compared with the outer doubles. Also, the asymmetries in the inner and outer doubles are not in the same sense. Although these trends need to be confirmed with larger samples, they are possibly a reflection of different environments due to different degrees of entrainment in the cocoons on opposite sides, coupled with any intrinsic jet asymmetries and effects of relativistic motion. However, the inner doubles appear marginally more collinear with the radio core than the outer doubles. This could arise due to the lobes of the outer double responding to large-scale density gradients, or motion of the parent optical galaxy during the two cycles of nuclear activity."
    },
    "0512/astro-ph0512439_arXiv.txt": {
        "abstract": "{ Scattering processes in the cosmic microwave background limit the propagation of ultra high energy charged particles in our Universe. For extragalactic proton sources resonant photopion production results in the famous Greisen-Zatsepin-Kuzmin (GZK) cutoff at about $4\\times10^{10}$ GeV expected in the spectrum observed on Earth. The faint flux of ultra high energy cosmic rays of less than one event per year and cubic kilometer and the large systematic uncertainties in the energy calibration of cosmic ray showers is a challenge for cosmic ray observatories and so far the GZK cutoff has not been unambiguously confirmed. We have investigated the possibility that the primaries of super-GZK events are strongly interacting neutrinos which are not subject to the GZK cutoff. For the flux of protons and neutrinos from extragalactic optically thin sources and a flexible parameterization of the neutrino-nucleon cross section we have analyzed the cosmic ray spectra observed at AGASA and HiRes taking also into account results from horizontal events at AGASA and contained events at RICE. We find that scenarios of strongly interacting neutrinos are still compatible with the data requiring a steep increase of the inelastic neutrino-nucleon cross section by four order of magnitude within one energy decade compared to the Standard Model predictions. We also discuss the impact of the preliminary cosmic ray spectrum observed by the Pierre Auger Observatory. } ",
        "introduction": "The origin and chemical composition of cosmic rays with energies above $10^9$ GeV - so called {\\it ultra high energy} (UHE) cosmic rays (CRs) - is still an open question in astroparticle physics. The large-scale isotropy~\\cite{Takeda:1999sg,Abbasi:2004ib} and also the shower characteristics~\\cite{Bergman:2004bk} of these events give reason to believe that CRs around the {\\it ankle} at about $10^{10}$ GeV are dominated by extragalactic protons. As charged particles, these protons are subject to scattering processes in the intergalactic photon gas. It has been shown~\\cite{Berezinsky:2002nc} (see also Refs.~\\cite{Fodor:2003bn,Berezinsky:2004fk,Ahlers:2005sn,Berezinsky:2005cq}) that the CR spectrum between $10^{9}$ GeV and $4\\times10^{10}$ GeV is in a remarkably good agreement with the propagated flux from extragalactic proton sources using a simple injection spectrum with two free parameters. In this case, the ankle can be identified as an $e^+e^-$ pair production {\\it dip} of protons scattering off photons of the cosmic microwave background (CMB) together with a {\\it pile-up} of protons due to photo-pion production.  Soon after the discovery of the CMB in 1965 it was pointed out by Greisen~\\cite{Greisen:1966jv}, Zatsepin, and Kuzmin~\\cite{Zatsepin:1966jv} that resonant photopion production of protons above $4\\times10^{10}$ GeV limits their propagation to about 50 Mpc.\\footnote{For heavier nuclei photo-disintegration in the CMB give a comparable or even stronger attenuation above this energy.} For an extra-galactic proton source this results in a cut-off expected in the CR spectrum. So far the experimental resolution for this feature is very poor due to low statistics and large systematic errors in energy calibration (cf.\\ Fig.~\\ref{alldataproc}). The next generation cubic-kilometer-sized detectors, notably the Pierre Auger Observatory (PAO) will soon have the necessary exposure to resolve this issue.  A continuation of the CR spectrum  beyond the GZK cutoff to extremely high energies (EHEs) seems only consistent with a proton dominance if the sources lie in our local cosmic environment. Only very few astrophysical accelerators can achieve the necessary energies~(see e.g.\\ Ref.~\\cite{Torres:2004hk} for a review) and so far none of the candidate sources have been confirmed. It has been speculated that decaying superheavy particles, which could be some new form of dark matter or remnants of topological defects, could be a source of UHE CRs, but also these proposals are not fully consistent with the gamma ray spectrum at lower energies~\\cite{Semikoz:2003wv}.  These missing parts of the CR {\\it puzzle} has led to speculations about a different origin of EHE CRs. In particular, a flux of neutral components would not be attenuated by the CMB and could be fueled by very distant sources. In the late 60s Berezinsky and Zatsepin~\\cite{Beresinsky:1969qj} proposed that cosmogenic neutrinos produced in photopion production of protons in the CMB could explain EHE events assuming a strong neutrino-nucleon interaction. The realization of such a behavior has been proposed abundantly in scenarios beyond the (perturbative) Standard Model (SM): e.g. arising through compositeness~\\cite{Domokos:1986qy,Bordes:1997bt,Bordes:1997rx}, through electroweak sphalerons~\\cite{Aoyama:1986ej,Ringwald:1989ee,Espinosa:1989qn,Khoze:1991mx,Ringwald:2002sw,Bezrukov:2003qm,Ringwald:2003ns,Fodor:2003bn,Han:2003ru}, through string excitations in theories with a low string and unification scale~\\cite{Domokos:2000dp,Domokos:2000hm,Burgett:2004ac}, through Kaluza-Klein modes from compactified extra dimensions~\\cite{Domokos:1998ry,Nussinov:1998jt,Jain:2000pu,Kachelriess:2000cb,Anchordoqui:2000uh,Kisselev:2003rz}, or through $p$-brane production in models with warped extra dimensions~\\cite{Ahn:2002mj,Jain:2002kf,Anchordoqui:2002it}, respectively (for recent reviews, see Refs.~\\cite{Fodor:2004tr,Anchordoqui:2005ey}).  Up to now, UHE cosmic neutrinos have been searched for in the Earth atmosphere (Fly's Eye~\\cite{FLYSEYE} and AGASA~\\cite{Yoshida:2001pw}), in the Greenland (FORTE~\\cite{FORTE}) and Antarctic ice sheet (AMANDA \\cite{AMANDA} and RICE~\\cite{RICE}), in the sea/lake (notably BAIKAL~\\cite{BAIKAL}), or in the regolith of the moon (GLUE~\\cite{GLUE}). A large neutrino-nucleon cross section is not necessarily in conflict with the very rare UHE neutrino events claimed by these experiments, if the transition from the weak SM interaction to a strong contribution of new physics proceeds over a sufficiently short energy range. Those neutrinos interacting very strongly will then be hidden in the nucleonic CR background . In our quantitative analysis of scenarios with strong neutrino-nucleon interactions we have used the search results from AGASA and RICE.  In this paper we will report on a recent statistical analysis~\\cite{Ahlers:2005zy} investigating scenarios with strong neutrino-nucleon interactions as an explanation of GZK excesses. For the flux of CRs and neutrinos from distant optically thin sources introduced in \\S 2 and a model for a strong neutrino-nucleon interaction introduced in \\S 3 we will show in \\S 4 the results of goodness-of-fit test for existing CR data from AGASA~\\cite{Takeda:2002at} and HiRes~\\cite{Bird:1994wp,Abbasi:2005ni,Abbasi:2005bw,HURL} as well as preliminary data from PAO~\\cite{Sommers:2005vs,PAO}. We also include the search results on neutrino-induced events from AGASA and RICE. Finally, we conclude in \\S 5.  \\FIGURE[t]{ \\begin{minipage}[t]{\\linewidth}\\center \\includegraphics[height=7cm]{./eps/alldataprocPAO.eps} \\caption{Cosmic ray spectrum observed by various experiments. Also shown is the size of the systematic error in energy calibration.} \\label{alldataproc} \\end{minipage}} ",
        "conclusions": "Our analysis shows that ultra high energy cosmic rays measured at AGASA and HiRes might be composed of extragalactic protons and strongly interacting neutrinos. The enhancement of the neutrino-nucleon cross section has to be rapid in order to agree with the experimental results from horizontal events at AGASA and contained events at RICE. We derived allowed regions of a flexible parameterization of the cross section using a goodness-of-fit test and showed the compatibility of our results with theoretical predictions from sphalerons, $p$-branes and string excitations. The good agreement between these fits with our simple parameterization indicate that out results (in particular Fig.~\\ref{csbound}) can be used as a {\\it quick test} for strongly interacting neutrino-nucleon cross sections.  Our results are based on the assumption that ultra high energy protons and neutrinos originate at optically thin sources. One should keep in mind that a more general injection spectrum of cosmic rays or optically thicker sources would certainly alter the prediction of the neutrino fluxes and accordingly our results for the range of neutrino-nucleon interactions. As a rule of thumb cosmic accelerators always produce cosmic rays, neutrinos and gamma rays with comparable luminosities. Also the top-down production of cosmic rays is always accompanied with high energy fluxes of other particles. Hence, the search for the origin of the highest energy cosmic rays involves essentially a multi-messenger analysis. It is in the focus of future neutrino experiments like ANITA, IceCube or the Pierre Auger Observatory (horizontal events) and gamma ray observatories like H.E.S.S.~\\cite{HESS} and MAGIC~\\cite{MAGIC} to resolve some of these model ambiguities in the analysis.  Notably the Pierre Auger Observatory combining the experimental techniques of AGASA and HiRes as a hybrid detector will have a great impact on scenarios with strongly interacting neutrinos. With a better energy resolution and a much higher statistics it will certainly help to clarify our picture on ultra high energy cosmic rays.   \\providecommand{\\href}[2]{#2}\\begingroup\\raggedright"
    },
    "0512/astro-ph0512325_arXiv.txt": {
        "abstract": "A number of deep, wide-field, near-infrared surveys employing new infrared cameras on 4m-class telescopes are about to commence. These surveys have the potential to determine the fraction of luminous dust-obscured quasars that may have eluded surveys undertaken at optical wavelengths. In order to understand the new observations it is essential to make accurate predictions of surface densities and number-redshift relations for unobscured quasars in the near-infrared based on information from surveys at shorter wavelengths. The accuracy of the predictions depends critically on a number of key components. The commonly used single power-law representation for quasar SEDs is inadequate and the use of an SED incorporating the upturn in continuum flux at $\\lambda \\sim 12\\,000\\,$\\AA \\ is essential. The presence of quasar host galaxies is particularly important over the restframe wavelength interval $8000 < \\lambda < 16\\,000\\,$\\AA \\ and we provide an empirical determination of the magnitude distribution of host galaxies using a low redshift sample of quasars from the SDSS DR3 quasar catalogue. A range of models for the dependence of host galaxy luminosity on quasar luminosity is investigated, along with the implications for the near-infrared surveys. Even adopting a conservative model for the behaviour of host galaxy luminosity the number counts for shallow surveys in the $K$ band increase by a factor of two. The degree of morphological selection applied to define candidate quasar samples in the near-infrared is found to be an important factor in determining the fraction of the quasar population included in such samples. ",
        "introduction": "There has been much debate in recent years regarding the existence, or otherwise, of a significant population of dust-obscured quasars that have eluded surveys based on an ultraviolet excess or optical colours. Since gas is required to fuel quasars and given the success of unified models, which invoke non-spherical structures that will certainly obscure the line-of-sight to the inner regions of certain active galactic nuclei (AGN), the existence of an obscured population of quasars is a natural consequence of our present understanding of AGN and quasars. Indeed, examples of obscured objects have been identified, from weakly reddened quasars, whose spectra show small deviations consistent with obscuration by dust at ultraviolet wavelengths (e.g. Richards et al. 2003), to more heavily obscured objects where emission at ultraviolet wavelengths is almost completely absent (e.g. Gregg et al. 2002). Since these objects have been definitively shown to exist, the debate turns to the question of their numbers in relation to the total quasar population.  Follow-up observations of the source populations identified from the very deep {\\it Chandra} X-ray fields (e.g. Alexander et al. 2003) have established that a large population of optically obscured AGN at bolometric luminosities $L_{Bol} \\la 10^{45}\\,{\\rm ergs}\\,{\\rm s}^{-1}$ exist at redshifts $z \\la 2$ (Barger et al. 2005). At higher luminosities, in the regime probed by the majority of large optically-selected quasar surveys (e.g. Hewett, Foltz \\& Chaffee 1995; Croom et al. 2004) the surface density of quasars is low and the X-ray catalogues contain relatively few objects, limiting the conclusions that can be drawn. In fact, Barger et al. (2005) find no evidence that a significant fraction of the quasar population is obscured at high luminosities, a finding apparently at variance with the conclusions of the widely cited study of Webster et al. (1995).  The {\\it Spitzer Space Telescope} has recently enabled deep, high-quality observations at infrared wavelengths to be performed. A type I quasar luminosity function (QLF) for the rest-frame 8$\\mu$m has been constructed by Brown et al. (2005), spanning $1<z<5$. The study specifically eliminated $z<1$ objects to avoid complications due to the presence of the quasar host galaxies. The shape of the derived infrared QLF is very similar to those constructed from optically-selected quasar samples. Comparison between surveys reveals that $\\sim$20 per cent of quasars present in the infrared sample are missing from optically-selected samples, but the missing fraction increases with decreasing luminosity. Most of the objects missing from the optical samples have red colours.  The conclusions drawn from a series of relatively modest samples of reddened quasars vary significantly. At one extreme, it has been suggested that the number of reddened objects far exceeds the number of optically-selected quasars (Webster et al. 1995), or that the obscured objects that have been found represent only `the most luminous tip of the red quasar iceberg' (White et al. 2003). However, other surveys (Francis, Nelson \\& Cutri 2004) claim that reddened quasars are not the dominant population. Targeted searches for type II quasars, the high-luminosity equivalents of the type II Seyfert population, predicted to exist in a number of unified-scheme models, have also begun to produce results (Zakamska et al. 2003, 2005) but the number of objects remains small.  The ability to undertake effective surveys for luminous quasars that are subject to a significant degree of reddening is about to be enhanced dramatically through the availability of large-format infrared cameras on 4m-class telescopes. In particular, the new Wide Field Camera (WFCAM) on the United Kingdom Infrared Telescope (UKIRT) on Mauna Kea has been commissioned successfully and the first observations for a series of photometric surveys under the banner of the UKIRT Infrared Deep Sky Survey (UKIDSS) have already begun. Of particular interest is the Large Area Survey (LAS) that will cover 4000 square degrees of the Sloan Digital Sky Survey (SDSS) to a depth more than three magnitudes fainter ($K \\simeq 18.5$) than the Two-Micron All Sky Survey (2MASS) (Cutri et al.  2003). Some 2000 square degrees of the LAS are scheduled for completion by mid-2007. Additional major wide-field near-infrared surveys will be undertaken by the WIRCam instrument on the Canada-France-Hawaii Telescope and the ESO VISTA telescope at Paranal.  For largely historical reasons almost all investigations of the form and evolution of the quasar luminosity function at redshifts $z<3$ from optically-selected surveys have been based on samples identified at blue wavelengths, employing flux limits in $B$, or, more recently, $b_J$ passbands. Quasar surveys in near-infrared passbands will almost certainly reveal the existence of a population of reddened objects that are, at best, under-represented in the traditional $B$-band samples. However, the predicted form of the number-magnitude and number-redshift relations for such surveys based purely on the known properties of quasars included in the $B$-band samples is necessary in order to interpret the new survey results. The combination of the very steep luminosity function at bright quasar luminosities, the form of the quasar SEDs (including strong emission lines) and the presence of host galaxies, which become increasingly prominent at longer wavelengths, means that the nature of the number-magnitude and number-redshift relations at near-infrared wavelengths are not as straightforward as they might at first seem. Certainly, the predictions are very different from those based on the approximation of a simple power-law SED for the quasars with a typical colour of $B-K \\simeq 2.5$ that is still commonly employed (e.g. Glikman et al. 2004).  The primary purpose of this paper is to take our current knowledge describing (i) the form and evolution of the quasar luminosity function from optical surveys, (ii) the shape of the ultraviolet through near-infrared SED of luminous unreddened quasars and (iii) the effect of host galaxy light on the magnitudes and colours of quasars, to produce predictions for the form of the number-magnitude and number-redshift relations expected from imminent wide-field near-infrared surveys. The predictions provide a reference against which the observational results may be compared, allowing the space density and properties of the new reddened population of quasars to be quantified.  The outline of the paper is as follows. Section 2 describes the input data required for the simulations, $\\S$3 discusses the estimation of quasar magnitudes from the SDSS data where the presence of a host galaxy leads to a non-stellar classification for the object in SDSS, and $\\S$4 briefly outlines some of the details of the simulation program. The results are presented in $\\S$5, with discussion and conclusions following in \\S6 and \\S7, respectively. A determination of the $b_J$ passband is described in Appendix A. Appendix B includes tabular material giving the predicted number-magnitude counts for a number of passbands. Concordance cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ is assumed throughout. Magnitudes on the Vega system are used throughout the paper, with the SDSS AB-magnitudes converted to the Vega system using the relations: $u=u_{\\rm AB}-0.93$, $g=g_{\\rm AB}+0.10$, $r=r_{\\rm AB}-0.15$, $i=i_{\\rm AB}-0.37$, and $z=z_{\\rm AB}-0.53$. ",
        "conclusions": "\\subsection{Initial Assessment}   Confidence in the results is gained by comparing the calculated number-magnitude and number-redshift relations with results from existing surveys. The excellent agreement between the predictions incorporating the distribution of host galaxy luminosities from the SDSS $i$-band sample, for 2QZ+6QZ in the $b_J$-band confirm that the presence of host galaxies has not biased significantly the determinations of the QLF from observations at blue wavelengths.  There is also good agreement between the predicted number counts in the $i$-band and the published SDSS results to $i\\sim 19.0$, which indicates that the fraction of reddened quasars missing from a $b_J$-selected sample but visible in the $i$-band, at nearly twice the observed-frame wavelength, is not large.  As seen in Figs. \\ref{fig:LF} through \\ref{fig:nzall}, employing an SED that represents the median quasar over a large range of wavelengths and adding a prescribed amount of host galaxy light mostly affect the low redshift ($z\\lesssim 1$) number counts, particularly at longer wavelengths. Hardly any effect is seen at $z\\gtrsim 1.5$ except for the UKIDSS LAS $K$-band simulation, where the faint magnitude limits enable intrinsically faint quasars to be detected to high redshifts. In fact, the UKIDSS LAS survey probes approximately one magnitude fainter down the QLF than does the 2QZ survey.  Morphological and colour restrictions applied to define quasar candidate samples need to be considered in order to avoid contamination from the many low-redshift galaxy-dominated objects while allowing the resolved quasars to be included. The results are relevant to the design of future surveys as depths and areas required to observe a given number of unobscured objects are now known. For example, the area of sky in the UKIDSS LAS needed to establish the existence of a given fraction of obscured quasars may be readily calculated.   \\subsection{Quasar SEDs}   The adoption of a single power-law approximation to the form of the restframe quasar continuum over $1200 < \\lambda < 5000\\,$\\AA \\ has proved largely adequate for surveys in the optical out to redshifts $z \\sim 2.5$. Extending observations into the near-infrared highlights the inadequacy of the single power-law approximation. Incorporation of the strong inflection in the SED at $\\lambda \\sim $12\\,000\\,\\AA\\ is essential in order to understand observations at near-infrared wavelengths. Fig. \\ref{fig:K2SED}(a) illustrates the dramatic difference in the K-correction at $z \\lesssim 1$. The effect on the optical--NIR colours of the quasars is large and the observed behaviour is sometimes misattributed to the presence of the host galaxy making the colours much redder, as in figure 7 of Glikman et al. (2004). Fig. \\ref{fig:SDSScolours} shows both the reddening of the optical--NIR colour from the SED at $z \\lesssim 1$ as well as the effect of the host galaxy, seen only at $z \\lesssim 0.3$. The shape of the SED and the presence of host galaxy clearly produce different signatures. Fig. \\ref{fig:K2SED}(b) shows the combined impact on the number of objects observed at $z\\lesssim 1$ at bright near-infrared magnitudes.  The simulation predictions are very conservative in the sense that an object contributing to the counts must have a minimum quasar absolute magnitude of $M_{b_J}=-21.0$, equivalent to $M_{K_{\\rm 2MASS}}=-24.6$. The objects identified as AGN in 2MASS by Cutri et al. (2002) and Francis et al. (2004) have no such limit applied. Furthermore, while a $J-K > 2.0$ is often taken as a definition of a `red' AGN, our default quasar SED has $J_{\\rm 2MASS}-K_{\\rm 2MASS}=2.5$ at $z=0.0$ and a non-negligible fraction of optically-selected quasars possess $J-K \\ge 2.0$ at low redshift. The impact of adopting fainter absolute magnitude limits for `quasar' samples, plus the effect of any additional colour selection criteria applied, need to be carefully assessed before reliable conclusions concerning any excess of objects over that predicted on the basis of optical samples can be reached.   \\subsection{Quasar--Galaxy Relationships}   Given our lack of knowledge concerning the distribution of host galaxy luminosities as a function of quasar luminosity the relations adopted in the simulations deliberately span an extended range in possible behaviour. One extreme, ($\\gamma$=1), is motivated by theoretical arguments from known relations between the mass of black holes and the mass of their host bulges. The distribution of $R_{gq}$ then approximates the range of accretion rates and efficiencies present in the population.  The other extreme, ($\\gamma$=0), where there is no relationship between quasar and host galaxy luminosity, approximates a situation in which any underlying relation between quasar and galaxy properties is overwhelmed by variations in accretion rates and efficiency (relative to the Eddington luminosity).  Intermediate between the two extremes, the $\\gamma$=0.42 case is observationally motivated (Croom et al. 2002). Such an observed relationship is perhaps not unexpected on theoretical grounds but the general applicability of the Croom et al. (2002) relationship remains to be established. The quasars included in their study cover a relatively modest range in redshift/luminosity and there must still be concerns over the exclusion of objects with substantial galaxy contributions. The fractional galaxy contribution was estimated from fibre spectra and Croom et al.'s claim that the fraction of host galaxy light that enters the fibre is effectively constant with redshift relies on assumptions concerning the host galaxy sizes and should be investigated quantitatively. However, we have chosen to adopt the $\\gamma$=0.42 case as the default for the simulations.  The key result of the simulations, in Section \\ref{qgrelation2}, is the significant impact of the presence of host galaxy light on the surveys in the near-infrared. The n(z) results for both of the $K$-band simulations are shown in Fig. \\ref{fig:K2KUlumgals}. The three prescriptions for adding host galaxy light affect different portions of the QLF at a given redshift. When $\\gamma$=1, the galaxies increase in brightness at the same rate as the quasars in this case, thus the n(z) relation is affected to higher redshifts. As can be seen in Fig. \\ref{fig:K2KUlumgals}(a) the n(z) at $z \\sim 2$ differ significantly among the host galaxy prescriptions, offering the prospect of constraining the relation between quasar and host galaxy luminosity at high redshifts.  The $\\gamma$=0 prescription is unique in that some relatively low-luminosity quasars have bright host galaxies. The faint quasars are then dominated by the host galaxy light, producing a dramatic steepening of the observed QLF at faint luminosities. The effect is most evident at low redshifts, $z < 1$, as the fractional contribution of host galaxy light decreases with increasing redshift, as the luminosity of the quasars increases faster than the passively evolving host galaxies. Whether such host galaxy-dominated objects feature prominently in a near-infrared survey depends on the morphological restriction imposed, since many very faint quasars are boosted above the faint magnitude limit of the simulation due almost entirely to light from the host galaxy. Requiring $R_{gq}<0.8$ ensures that host galaxy-dominated objects are excluded. Interpreting the results of surveys where host galaxy dominated objects are detected depends on the ability to estimate accurately the luminosity of the central source alone, so that relatively weak AGNs located in bright galaxies are not classified as quasars (c.f. Section \\ref{PSFMag}). As an illustration of the importance of accurately separating quasar from host galaxy light, the study of White et al. (2003) found several red objects in an $i$-band selected sample. However, two of the four objects with the largest derived $E(B-V)$ values from their table 1, (J1011+5205, J1021+5114), were classified as non-stellar by the SDSS photometry. As both objects are at $z\\sim1$, they must be dominated by their host galaxies in order to be resolved, and the red colours of the objects almost certainly result in part from host galaxy light.   \\subsection{Selection Criteria and Survey Design}\\label{select}   Morphological and colour selection applied to an object catalogue usually form an integral part of the identification of quasar candidates. A morphological restriction on object selection is approximated in the simulations by monitoring the fraction of host galaxy flux that is contributing to the total flux of each object, denoted $R_{gq}$. This restriction is only an approximation, as in practice the $R_{gq}$ value separating morphological classifications will be dependent on both the redshifts of the targets and the seeing conditions of the observations.  In the case of the 2QZ+6QZ survey, based on (by modern standards) relatively poor quality photographic imaging, the adoption of strict morphological selection turns out to have little impact on the fraction of the quasar population included. Changing the morphological restriction from $R_{gq}<0.8$ to $R_{gq}<0.3$ has almost no effect on the counts. The insensitivity to the presence of host galaxies arises due to a combination of the increasingly large contrast between quasar and galaxy flux at ultraviolet wavelengths and the relatively high lower redshift limit of $z_{low}=0.4$ imposed by Croom et al. (2002).  The same is not true for surveys in the near-infrared, as seen in Fig. \\ref{fig:K2KUselect}. Excluding extended sources will certainly eliminate many low redshift quasars. The effect can also be seen by using the 2MASS magnitudes for the objects in the SDSS DR3 quasar catalogue. If the restriction of requiring point-like morphology is imposed on the DR3 quasars, n(m) and n(z) counts in $K_{\\rm 2MASS}$ agree well with the quasar-only $K_{\\rm 2MASS}$ simulation, further reinforcing the importance of adopting a realistic quasar SED parametrization. Relaxing the point-like restriction, to include resolved objects, results in n(m) and n(z) counts in $K_{\\rm 2MASS}$ that agree well with the $K_{\\rm 2MASS}$ simulation that includes host galaxy light limited at $R_{gq}<0.8$. The study by Francis et al. (2004), searching for dust-reddened quasars, selected objects from the 2MASS catalogue and found a higher surface density of $1\\pm{0.3}$ quasar deg$^{-2}$ to $K_{\\rm 2MASS}=15.0$. They note that some of their `quasars' have optical colours indistinguishable from galaxies, and the objects are dominated by their host galaxies. They impose no morphological restriction on their sample, and some of their objects could easily have $R_{gq}$ values approaching unity and would thus be removed from our simulations.  It is clear from the discussion above that the inclusion of certain non-stellar objects in quasar candidate samples selected in the near-infrared will be necessary for many investigations. Of particular interest is the impact of the presence of host galaxy light on the colours of quasars and, thereby, on the effectiveness of colour-based selection schemes designed to isolate particular elements of the quasar population. For example, a key goal, now that faint $K$-band data from UKIDSS LAS and other surveys are becoming available, is to establish the fraction of reddened luminous quasars.  Warren, Hewett \\& Foltz (2000) introduced the KX-method of selection (by analogy with the UV-excess selection, or UVX) specifically to address the problem of isolating reddened quasars using a combination of optical and near-infrared colours. However, they did not consider the impact of host galaxy light on the proposed selection scheme. Fig. \\ref{fig:ccgal} shows the colour-colour diagram advocated by Warren et al. (2000). Shown are the default quasar and elliptical galaxy loci for redshifts $0.1<z<3.0$ and $0.1<z<2.0$, respectively, the locus of galactic dwarf stars, as well as quasar tracks with $0<R_{gq}<1$ at several redshifts. The crucial aspect of the KX-method is that reddened quasars move further away from the locus of galactic stars, which will be present in enormous numbers relative to the much rarer quasars. It can be seen immediately that the presence of the host galaxies moves the quasars redward in $r-J$. Of more potential concern is the horizontal displacement in $J-K$. At a fixed redshift, the locus of objects with $0<R_{gq}<1$ is indicated by the dotted line, joining the track (blue line) for the pure quasar, $R_{gq}=0$, to the track (red) for the pure galaxy, $R_{gq}=1$.  At low redshifts, $z \\lesssim 0.3$, increasing the amount of host galaxy light makes the objects more blue in $J-K$, so very low redshift, $z < 0.1$, galaxy-dominated objects will be excluded with a moderate colour cut, like the $J-K>1.2$ selection employed by Francis et al. (2004). Fortunately, the effect of the host galaxy light diminishes rapidly with increasing redshift and only for redshifts $z<0.1$ is proximity to the stellar locus an issue for the quasar selection. At $z\\gtrsim 0.3$, increasing host galaxy contribution causes the objects to appear redder in $J-K$, improving the effectiveness of the KX-selection scheme. Note however that the sense of the colour changes for quasars with $0.5 \\lesssim z \\lesssim 2$ are very similar to those due to simply reddening a pure quasar SED, highlighting the importance of quantifying the fraction of host galaxy light present when interpreting surveys for reddened quasars.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig15.ps}} \\caption{Colour-colour diagram showing the evolution of the quasar locus (top) for $0.1<z<3.0$ and galaxy locus (bottom) for $0.1<z<2.0$, as well as the location of the stellar locus for dwarf stars from the Bruzual-Persson-Gunn-Stryker Atlas included in the STSDAS {\\tt SYNPHOT} package. The dotted lines are for objects with $R_{gq}$ increasing from zero to unity, for redshifts 0.1, 0.2, 0.3, 0.4, 0.5, 0.6, and 0.7. Arrows indicate the direction of increasing redshift for the quasar and galaxy tracks, with the filled circles on the galaxy track marking redshifts 0.1, 1.0 and 2.0.} \\label{fig:ccgal} \\end{figure}"
    },
    "0512/astro-ph0512055_arXiv.txt": {
        "abstract": "We present a study of the dynamical evolution of Ultraluminous Infrared Galaxies (ULIRGs), merging galaxies of infrared luminosity $>10^{12}$ \\lsun. During our Very Large Telescope large program, we have obtained ISAAC near-infrared, high-resolution spectra of 54 ULIRGs (at several merger phases) and 12 local Palomar-Green QSOs, to investigate whether ULIRGs go through a QSO phase during their evolution. One possible evolutionary scenario is that after nuclear coalescence, the black hole radiates close to Eddington to produce QSO luminosities. The mean stellar velocity dispersion that we measure from our spectra is similar ($\\sim$160 km/s) for 30 post-coalescence ULIRGs and 7 IR-bright QSOs. The black holes in both populations have masses of order 10$^7$-10$^8$ \\msun\\ (calculated from the relation to the host dispersion) and accrete at rates $>0.5$ Eddington. Placing ULIRGs and IR-bright QSOs on the fundamental plane of early-type galaxies shows that they are located on a similar region (that of moderate-mass ellipticals), in contrast to giant ellipticals and radio-loud QSOs. While this preliminary comparison of the ULIRG and QSO host kinematical properties indicates that (some) ULIRGs may undergo a QSO phase in their evolutionary history before they settle down as ellipticals, further data on non-IR excess QSOs are necessary to test this scenario. ",
        "introduction": "\\label{intro} In the local Universe, the best laboratories for studying violent merging events (considered key-mechanisms in driving galaxy evolution) are the ultraluminous infrared galaxies (ULIRGs). Several studies indicate that ULIRGs transform gas-rich disks into moderate mass ellipticals through merger-induced dissipative collapse \\citep{kor92,mihos96}. Photometric analysis of ultraluminous merger remnants by \\citet{veilleux02} and \\citet{sco00} indicates that most of the sources are are well-fit by an elliptical-like r$^{1/4}$ light profile. Specrtoscopic analysis by \\citet{genzel01} and \\citet{tacconi02} indicates that their stellar kinematics resemble those of dispersion-supported systems.  While the end products of galactic mergers can be considered understood, several details of the merging process (often related to the physics of the gas) are still very uncertain, even in the local Universe. A plethora of numerical models \\citep{mihos96,springel05} and observations \\citep{sami96,kim02} indicates that major starburst episodes occur after the first galactic encounter and can be present after nuclear coalescence, before complete relaxation sets in. According to \\citet{dimatteo}, the main accretion event happens after the coalescence. However, the relative contributions of starburst and AGN to the infrared (IR) emission during each phase of a gas-rich merger are still unclear. Whether ULIRGs may go through a QSO phase after the nuclear coalescence also needs to be confirmed.  One way to investigate the physical details of mergers is to determine the evolution of the ULIRG kinematic properties as a function of time. Hence, we obtained high-resolution, H- and K- band ISAAC spectroscopic data for 54 ULIRGs (spanning a wide range of merger phase) during our European Southern Observatory Very Large Telescope (VLT) large program (ID 171.B-0442). Of these sources (at $0.018<z<0.268$), 23 are binary (not fully merged) ULIRGs and 30 are (relaxed) remnants. We extract the stellar dispersion $\\sigma$ and rotational $V_{\\rm rot}$ velocity from the CO rovibrational bandheads that appear in our spectra using the Fourier correlation quotient technique \\citep{bender90}; this method provides the line-of-sight (LOS) broadening function with which a stellar template has to be convolved to produce the observed spectrum. ",
        "conclusions": "From the high-resolution NIR spectrocopic data we have obtained for local ULIRGs and QSOs, we have found that \\begin{itemize} \\item For IR luminosity $>10^{12}$ \\lsun\\ to be triggered during a gas-rich merger, encounters of comparable mass galaxies are (typically) required. \\item Evolution (increase) of the host dispersion is observed as the merger advances from pre- to post- coalescence. \\item The merger remnants resemble moderate-mass ($\\sim$10$^{10}$-10$^{11}$ \\msun) ellipticals. The black holes they host are of the order 10$^{7}$-10$^{8}$ \\msun\\ and, on average, accrete at high Eddington rates ($\\ge$0.5). \\item The IR bright QSO dispersions and black hole masses, being of the order 10$^7$-10$^8$ \\msun, resemble those of ULIRG remnants and indicate a possible link between the two populations. Our IR-bright sources differ from QSOs that host supermassive black holes and accrete at low rates. \\item Imaging and spectroscopy of PG 1426+015 show that it is a binary system of nuclear separation $\\sim$4 kpc. Already at this early merger phase, one of the components is a powerful QSO (with strong dust continuum). \\end{itemize}"
    },
    "0512/astro-ph0512263_arXiv.txt": {
        "abstract": "A number of radio interferometers are currently being planned or constructed to observe 21 cm emission from reionization.  Not only will such measurements provide a detailed view of that epoch, but, since the 21 cm emission also traces the distribution of matter in the Universe, this signal can be used to constrain cosmological parameters at $6 \\la z \\la 20$.  The sensitivity of an interferometer to the cosmological information in the signal may depend on how precisely the angular dependence of the 21 cm 3-D power spectrum can be measured. Utilizing an analytic model for reionization, we quantify all the effects that break the spherical symmetry of the 3-D 21 cm power spectrum and produce physically motivated predictions for this power spectrum.  We find that upcoming observatories will be sensitive to the 21 cm signal over a wide range of scales, from larger than $100$ to as small as $1$ comoving Mpc.  Next, we consider three methods to measure cosmological parameters from the signal: (1) direct fitting of the density power spectrum to the signal (this method can only be applied when density fluctuations dominate the 21 cm fluctuations), (2) using only the velocity field fluctuations in the signal, (3) looking at the signal at large enough scales such that all fluctuations trace the density field.  With the foremost method, the first generation of 21 cm observations should moderately improve existing constraints on cosmological parameters for certain low-redshift reionization scenarios, and a two year observation with the second generation interferometer MWA5000 in combination with the CMB telescope Planck can improve constraints on $\\Omega_w$ (to $\\pm 0.017$, a $1.7$ times smaller uncertainty than from Planck alone), $\\Omega_m \\, h^2$ ($\\pm 0.0009$, $2.5$ times), $\\Omega_b \\,h^2$ ($\\pm 0.00012$, $1.5$ times), $\\Omega_{\\nu}$ ($\\pm0.003$, $3$ times), $n_s$ ($\\pm 0.0033$, $1.4$ times), and $\\alpha_s$ ($\\pm 0.003$, $2.7$ times).  Larger interferometers, such as SKA, have the potential to do even better.  If the Universe is substantially ionized by $z \\sim 12$ or if spin temperature fluctuations are important, we show that it will be difficult to place competitive constraints on cosmological parameters from the 21 cm signal with any of the considered methods. ",
        "introduction": "\\label{intro}  The reionization of the Universe involves many poorly understood astrophysical phenomena such as the formation of stars, the escape of ionizing photons from star-forming regions, and the evolving clumpiness of the gas in the intergalactic medium (IGM). However, reionization imprints signatures onto 21 cm emission from high-redshift neutral hydrogen, as will be studied with the instruments PAST, LOFAR, and MWA, in a manner that is sensitive to these processes.\\footnote{For more information, see http://www.lofar.org/, http://web.haystack.mit.edu/arrays/MWA/, and \\citet{pen04}.} Moreover, the 21 cm emission encodes information pertaining to fundamental cosmological parameters.  Due to all the overlying astrophysics, it is uncertain whether or not 21 cm observations can be competitive with other cosmological probes.  Several authors have discussed using the 21 cm signal from reionization to study cosmology, in addition to mapping out the epoch of reionization (EOR); \\citep{scott90, tozzi00, iliev02}. Recently, \\citet{barkana04a} show that redshift-space distortions from peculiar velocities allow for the decomposition of the observed 21 cm 3-D power spectrum into terms that are proportional to $\\mu^0, \\mu^2$ and $\\mu^4$, where $\\mu$ is the cosine of the angle between a mode $\\kvec$ and the line-of-sight (LOS). In principle, this decomposition makes it possible to separate the contribution from reionized bubbles from that owing to a fundamental cosmological quantity, the linear-theory density power spectrum.  Even if the signal from the ionized bubbles dominates over the cosmological one, \\citet{nusser04} shows that one can look for an asymmetry between the 21 cm signal measured in depth and that measured in angle.  The presence of this asymmetry may imply that the cosmology assumed in the analysis is incorrect [the Alcock-Paczynski (AP) effect].  This effect could help further constrain $\\Omega_m$ and $h$, as well as dark energy models \\citep{nusser04}.  It might be possible to distinguish the AP effect because it creates a $\\mu^6$ dependence in the 3-D power spectrum, which is distinct from the behavior that arises from velocity-field effects alone \\citep{nusser04, barkanaAP}.  For both the $\\mu$ decomposition of the 21 cm power spectrum and the AP effect, the feasibility of inferring cosmological parameters using future surveys depends on how sensitive these surveys are to deviations from spherical symmetry in the 3-D power spectrum. \\citet{morales05} suggests that 21 cm observations should spherically average $\\bfk$-modes over a shell in Fourier space to increase the signal-to-noise ratio.  In addition to losing the angular information contained in the signal, such averaging would significantly bias upcoming measurements: The power spectrum is not close to spherically symmetric and 21 cm interferometers will be most sensitive to modes oriented along the LOS.  Array design also factors into the sensitivity to certain features in the signal.  The first generation of EOR arrays are still being planned, and so it is important to understand different design trade-offs.  Authors have considered the 21 cm emission signal in several limits, such as when the spin temperature fluctuations are still important and the HII regions are small \\citep{barkana05, ahn05} or when the spin temperature is much larger than the cosmic microwave background (CMB) temperature \\citep{zald04, furlanetto04a, santos05}.  In reality, we do not know how quickly X-rays from the first stars, black holes and shocks will heat the gas and how quickly the spin temperature can couple to the gas temperature via the Wouthuysen-Field effect \\citep{wout, field58}.  Previous work suggests that these processes act quickly after the first stars turn on \\citep{oh01, venkatesan01, chen04,ciardi03-21cm}.  Moreover, it is expected that quasars or stellar sources contribute the bulk of the ionizing radiation.  If this is the case, reionization will be a patchy process, in which the HII regions grow around these sources as the ionization fraction increases.  Upcoming observations will be most sensitive to lower redshifts ($z \\sim 6-12$) during reionization \\citep{bowman05}.  At these low redshifts, it is likely that the spin temperature is greater than the CMB temperature and that the ionized fraction is of order unity and very patchy.  This is the regime that we consider for much of this paper. It is also possible that the ionization fraction is near zero for a period at these low redshifts, which will facilitate cosmological parameter estimation.  We consider this case as well.  The organization of this paper is as follows.  In \\S \\ref{21cm} we make physically motivated predictions for the form of the 3-D 21 cm power spectrum.  We then generalize the detector noise calculation of \\citet{morales05} to a power spectrum that is not spherically symmetric (\\S \\ref{noise}) and incorporate realistic foregrounds into this sensitivity calculation (\\S \\ref{foregrounds}).  These calculations allow us to estimate the sensitivity of upcoming interferometers to the 21 cm power spectrum (\\S \\ref{detectors}).  We conclude with a discussion of how the 21 cm signal can be used to measure fundamental cosmological parameters as well as a Fisher matrix analysis to estimate how precisely future observations can constrain these parameters (\\S \\ref{cosmology}).  In our calculations, we assume a cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.046$, $H=100 \\, h \\, {\\rm km \\; s^{-1} \\; \\Mpc^{-1}}$ (with $h=0.7$), $n=1$, and $\\sigma_8=0.9$, consistent with the most recent determinations \\citep{spergel03}.  All distances are measured in comoving coordinates. ",
        "conclusions": "\\label{Discussion}  In this paper, we have used the FZH04 analytical model of reionization to calculate the power spectrum of 21 cm brightness temperature fluctuations $P_{\\Delta T}(\\bfk)$, extending this calculation beyond calculations done in FZH04 by including redshift-space distortions. When $\\bar{x}_i \\la 0.5$ or on smaller scales than the effective bubble size, these distortions are quite important and give the power spectrum a substantial anisotropy between modes parallel and perpendicular to the LOS. These distortions not only increase the signal, but may allow us to separate $P_{xx}$, $P_{x\\delta}$ and $P_{\\delta \\delta}$, facilitating the measurement of the size and bias of the HII bubbles and perhaps the spectrum of density fluctuations in the Universe.  We show that higher order terms may complicate the separation of $P_{\\Delta T}$ when the ionized fraction is significant.  To quantify the detectability of $P_{\\Delta T}(\\bfk)$, and in particular $P_{\\delta \\delta}$, we make realistic sensitivity estimates for LOFAR, MWA and SKA.  The most important parameter for these interferometers is the collecting area.  But, this is not to say that the other factors that go into the design are unimportant.  We agree with the conclusion of \\citet{bowman05} that, everything else being equal, arrays with denser cores will be more sensitive to $P_{\\Delta T}(\\bfk)$.  This is because modes along the LOS can be detected by even the shortest baselines, and arrays with cores have more of these shorter baselines.  The antenna size can also have a similar effect: arrays that have large antennae cannot pack them as closely together as arrays with small antennae.  As a result, they will not sample the shorter baselines as well.  Smaller antennae also provide a larger FOV, which will aid statistical detections of the signal.  Because the current design for LOFAR does not include the shorter baselines that the design for MWA has and because the design for LOFAR results in a much smaller FOV, we find that, despite differences in collecting area, LOFAR and MWA will be comparably sensitive to $P_{\\Delta T}(\\bfk)$ at most redshifts. This is not to say that this will be the case when these instruments are actually deployed.  Since none of the discussed 21 cm arrays have begun construction, their designs can still be optimized.  Even with an optimally constructed radio interferometer, the removal of foregrounds that are $10^4$ times larger than the 21 cm fluctuations will be a serious challenge.  In this paper, we find that foregrounds will contaminate the signal on scales greater than the depth of the slice used to construct the 21 cm power spectrum.  At most scales smaller than this, we are optimistic that foregrounds can be cleaned below the signal.  On such scales, we show that fitting a quadratic or cubic polynomial to the observed visibilities in the frequency direction has little difficulty cleaning a realistic model for foregrounds.  It does not appear to be the case, as was claimed by \\citet{oh03} and \\citet{gnedin04}, that foregrounds will contaminate all angular modes beyond repair.  Applying our calculation for the detector noise and foreground power spectrum, we find that MWA and LOFAR will not be sensitive to the $P_{\\mu^4}$ component of $P_{\\Delta T}(\\bfk)$.  This component is particularly interesting because it traces the linear-theory density power spectrum.  However, these interferometers will be sensitive to $P_{\\mu^0}$, which probably will tell us more about the astrophysics of reionization than about cosmology, except perhaps for very small ionization fractions.  MWA5000 and SKA will be moderately sensitive to $P_{\\mu^4}$ and $P_{\\mu^6}$, but not sensitive enough to provide competitive constraints on cosmological parameters. We find that only if there exists a time when density fluctuations dominate $P_{\\Delta T}$, will upcoming probes of 21 cm emission be able to place competitive constraints on cosmological parameters.  In addition, planned 21 cm interferometers will not be very sensitive to the signal for $z \\ga 12$.  The is primarily because detector noise fluctuations are proportional to $T_{\\rm sky}(z)$, which scales as $(1 + z)^{2.6}$.  If there is a period where density fluctuations dominate $P_{\\Delta T}$, a $2$ yr observation with MWA5000 plus Planck can give the constraints $\\delta \\Omega_{w} = 0.0017$ (a $1.7$ times smaller uncertainty than from Planck alone), $\\delta w = 0.06$ ($1.5$ times), $\\delta \\Omega_m h^2 = 0.0009$ ($2.5$ times), $\\delta \\Omega_b h^2 = 0.00012$ ($1.5$ times), $\\delta n_s = 0.0033$ ($1.4$ times), $\\delta \\alpha_s = 0.003$ ($2.7$ times), $\\delta \\Omega_{\\nu} = 0.003$ ($3$ times) and $\\delta x_H = 0.03$. SKA plus Planck yield similar constraints and MWA50K can do even better. However, if $\\tau = 0.17$, as suggested by WMAP, and reionization began at $z \\approx 20$, observations of the signal at scales much larger than the effective bubble size may be the most promising direct method to probe cosmology (\\S \\ref{largescales}).  Observations must overcome many additional challenges beyond those that have been discussed in this paper.  Issues that we have not addressed include contamination by radio recombination lines, terrestrial radio interference, the residuals left from wave front corrections for a turbulent atmosphere, and the enormous data analysis pipeline needed to analyze potentially larger data sets than those from current experiments.  The 21 cm signal will also be affected by gravitational lensing by intervening material \\citep{zahn05, mandel05}.  If taken into account, this effect can further improve cosmological parameter estimates \\citep{zahn05}.  Cosmic variance sets a limit on how well we can constrain cosmological models with the CMB.  Because 21 cm emission can be observed as a function of redshift, this signal allows us to measure many more independent modes than is possible with the CMB.  We have seen that cosmological parameters are extractable from 21 cm emission.  In an ideal case in which reionization begins at relatively low redshifts, upcoming interferometers may be able to compete with future CMB experiments such as Planck.  If reionization begins at higher redshifts or if the spin temperature fluctuations are important, a more sensitive interferometer will be required than those that are currently planned to be able to compete with CMB parameter constraints.  Regardless of how reionization actually proceeded, high-redshift 21 cm emission has the potential to become a valuable probe of cosmology.\\\\ \\\\   We thank Judd Bowman, Bryan Gaensler, Adam Lidz and Miguel Morales for useful discussions.  This work was supported in part by NSF grants ACI 96-19019, AST 00-71019, AST 02-06299, and AST 03-07690, and NASA ATP grants NAG5-12140, NAG5-13292, and NAG5-13381.  \\begin{appendix}"
    },
    "0512/physics0512208_arXiv.txt": {
        "abstract": "General stability criterions of two-dimensional inviscid parallel flow are obtained analytically for the first time. First, a criterion for stability is found as $\\frac{U''}{U-U_s}>-\\mu_1$ everywhere in the flow, where $U_s$ is the velocity at inflection point, $\\mu_1$ is eigenvalue of Poincar\\'{e}'s problem. Second, we also prove a principle that the flow is stable, if and only if all the disturbances with $c_r=U_s$ are neutrally stable.  Finally, following this principle, a criterion for instability is found as $\\frac{U''}{U-U_s}<-\\mu_1$ everywhere in the flow. These results extend the former theorems obtained by Rayleigh, Tollmien and Fj\\o rtoft and will lead future works to investigate the mechanism of hydrodynamic instability. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512145_arXiv.txt": {
        "abstract": "{We present the first optical photometry of the counterpart to the candidate intermediate polar RX J0153.3+7446. This reveals an optical pulse period of $2333\\pm5$~s. Reanalysis of the previously published {\\em ROSAT} X-ray data reveals that the true X-ray pulse period is probably $1974\\pm30$~s, rather than the 1414~s previously reported. Given that the previously noted orbital period of the system is 3.94~h, we are able to identify the X-ray pulse period with the white dwarf spin period and the optical pulse period with the rotation period of the white dwarf in the binary reference frame, as commonly seen in other intermediate polars. We thus confirm that RX J0153.3+7446 is indeed a typical intermediate polar. ",
        "introduction": "Intermediate polars are cataclysmic variables, characterised by pulsed X-ray emission which reflects the rotation period of an accreting magnetic white dwarf. Over two dozen confirmed systems are currently known (for a recent list see Norton, Wynn \\& Somerscales 2004) and of these, around a quarter were discovered by {\\em ROSAT} during its all sky survey. In an important paper, Haberl \\& Motch (1995) reported the first {\\em ROSAT} observations of six intermediate polar candidates. Five of these systems have gone on to be well studied. RE0751+14 (PQ Gem; see Duck et al 1994, Potter et al 1997), RX J0028.8+5917 (V709 Cas; see Norton et al 1999, de Martino et al 2001) and RX J0558.0+5353 (V405 Aur; see Allan et al 1996, Evans \\& Hellier 2004) are now confirmed as typical intermediate polars with a wealth of published observations; RX J1712.6--2414 (V2400 Oph;  see Buckley et al 1997; Hellier \\& Beardmore 2002) is the first confirmed stream-fed intermediate polar; and RX J1914.4+2456 (V407 Vul; see Cropper et al 1998, Ramsay et al 2002, 2005, Norton, Haswell \\& Wynn 2004) has excited much controversy as a possible ultra-compact binary, although its nature is still in doubt.  However, the sixth object reported by Haberl \\& Motch (1995), namely RX J0153.3+7446, has been virtually ignored ever since. They showed an X-ray pulse profile of the system, claiming the pulse period as 1414s, and mentioned that a publication on optical observations of the object was in preparation. Such a publication has never appeared, although the on-line cataclysmic variable catalogue (Downes et al 2001) lists an optical counterpart in Cassiopeia with V=16.4 at RA 01:53:20.9, Dec +74:46:22 and mentions an orbital period of 0.16415~d credited to a private communication from John Thorstensen. The only published observation of RX J0153.3+7446 is an identification spectrum by Liu \\& Hu (2000) which confirms its cataclysmic variable nature by virtue of the typical emission line spectrum.  Optical photometry of intermediate polars usually shows modulation at the X-ray pulse period (i.e. the spin period of the white dwarf in most cases) and/or the beat period (i.e. the spin period of the white dwarf in the binary reference frame). The latter modulation, where seen, is presumed to be due to reprocessed X-ray emission originating from the illuminated surface of the donor star or some other structure fixed in the binary reference frame. An orbital photometric modulation is often seen as well, presumably depending on the inclination angle of the system. ",
        "conclusions": "We have detected a strong optical photometric modulation from RX J0153+7446 at a period of $2333\\pm5$s. Having reanalysed the {\\em ROSAT} X-ray observation of the source, we find that the true X-ray pulse period is probably $1974\\pm30$sec. Assuming that the X-ray period represents the white dwarf spin period and the optical period represents the rotation period of the white in the binary reference frame, then the two are in agreement with the reported orbital period for this system of 3.94h. RX J0153+7446 is thus confirmed as an intermediate polar.  The spin to orbital period ratio of RX J0153.3+7446 is therefore 0.14, which is one of the highest amongst intermediate polars above the period gap. From the results presented in Norton, Wynn \\& Somerscales (2004) we can estimate the magnetic moment of the white dwarf in this systems as $\\sim 6 \\times 10^{33}$~G~cm$^{3}$. This places RX J0153.3+7446 at the high end of the intermediate polar magnetic moment distribution and suggests it may be a candidate for exhibiting polarized emission and also episodes of stream-fed accretion. Further X-ray observations are encouraged to confirm and refine the X-ray pulse period and orbital period of this system."
    },
    "0512/astro-ph0512003_arXiv.txt": {
        "abstract": "We report the \\chandra\\ detection of \\ovii\\ K$\\alpha$ absorption at $z=0$ in the direction of the $z=0.03$ Seyfert 1 galaxy Mkn 279.  The high velocity cloud Complex C lies along this line of sight, with \\hi\\ 21-cm emission and \\ovi\\ 1032\\AA\\ absorption both observed at velocities of $\\sim -150$\\kms\\ relative to the local standard of rest.  We present an improved method for placing limits on the Doppler parameter and column density of a medium when only one unresolved line can be measured; this method is applied to the \\ovii\\ absorption seen here, indicating that the \\ovii\\ Doppler parameter is inconsistent with that of any low--velocity (Galactic thick disk) or high--velocity \\ovi\\ (\\ovihv) component. Direct association of the \\ovii\\ with the \\ovihv\\ is further ruled out by the high temperatures required to produce the observed \\ovihv/\\ovii\\ ratio and the significant velocity difference between the \\ovii\\ and \\ovihv\\ lines.  If the \\ovii\\ absorption is associated with a very broad, undetected \\ovi\\ component, then the absorption must be broadened by primarily nonthermal processes.  The large velocity dispersion and possible slight redshift of the \\ovii\\ absorption (as well as limits on the absorber's temperature and density) may be indicative of a local intergalactic medium origin, though absorption from a hot, low--density Galactic corona cannot be ruled out. ",
        "introduction": "The advent of high--resolution X-ray spectroscopy with the \\chandra\\ X-ray Observatory and XMM--Newton, as well as far--ultraviolet spectroscopy with the Far Ultraviolet Spectroscopic Explorer (FUSE) and the Space Telescope Imaging Spectrograph on HST, has shed a great deal of light on the local warm--hot intergalactic medium (WHIM).  Hydrodynamic simulations predict that this tenuous web of $\\sim 10^5-10^7$\\,K gas should contain about half of the baryons in the nearby universe \\citep{cen99,dave01}, appearing as a forest of highly ionized metal absorption lines in high signal--to--noise X-ray spectra of background sources \\citep{hellsten98,perna98}. Indeed, recent \\chandra\\ grating observations of this ``X-ray forest'' toward bright blazars have confirmed these predictions \\citep[][Nicastro et al., in preparation]{nicastro05a,nicastro05b}.  In addition to the intervening absorption systems, similar metal absorption lines (primarily \\ovii) are observed at redshifts consistent with zero toward several background quasars, such as Mrk~421 \\citep[][hereafter W05]{williams05}, PKS~2155-304 \\citep[][Williams et al., in preparation]{nicastro02,fang02}, and 3C 273 \\citep{fang03}; larger archival analyses have also been performed for various instruments \\citep[e.g.][]{mckernan04,fang06}.  This nearby absorption presents a  unique puzzle: since no morphological or distance information is known about the X-ray absorption (other than limited spatial data along the very few quasar pencil beams where it has been detected), it is still unknown whether this warm--hot gas originates within the Galactic halo or rather is part of the Local Group WHIM, or a combination of the two. While some \\ovii\\ absorption has been detected within 50\\,kpc of the Galaxy \\citep{wang05}, this is unlikely to be uniformly distributed. Additionally, constrained simulations of the Local Supercluster predict high column densities of \\ovi\\ and \\ovii\\ in some directions \\citep{kravtsov02}.  It is thus likely that Galactic and extragalactic phenomena both contribute to the local X-ray absorption, but the question of which sightlines are dominated by which phenomena is, by and large, unanswered.  Further complicating the issue is the presence of other gaseous components of unknown origin.  \\ion{H}{1} high--velocity clouds (HVCs) have a velocity distribution inconsistent with the Galactic rotation and therefore are thought to be either neutral gas high in the Galactic halo (perhaps from tidally stripped dwarf galaxies) or cooling, infalling gas from the surrounding IGM.  Along many lines of sight studied with FUSE, high--velocity \\ovi\\ absorption lines (\\ovihv) at velocities coincident with the \\ion{H}{1} HVCs are seen, while in some other directions \\ovihv\\ is present even in the absence of any significant \\ion{H}{1} 21-cm emission at the same velocity \\citep{sembachetal03}.  The question of whether these isolated \\ovi\\ HVCs arise in an extended, hot Galactic corona or at Local Group scales -- and their relation to the $z=0$ X-ray absorption lines -- is a subject of continuing debate.  There is some evidence for such a Galactic corona \\citep[discussed in detail by][and references therein]{sembach03}; indeed, a significant fraction of HVCs appear to exhibit low--ionization absorption lines (such as C\\,{\\sc ii}--{\\sc iv} and Si\\,{\\sc ii}--{\\sc iv}), unlikely to arise in a low--density, warm--hot photoionized medium \\citep{collins05}. However, \\citet{nicastro03} showed that the velocity distribution of these unassociated \\ovihv\\ clouds is minimized in the Local Group rest frame, indicating that their origin is extragalactic. Additionally, W05 found that the \\ovihv\\ toward Mrk 421 cannot be associated with the \\ovii\\ along that sightline (assuming a single temperature/density phase). However, the links between the \\ovihv, \\ion{H}{1} HVCs, and local \\ovii/\\oviii\\ absorption in the context of the Galactic corona and local WHIM are to a large degree unknown.  Determining the origin of these components, and the relations between them, is a crucial step in our understanding of the ongoing process of galaxy formation  The nearby ($z=0.03$), X-ray bright Seyfert galaxy Mrk~279 lies in the direction of the \\ion{H}{1} HVC Complex C, thus providing a particularly valuable background source that can be used to study these gaseous components.  Here we report on our analysis of deep \\chandra\\ and FUSE spectra of this object, the detected UV and X-ray absorption, and the implications for gas in the Galaxy and Local Group. ",
        "conclusions": "Long--duration \\chandra\\ grating observations of the bright AGN Mrk~279 reveal the presence of strong \\ovii\\ K$\\alpha$ absorption at a redshift consistent with zero.  A FUSE spectrum of the same source shows several additional \\ovi\\ components at velocities near zero.  Through kinematic, curve--of--growth, and ionization balance modeling, we conclude the following: \\begin{enumerate} \\item{A direct $\\chi^2$ analysis of the \\chandra\\ spectrum coupled with absorption line models constrains the Doppler parameter of the \\ovii\\ absorption to be $b>74$\\kms\\ and $b<24$\\kms. This latter range is unlikely due to the extremely high \\ovii\\ column densities required to produce the strong absorption feature.} \\item{The \\ovii\\ Doppler parameter limits are inconsistent with the measured $b$ values for any of the $v\\sim 0$ \\ovi\\ absorption components.  Additionally, the centroid of the \\ovii\\ K$\\alpha$ line is inconsistent (at the $2.5\\sigma$ level) with that of the \\ovihv, indicating that the \\ovii\\ is not associated with \\emph{any} local \\ovi\\ component.} \\item{If the \\ovii\\ absorption is associated with a broad, undetected \\ovi\\ absorption line, then a large Doppler parameter ($b\\sim 200$\\kms) is required to provide a single--phase solution for the \\ovi, \\ovii, and \\oviii\\ column densities.  This large value of $b$ could be a result of either microturbulence, velocity shear from infalling gas, or broadening due to the Hubble expansion over a path length of a few Mpc.  If the line is purely Hubble--broadened, at $b=200$\\kms\\ a pathlength of 3\\,Mpc and density of $\\log n\\sim -5$ is implied (assuming an oxygen abundance of 0.3 times solar).} \\item{The large velocity dispersion, possible redshift, and lack of association with any Galactic absorption components (as well as the proximity of HVC Complex C) indicates that this X-ray absorption may be from a large--scale nearby WHIM filament; however, a Galactic corona origin cannot be ruled out with the current data.} \\end{enumerate}"
    },
    "0512/astro-ph0512529_arXiv.txt": {
        "abstract": "We present the results of a systematic search in 8.5 years of {\\it Rossi X-ray Timing Explorer} All-Sky Monitor data for evidence of periodicities.  The search was conducted by application of the Lomb-Scargle periodogram to the light curves of each of 458 actually or potentially detected sources in each of four energy bands (1.5--3 keV, 3-5 keV, 5-12 keV, and 1.5--12~keV).  A whitening technique was applied to the periodograms before evaluation of the statistical significance of the powers.  We discuss individual detections with focus on relatively new findings. ",
        "introduction": "Periodic modulation of the X-ray intensity has been observed in many X-ray binary systems (see a general discussion in White, Nagase, \\& Parmar 1995).  The periods of the detected signals range from a few milliseconds~\\citep{wvk98} to hundreds of days~\\citep{hugo00}, and are, in general, associated with spin of a neutron star, binary orbital motion, or precession of a tilted accretion disk or other phenomenon in which the period usually exceeds the orbital period.  In addition, very long quasi-periods have been found to be associated with the outbursts of X-ray transients.  The periodicities attributable to anisotropic emission from rotating neutron stars are found mostly in high-mass X-ray binaries (HMXBs) and in a few low-mass X-ray binaries (LMXBs).  The pulse periods range from a few milliseconds to at least tens of minutes, and possibly hours (e.g., Hall et al. 2000).  The known orbital periods of LMXBs range from 0.19 h to 398 h, with the majority between 1 h and 2 d \\citep{white95}. The observed properties of the orbital modulation depend on the viewing angle.  At low orbital inclination angles ($< 70^\\circ$), X-ray orbital periods are rarely observed.  At intermediate inclinations, periodic dipping behavior, and in a few cases, brief eclipses of the X-rays by the companion may be observed.  The dips may be caused by obscuration of materials splashed above the disk plane when the gas stream from the companion hits the accretion disk.  For high inclination systems ($>80^\\circ$), X-ray eclipses have been observed in several systems but the number is smaller than what might be expected if the systems simply consist of a dwarf companion overflowing its Roche lobe and transferring material to a compact object via a thin accretion disk~\\citep{joss79}.  To resolve this discrepancy, it has been proposed~\\citep{milgrom78} that the central X-ray source is hidden behind a thick accretion disk rim while the X-rays are scattered via a photo-ionized corona above the disk and can still be seen. The source thus appears extended and may be partially eclipsed periodically by the bulge at the disk rim formed by the impact of the accretion flow on the disk.  The known orbital periods of HMXBs range from a few hours to hundreds of days, with the majority above 1 day.  All but one of the identified supergiant systems have nearly circular orbits with periods less than 15 d, whereas most Be star systems have eccentric orbits and periods of several tens or hundreds of days \\citep{white95}.  The orbital periods are evident by means of several effects including eclipse of the X-rays by the companion, phase-dependent absorption and scattering in a stellar wind from the companion, and absorption dips caused by non-axisymmetric structure of an accretion disk, accretion streams, or other structures.  In Be star systems, the orbital periods may be manifest as periodic X-ray outbursts that occur when the accretion rate onto the compact object is enhanced around periastron passage.  Superorbital periods have been observed in both HMXBs and LMXBs.  In a handful of cases, i.e., the on-off cycles of Her X-1 (35 d period), LMC X-4 (30 d), and SMC X-1 ($\\sim55$ d), and the 164 d jet precession cycle of SS 433, the modulation is very likely due to the presence of a precessing tilted accretion disk that periodically occults the X-ray source or guides the directions of the emitted jets \\citep[e.g.,][]{katz73,levine82,pringle96}. The disk precession can be sufficiently regular to produce quasi-coherent or relatively narrow features in a Fourier spectrum made from a few years of data.  In other sources, the causes of the modulation corresponding to reported superorbital periods are not well-established.  Indeed, the modulations may be weak and/or long compared with the duration of the observations, so the stability and significance of the effects are difficult to quantify.  Very long-term quasi-periodicities have been seen in several recurrent X-ray transients. The X-ray light curves in these transients are characterized by prominent outbursts separated by long periods of quiescence.  The physical mechanism underlying these outbursts is still unclear.  The favored scenario at present is some variation of the disk instability model \\citep[see the review by][]{lasota01}. Transient behavior as well as quasi-periodicities of outburst occurrence times are known to occur in both neutron star systems and systems comprising black hole candidates.  Among them, the Rapid Burster (X1730$-$333), a neutron star system, is known to exhibit outbursts that sometimes recur every 200 d or so (but at other times the recurrence interval has been as short as $\\sim100$ d), and the black-hole candidate 4U 1630$-$47 is known to exhibit outbursts at 500 to 700 day intervals.  A summary of detections of transient X-ray sources observed with the \\RXTE/ASM prior to the year 2000 can be found in \\cite{bradt00}.  Most of the known periodicities have been found in studies of one or a few sources, but there have also been a few more systematic searches for periodic signals.  \\citet{priedhorsky83} analyzed data obtained between 1969 and 1976 by the Vela 5B satellite on 4 sources in the Centaurus region and confirmed the 41.5 day period of GX301$-$2, strengthened the evidence for a 187 d period in 4U1145$-$619, and reported a period of 132.5 days in GX304$-$1.  Another search using Vela~5B data focused on 9 galactic X-ray sources in the Aquila-Serpens-Scutum region and uncovered evidence for a 199 day period in X1916$-$053, and a 41.6 day period in X1907+09 as well as a 122--125 day cycle in the outbursts of the recurrent transient Aql X-1 \\citep{priedhorsky84a}.  \\citet{smale92} searched for periodicities with periods longer than 1 day in 17 confirmed or suspected low-mass X-ray binaries also using data from Vela 5B.  They confirmed the $\\sim175$ day period in X$1820-303$ \\citep{priedhorsky84b} and found evidence for a $\\sim 77$ day period in Cyg X-2, but reported that they did not convincingly detect the 199 day period in X1916$-$053 nor were long-term periods detected in 13 of the other systems. They found evidence of a 333 d period in Cyg~X-3 which could also be explained by non-periodic variability \\citep[see also][]{pt86}.  \\citet{smale92} concluded that long-term cyclic variability is rare in LMXBs. Another systematic periodicity search was done by \\citet{priedhorsky95} on 8 bright X-ray binaries detected by WATCH/Eureca, but they did not find any previously unreported periodicities.  Herein we report on a search for periodicities in data obtained with the All-Sky Monitor (ASM) on the \\rxte (\\RXTE).  Since commencing regular observations in 1996 March, the ASM has monitored over 400 X-ray sources.  Among these, about 150 X-ray sources have been detected with an average intensity above 10 mCrab for at least one week.  We have conducted a global search for evidence of periodic behavior in 458 sources using results from the ASM produced over its first 8.5 years of operation.  This paper is organized as follows.  Section~\\ref{data} describes the data and observational constraints.  Section~\\ref{tech} describes technical details of the Lomb-Scargle (L-S) periodogram. Section~\\ref{strategy} describes the implemented search strategies.  Section~\\ref{results} gives a summary of results and brief discussions of selected detections. ",
        "conclusions": "We have presented results from a systematic periodicity search through the first 8.5 years of source intensity measurements made using the \\RXTE/ASM.  The ASM data base is an incredibly valuable resource that has allowed us to carry out this search in a uniform fashion using the intensity records for a large number of X-ray sources.  Our search has served the purposes of discovering and measuring new periodicities and providing a means for improving the level of knowledge of previously known periodic phenomena in X-ray sources.  The detection strategy we have adopted is conservative in general.  It, hopefully, can be useful in clarifying the authenticity of some previously claimed periodicities.  The results of our search demonstrate that orbital modulation is more readily detected in HMXBs than in LMXBs \\citep{joss79}. A total of 33 orbital periods have been securely detected in our search.  Only eight of them are believed to be in LMXBs, while twenty-four are in HMXBs with sixteen in supergiant systems and eight in Be-star systems.  Out of the $\\sim$ 400 X-ray sources monitored with the ASM, over 100 of them are believed to be LMXBs and only $<60$ are believed to be HMXBs. Yet it is apparent from the present results, that there are far fewer orbital periodicities detected in the LMXBs than in the HMXBs.  In particular, the fraction of eclipses observed in LMXBs is $<3\\%$, well below that of the HMXB systems, consistent with previous observations that the number of eclipses detected in LMXBs is well below the expected rate if the companion stars are dwarf stars that fill their Roche lobe and if the accretion disks are thin."
    },
    "0512/astro-ph0512235_arXiv.txt": {
        "abstract": "We describe high resolution Smoothed Particle Hydrodynamics (SPH) simulations of three approximately $M_*$ field galaxies starting from \\LCDM initial conditions. The simulations are made intentionally simple, and include photoionization, cooling of the intergalactic medium, and star formation but not feedback from AGN or supernovae. All of the galaxies undergo an initial burst of star formation at $z \\approx 5$, accompanied by the formation of a bubble of heated gas. Two out of three galaxies show early-type properties at present whereas only one of them experienced a major merger. Heating from shocks and -PdV work dominates over cooling so that for most of the gas the temperature is an increasing function of time. By $z \\approx 1$ a significant fraction of the final stellar mass is in place and the spectral energy distribution resembles those of observed massive red galaxies. The galaxies have grown from $z=1 \\rightarrow 0$ on average by 25\\% in mass and in size by gas poor (dry) stellar mergers. By the present day, the simulated galaxies are old ($\\approx 10 \\;{\\rm Gyrs}$), kinematically hot stellar systems surrounded by hot gaseous haloes. Stars dominate the mass of the galaxies up to $\\approx 4$ effective radii ($\\approx 10$ kpc). Kinematic and most photometric properties are in good agreement with those of observed elliptical galaxies. The galaxy with a major merger develops a counter-rotating core.  Our simulations show that realistic intermediate mass giant elliptical galaxies with plausible formation histories can be formed from \\LCDM initial conditions even without requiring recent major mergers or feedback from supernovae or AGN. ",
        "introduction": "There are many puzzles encountered in understanding the formation and evolution of elliptical galaxies and the spheroidal components of spiral galaxies. On the one hand a naive reading of the hierarchical theory of structure formation in a $\\Lambda$CDM universe would argue that, since massive halos form later than less massive ones, massive ellipticals which reside in the centers of these massive halos should also form late.  But there is strong observational evidence that old, massive, red and metal rich proto-ellipticals are already in place at $z=2-3$ and that present day early-type galaxies formed most of their stars well before a redshift $z=1$ \\citep{1973ApJ...179..427S,2000ApJ...536L..77B,2005ApJ...633..174T, 2005ApJ...631..145V}. We also know from evidence dating to the 1970's that current rates of star formation in these systems are quite low, the rates increasing sharply into the past as $z^1$ \\citep{1978ApJ...219...18B,1980ApJ...236..351D,1995ApJ...439...47R}. Further, with regard to the evolution of this population, there appears to be a significant increase in the total stellar mass in these old objects from $z=1$ to the present \\citep{2004ApJ...608..752B,2004ApJ...608..742D, 2005ApJ...620..564C,2005astro.ph..6044F}, but this cannot easily be accounted for by the fading of younger, bluer star forming galaxies present at $ z =1$ \\citep{2004ApJ...608..742D, 2005astro.ph..6044F}.  This has led to the plausible and popular idea that `dry merging' (i.e. merging of predominantly stellar systems) among elliptical systems \\citep{2005ApJ...627L..25T,2005AJ....130.2647V,2006astro.ph..2038B, 2006ApJ...640..241B} pushes more and more of the mass over the observational cutoff to provide an increased number of 0.5-2$L_*$ galaxies.  This attractive idea gives reasonable explanations for some observed correlations of the mass of ellipticals and their kinematics, sizes or isophotal shape \\citep{2005MNRAS.359.1379K,2006ApJ...636L..81N,2006MNRAS.369.1081B}. However, noting that there are tight relations among luminosity, color, age and metalicity, puts severe limits on the amount of late dry merging that could acceptably occur without destroying the aforementioned constraints.  On the theoretical side it is now possible to perform accurate high resolution simulations of the gravitational evolution of the dark matter distribution \\citep{1998ApJ...499L...5M,2005Natur.435..629S}. In contrast, the numerical simulation of galaxy formation, including a hydrodynamic treatment of the baryonic component, is still in its infancy. Very few high resolution simulations from realistic cosmological initial conditions have been done so far and most of these have concentrated on the formation of disc galaxies rather than early-type spheroidal systems \\citep{2003ApJ...596...47S,2003ApJ...597...21A,2004ApJ...607..688G,2004ApJ...606...32R}.  This is surprising since spheroidal systems are of interest in their own right, as they contain more than half of the total stellar mass in the local universe \\citep{1998ApJ...503..518F}. The most massive galaxies known,  the giant ellipticals, are spheroidal systems which  predominantly consist of old stars (see e.g. \\citealp{2005ApJ...621..673T}) and so must have formed at high redshift. They are therefore likely to be good probes of galaxy assembly, star formation and metal enrichment in the early Universe. \\begin{figure} \\centering \\includegraphics[width=8cm]{./f1.eps} \\caption{Circular velocity curves for galaxy A at four different numerical resolutions: $40^3$, $50^3$, $100^3$, and $200^3$ SPH particles and collisionless dark matter particles, respectively. Note how the rotation curves become increasingly flat as the resolution increases. \\label{rotcurve_haloA}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=8cm]{./f2.eps} \\caption{Star formation rate (SFR) histories, computed from stellar ages, of galaxy A versus lookback time at four different numerical resolutions: $40^3$, $50^3$, $100^3$, and $200^3$ SPH particles and collisionless dark matter particles, respectively. There is a strong trend that the low redshift star formation rate is reduced in higher resolution simulations. \\label{sfr}} \\end{figure}  Most numerical work on early-type galaxy formation has used either very idealized initial conditions \\citep{2004MNRAS.347..740K}, or had insufficient spatial and mass resolution to resolve the internal structures of galaxies \\citep{2003MNRAS.346..135K,2004ApJ...601L.131S}. An exception is the simulation discussed by \\citet{2003ApJ...590..619M}. These authors used an SPH simulation that included feedback from supernovae to follow the formation of a single spheroidal galaxy self-consistently from CDM initial conditions. The spatial resolution of this simulation was high enough to resolve the region within an effective radius for a typical real  elliptical galaxy. However, the final stellar system formed in this simulation was far too dense, with an effective radius about an order of magnitude smaller than real elliptical galaxies of the same brightness. \\citet{2003ApJ...590..619M} speculate that this discrepancy may be a consequence of their star formation and feedback algorithm and that it might be possible to produce less concentrated systems if more aggressive stellar feedback were implemented to prevent star formation in high density sub-units at high redshift.  In this paper we present high resolution hydrodynamical simulations based on cosmological initial conditions admitting, intentionally, only bare-bones prescriptions for the physics involved (e.g. no ``feedback'' from supernovae or AGN), to see if some resolution of these paradoxes can be derived.  The rationale for simplyfing the simulations is straightfoward. As we will show in Section \\ref{MASSIVE}, at least $100^3$ gas particles are required for hydrodynamical simulations to `converge', making them very expensive to run. With present generation computers it is impossible to run a large ensemble of simulations of this size to properly explore a huge parameter space. Our point of view, therefore, is to keep the physics of the simulations as simple as possible and to get an understanding of the behaviour of a simplified problem before investigating additional complexities such as supernova and AGN feedback. Our goalin this paper, therefore, is to investigate the formation of a number of massive ($\\approx M_*$) isolated galaxies starting from realistic initial conditions and to see how variable the final systems and  whether they resemble real galaxies.  What we find can be summarized simply.  The initial cooling and collection of cold gas from infalling smaller scale perturbations happens very rapidly and easily within the most massive systems leading to a very rapid burst of star formation beginning at $z \\sim 6$ and then falling off exponentially on a time-scale of roughly 1.5Gyrs. This phase is terminated as the star forming region is enveloped in an expanding hot bubble which prevents new, infalling cold gas mass elements from reaching the central regions. This early phase - reminiscent in some respects to a modernized version of the ``monolithic collapse'' picture of galaxy formation, prominent in the 1960's through 1980's \\citep{1974MNRAS.169..229L,1982MNRAS.201..939V}- produces a sequence of objects with effective radii of 1-2kpc, which might satisfy the tight relations observed among the red metal rich old cores of elliptical galaxies.  For simulations with early-type properties at $z=0$ stellar accretion or mergers (the choice of the appropriate term for the process is arbitrary and dependent on the relative mass of the infalling stellar objects) add to the growing stellar envelope of relatively blue, old and metal poor stars. This accounts for the growth in size seen for ellipticals as well as the growth in total mass in the time frame $z=1 \\rightarrow 0$.  Furthermore, since this assembly of added stellar mass is not accompanied by much in situ star formation (the coterminously added hot and cold gas being simply added to the expanding hot bubble surrounding the central galaxy) the growth of the stellar population occurs without the presence of young stars.  Existing work has shown that the tightness of the elliptical color-magnitude relation puts strong constraints on dry merger scenarios \\citep{1998MNRAS.299.1193B,2005MNRAS.360...60K}. But we find that in our simulations that minor mergers or accretion events do not typically add much stellar mass to the central region, because of the relatively large angular momentum of the infalling systems - and thus the central, tight color-magnitude relations are not expected to be overly perturbed even if as much as 40\\% of the total final mass is added during this phase.  As an additional byproduct of this assembly scenario, we find it readily understandable how so much of the stellar mass in these systems resides in regions of such low density, where star formation would always have been difficult to contemplate. The stars seen at $R \\sim R_e$ were not born there, but rather were formed in the central regions of much lower mass systems that have been accreted and shredded in the envelopes of the giant elliptical galaxies.  \\begin{table*} \\caption{Properties of galaxy A at $r < r_{\\mathbf{vir}}$ } \\label{A_global} \\centering \\begin{tabular}{c| c c c c c c | c c c c c c}     % \\hline\\hline Resolution & $M_{\\mathbf{vir}}$ \\tablenotemark{(a)} & $M_{\\mathbf{stars}}$ &$M_{\\mathbf{gas}}$ & $M_{\\mathbf{dark}}$ &  $r_{\\mathbf{vir}}$ \\tablenotemark{(b)}&  $v_{\\mathbf{max}}$  \\tablenotemark{(c)}&  $m_{\\mathbf{stars}}$ \\tablenotemark{(d)} & $m_{\\mathbf{dark}}$& $\\epsilon_{\\mathbf{stars}}$ \\tablenotemark{(e)}& $\\epsilon_{\\mathbf{dark}}$  \\\\ \\hline $40^3$  & 242 &  25.3 & 20.2 & 197 & 434 & 341 &  161 & 1288 & 0.625& 1.3\\\\ $50^3$  & 241 &  27.0 & 20.1 & 193 & 433 & 356 &  82.5 & 659  & 0.5  & 1.0\\\\ $100^3$ & 230 &  23.8 & 20.9 & 184 & 427 & 270 &  10.5 &  82   &0.25 & 0.5\\\\ $200^3$ & 225 &  24.9 & 18.4 & 182 & 424 & 232 &  1.3 & 10.3  &0.125& 0.25 \\\\ \\hline \\end{tabular} \\tablecomments{(a) Total masses $M$ in$10^{10}M_{\\odot}$; (b) Virial radius in kpc; (c) Maximum circular velocity in km/s; (d) Particle masses $m$ in $10^{5}M_{\\odot}$; (e) Gravitational softening lengths in kpc}  \\end{table*}  \\begin{table} \\caption{Properties of galaxy A at $r < 30$kpc} \\label{A_30} \\centering \\begin{tabular}{c| c c c}     % \\hline\\hline Resolution & $M_{\\mathbf{stars}}$ \\tablenotemark{(a)}& $M_{\\mathbf{gas}}$ &$M_{\\mathbf{dark}}$  \\\\ \\hline $40^3$  & 16.3 &  1.0 & 27.7 \\\\ $50^3$  & 16.6 &  1.0 & 28.7 \\\\ $100^3$ & 12.2 &  0.6 & 20.4 \\\\ $200^3$ & 11.8 &  0.5 & 23.1 \\\\ \\hline \\end{tabular} \\tablecomments{(a) Total masses $M$ in$10^{10}M_{\\odot}$} \\end{table} ",
        "conclusions": "We have followed the formation and evolution of three $\\approx M_*$ field galaxies using numerical simulations 'ab initio' from cosmological initial conditions at high redshift. All galaxies start to form when cold gas that has collapsed in subunits is effectively consumed by star formation. At early times the remaining diffuse gas is predominantly heated by shocks. In general the stars galaxies assemble both by in situ star formation as well as major/minor mergers and accretion. During the early formation phase at $2<z<8$ the assembly of the galaxies is dominated by mergers of gas rich subcomponents and  in situ star formation. Although some fraction of the stars is accreted this phase has  the characteristics of a dissipative collapse. Thereafter stellar accretion or minor/major mergers become more important and , for early-type galaxies, tends to dominate for $z < 1$.  After z=1 the further assembly of the galaxies varies from halo to halo. The evolution of one galaxy is dominated by gas infall and in situ star formation at rates of $5-8 M_{\\odot}$/yr. It develops a massive, however too compact, stellar disk and has  properties more similar to early-type disk galaxies.  Two other galaxies have present day properties more similar to elliptical galaxies and their assembly is dominated by minor and major predominantly stellar mergers. One of the galaxies grows by minor mergers and does not experience  a major merger after its formation phase. At redshift $z \\approx 1$ it has formed a massive stellar spheroid (see e.g. \\citealp{2005ApJ...618...23N}) which resembles observed massive evolved galaxies and their inferred formation histories \\citep{2000A&A...362L..45D}. The system evolves into an old stellar system surrounded by a halo of hot gas at the present day. The second early-type galaxy had a major merger at z=0.6 that can be classified as a gas poor or 'dry' merger as it was not accompanied by a strong  burst of star-formation. From those two examples we find that $\\approx$ 15\\% - 35\\% of the final stellar mass was assembled by accretion of stars that have formed outside the galaxy (see \\citealp{2006MNRAS.370..902K}). This increase will contribute to the evolution of the luminosity function of elliptical galaxies from $z=1\\rightarrow 0$ in addition to a fading population of blue galaxies and mergers \\citep{2004ApJ...608..752B,2004ApJ...608..742D,2005ApJ...620..564C, 2005astro.ph..6044F}. The size evolution of the simulated early-type galaxies is determined by the accreted stars. While the half-mass radius of the stars formed in-situ remains almost constant independent of redshift at $r_{1/2}= 1-2$ kpc the accreted stars create an envelope whose half-mass radius is increasing with decreasing redshift (see \\citealp{2006ApJ...648L..21K}, for a semi-analytical approach). Interestingly, a significant fraction of the stars at the final half-mass radius of the galaxies were not born there but accreted. This process might provide a good explanation for a possible size evolution of early-type galaxies \\citep{2006ApJ...650...18T,2006astro.ph.10241L} without changing the fundamental parameter relations which are determined by the central properties of the galaxies and should have been established during the early formation phase.   At the present day most investigated kinematical and photometric properties (see also \\citealp{2003ApJ...596...47S}) are consistent with observed rotating disky and boxy early-type galaxies at similar luminosities. In particular, the concentration is in good agreement with observations. We believe that the discrepancy with \\citet{2003ApJ...590..619M} is due to the different star formation algorithm which in their case allows predominantly gaseous subunits to merge and develop a high concentration before forming stars. In our simulations gas is converted into stars efficiently in subunits which merge at early times, prior to z=2, and so the morphology of the systems is elliptical if its late accretion is dominated by stars. Interestingly, the early-type galaxy with the late major merger develops a counter rotating core which previously only have been successfully simulated in isolated equal-mass disk-disk mergers \\citep{1991Natur.354..210H,2006astro.ph..6144J} and are supposed to be signposts of merger events.   Feedback from supernovae and/or AGN, which was not included in this simulation, would provide additional heat sources but appear neither vital to stop the star formation and produce a massive red galaxy at higher redshifts nor to produce a present day stellar system with consistent kinematical properties. Theoretical arguments (see {\\it e.g.} \\citealp{1986ApJ...303...39D}; \\citealp{2000MNRAS.317..697E}) suggest that stellar feedback rapidly becomes ineffective in systems with velocity dispersions greater than 100 km/s. However, AGN feedback may be important in galaxies with high velocity dispersions \\citep{1998A&A...331L...1S,1999MNRAS.308L..39F,2001ApJ...551..131C}. In our case AGN feedback might have been able to prevent further star formation at low redshift by heating the cold gas at the center of the galaxy (see e.g. \\citealp{2005ApJ...620L..79S,2005MNRAS.358..168S}). On the other hand many field ellipticals show the related E+A phenomenon which seems to be suppressed in the centers of clusters as any residual cold gas will be blown away \\citep{1993AJ....106..473C}. For more massive systems than the ones studied here it has been argued that feedback from AGN might have a bigger impact and be responsible for the sharp cutoff at the bright end of the luminosity function of ellipticals \\citep{2005MNRAS.356.1155C}. However, only full cosmological simulations at high resolution of the formation of more massive systems will help to fully understand the importance of stellar and AGN feedback on the formation and evolution of elliptical galaxies.  In a hierarchical universe most galaxies are expected to experience mergers but it has been argued for a long time that dissipation is essential to explain properties of elliptical galaxies and in particular intermediate mass ellipticals (\\citealp{1989ApJ...342L..63K} and references therein; \\citealp{1996ApJ...464L.119K,2006MNRAS.372..839N}). The cosmological formation processes presented here combines characteristics of both stellar mergers and dissipational collapse/mergers whereas the latter dominates the early formation phase. It naturally meets many observational constraints on the evolution of massive spheroidal galaxies (e.g. \\citealp{2005ApJ...633..174T,2005ApJ...621..673T, 2005ApJ...631..145V}). It might be equally or even more important than major mergers of disk galaxies \\citep{2006ApJ...641...21R,2006ApJS..163....1H} which can not reproduce the properties the most massive ellipticals \\citep{2003ApJ...597..893N,2006astro.ph..7446C} whose evolution could be driven by mergers of evolved ellipticals \\citep{1977ApJ...217L.125O,2003ApJ...597L.117K,2006ApJ...640..241B,2006ApJ...636L..81N,2006MNRAS.369.1081B}."
    },
    "0512/astro-ph0512079_arXiv.txt": {
        "abstract": "The origin of the micro-Gauss magnetic fields in galaxy clusters is one of the outstanding problem of modern cosmology. We have performed three-dimensional particle-in-cell simulations of the nonrelativistic Weibel instability in an electron-proton plasma, in conditions typical of cosmological shocks.  These simulations indicate that cluster fields could have been produced by shocks propagating through the intergalactic medium during the formation of large-scale structure or by shocks within the cluster. The strengths of the shock-generated fields range from tens of nano-Gauss in the intercluster medium to a few micro-Gauss inside galaxy clusters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512553_arXiv.txt": {
        "abstract": "In this comment we question some arguments presented in astro-ph/0512048 to refuse the presence of an singular mass surface layer. In particular, incorrect expressions are used for the disk's surface mass density. We also point out that the procedure of removing the descontinuity on the $z=0$ plane with a region of continuous density gradient generates other two regions of descontinuities with singular mass surface layers making the  model unrealistic. ",
        "introduction": "In \\cite{Cooperstock1} the authors make some comments about their previous work \\cite{Cooperstock2} and also reply to certain issues that have been raised by some authors \\cite{Korzynski,Vogt,Garfinkle} concerning the nature of the matter distribution and the asymptotic behaviour of the spacetime in their model. In particular, \\cite{Korzynski} pointed out that the use of a $|z|$ function in the solution of the field equations introduces an additional disk on the $z=0$ plane and later \\cite{Vogt} showed that the disk was made of exotic matter.  Consequently, \\cite{Cooperstock1} examine the question in several ways and present some arguments to dismiss the existence of a physical mass layer present on the $z=0$ plane. We would like to notice some inconsistencies in their paper. The first is related to the expressions used for the surface density of the singular disk, and is presented in Sec.\\ \\ref{sec_disk}. In Sec.\\ \\ref{sec_smooth} we question the smoothing procedure discussed in  Section 4  of \\cite{Cooperstock1}.  We find that one cannot eliminate disk singularities, moreover one introduces new singularities that make the model unrealistic. ",
        "conclusions": ""
    },
    "0512/astro-ph0512286_arXiv.txt": {
        "abstract": "Polarized diffuse emission observations at 1.4-GHz in a high Galactic latitude area of the northern Celestial hemisphere are presented. The $3.2\\degr\\times3.2\\degr$ field, centred at RA~=~$10^{h}$$58^{m}$, Dec~=~$+42^{\\circ}$$18'$ (B1950), has Galactic coordinates $l$~$\\sim$~$172^\\circ$, $b$~$\\sim$~$+63^\\circ$ and is located in the region selected as northern target of the BaR-SPOrt experiment. Observations have been performed with the Effelsberg 100-m telescope. We find that the angular power spectra of the $E$-- and $B$--modes have slopes of $\\beta_E = -1.79\\pm0.13$ and $\\beta_B = -1.74\\pm0.12$, respectively. Because of the very high Galactic latitude and the smooth emission, a weak Faraday rotation action is expected, which allows both a fair extrapolation to Cosmic Microwave Background Polarization (CMBP) frequencies and an estimate of the contamination by Galactic synchrotron emission. We extrapolate the $E$--mode spectrum up to 32-GHz and confirm the possibility to safely detect the CMBP $E$--mode signal in the Ka band found in another low emission region \\citep{carretti05b}. Extrapolated up to 90-GHz, the Galactic synchrotron $B$--mode looks to compete with the cosmic signal only for models with a tensor-to-scalar perturbation power ratio $T/S < 0.001$, which is even lower than the $T/S$ value of 0.01 found to be accessible in the only other high Galactic latitude area investigated to date. This suggests that values as low as $T/S = 0.01$ might be accessed at high Galactic latitudes. Such low emission values can allow a significant {\\it red}-shift of the best frequency to detect the CMBP $B$--mode, also reducing the contamination by Galactic dust, and opening interesting perspectives to investigate Inflation models. ",
        "introduction": "The polarization of the Cosmic Microwave Background (CMB) is a powerful probe to investigate the early Universe. Inflation,  scalar (density) and tensorial (gravitational waves) primordial perturbations, and formation processes of first stars and galaxies can be effectively studied through its angular power spectra (e.g. \\citealt*{zaldarriaga97,kinney99,cen03}). The CMB Polarization (CMBP) can be expanded in a curl-free and a curl component called $E$-- and $B$--mode, respectively (\\citealt{zase97}; for an equivalent formalism see \\citealt*{KKS}). These fully describe the polarized emission and have the advantage to be scalar quantities, differently from the Stokes parameters $Q$ and $U$ that are components of a 2-spin tensor.  The emission level of the $E$--mode is a few per~cent of the CMB temperature anisotropy signal. First detections have been recently performed \\citep{kovac02,readhead04,barkats04,leitch05,montroy05}, but a full characterization is far from being completed. On subdegree scales, level and behaviour of the $E$-mode are well predicted by the standard model, provided a number of cosmological parameters are determined (e.g. those of the {\\it concordance} model after the WMAP data set: \\citealt{spergel03}). A precise measurement of the power spectrum thus represents a powerful check of the overall theoretical framework based on both the standard model and the Inflation paradigm.  Foreground noise sources can compete with the cosmic signal and disturb the precise measurements needed for CMBP goals. Estimates of contaminant emissions are then crucial to define the best observing conditions, as the identification of the frequency window in which the cosmic signal is accessible as a function of the sky position. These estimates also allow the determination of the foreground properties, which, in turn, greatly improve the effectiveness of foreground cleaning algorithms (e.g. \\citealt{tegmark00,tucci05,verde05}). The diffuse synchrotron emission of the Galaxy is considered one of the most relevant polarization contaminant and is expected to dominate the foreground contribution at frequencies less than 100-GHz. Recent measurements performed in a low emission area at both 1.4 and 2.3-GHz indicate that the synchrotron signal does not represent an obstacle for the $E$--mode at high Galactic latitude and at frequencies higher than 30-GHz \\citep{bernardi03,carretti05a,carretti05b}. However, this is just one sample that may not be representative of the situation throughout high Galactic latitudes and other measurements in different areas are required to confirm this first evaluation.  The second CMBP component, the $B$--mode, is even fainter and no information which can help predict its level are available so far. In fact, within the Inflation paradigm its level on degree scales is related to the amount of primordial tensor perturbations (Gravitational Waves, GW) generated in the early whiles of the Universe by the Inflation event (e.g. see \\citealt*{ka98,kinney99,boyle05}). Usually, this GW amount is measured by the tensor-to-scalar perturbation power ratio $T/S$ and only upper limits are available so far ($T/S < 0.71$, 95~per~cent confidence level: \\citealt{spergel03}). The detection of the CMB $B$--mode and the measure of its value are thus fundamental to disentangle among different Inflation models and to measure their relevant parameters (e.g. \\citealt{kinney99}). Estimates of the $B$--mode emission of the foreground contaminants (mainly synchrotron and dust emissions of the ISM) are crucial not only to identify clean sky areas, but also to understand the range of Inflation models which can be explored through CMB measurements. The low emission area surveyed by \\citet{carretti05b} was found promising also for the detection of this CMBP component. In fact, these authors find that models with a GW amount down to $T/S=0.01$ should be accessible in that area. However, as for the $E$--mode, this is just one sample of the high Galactic latitude conditions and measurements in other sky areas are needed to check whether this trend is general.  High Galactic latitude fields are promising for CMBP investigations for their low total intensity emission both by ISM synchrotron and dust. Among them, \\citet{carretti02} identify a couple of wide fields (about $30^\\circ\\times30^\\circ$) as observing targets for the BaR-SPOrt experiment \\citep{cortiglioni03}, one for each of the northern and southern Celestial hemispheres.  An area in the southern field has been observed by \\citet{bernardi03} and \\citet{carretti05b}, providing the results briefly reported above.  In this paper we present deep 1.4-GHz observations of an area in the northern field. Centred at RA~=~$10^{h}$$58^{m}$, Dec~=~$+42^{\\circ}$$18'$ (B1950), the observed area features very high Galactic latitudes ($l$~$\\sim$~$172^\\circ$, $b$~$\\sim$~$+63^\\circ$), which should prevent the significant alterations of the polarized emission  by Faraday rotation possible at this frequency for $|b| < 40^\\circ$--$50^\\circ$\\citep{carretti05a}. This leads to estimates of the $E$-- and $B$--mode power spectra which can be safely extrapolated to the frequency range relevant for CMB studies.  The paper is organized as follows: in Section~\\ref{obsSec} we present details of the observations along with the discussion of the obtained maps.  In Section~\\ref{specSec} we provide the analysis of polarized angular power spectra. Finally, in Section~\\ref{discSec} we discuss the implications for CMBP measurements. ",
        "conclusions": "\\label{discSec}  In Section~\\ref{specSec} we find indications that the Effelsberg area should not be affected by significant Faraday rotation effects. Therefore, it is likely that the power spectra we find are not significantly modified and can be safely scaled up to the CMBP frequency range. It is worth noting that, in any case, Faraday effects would enhance the power on these angular scales.  Therefore, our results would be an upper limit of the real synchrotron emission, providing us with a worst case analysis of the contamination of the CMB by this foreground.  Extrapolations have been performed assuming the brightness temperature frequency slope $\\alpha = -3.1$ typical of the synchrotron emission in the 1.4--23~GHz range \\citep{bernardi04}. Considering that \\citet{bennett03} find a continuous steepening of this slope from K (23-GHz) to W band (94-GHz), our assumption should provide worst case estimates when applied to the 30--90~GHz range. Brightness temperatures are converted into CMB thermodynamic temperatures by the factor \\begin{equation} c = \\left[{2\\,\\sinh(x/2) \\over x}\\right]^2, \\hskip 1cm x = h\\nu / kT_{\\rm cmb} \\end{equation} where $\\nu$ is the frequency and $T_{\\rm cmb} = 2.725$~K \\citep{mather99}.  Figure~\\ref{specCMB_E_Fig} reports the $E$--mode spectrum extrapolated to 32-GHz, which is the lower frequency of the BaR-SPOrt experiment and at the low end of the frequency range relevant for CMBP aims [both the DASI \\citep{leitch05} and CBI experiment \\citep{readhead04} observed in such a band]. The synchrotron contribution is well below the cosmic component throughout the peak region ($\\ell = 150$--1500). Even the low out-of-horizon feature at $\\ell$~$\\sim$~$150$ looks well accessible.  A comparison with the southern region is done using the 2.3-GHz data taken in that area, more reliable for extrapolation than the 1.4-GHz ones. The contamination in the two areas look similar, apart from lower values of the northern one on smallest scales, so that our results substantially confirm the conclusions of \\citet{carretti05b}. This result is also supported by the analyses of both the DASI and CBI experiment teams, who find indications that, in other two regions, the emission of this foreground is not dominant at 30-GHz. This scenario suggests the detectability of CMBP $E$--mode is achievable at high Galactic latitudes even in the Ka-band.  \\begin{figure} \\centering \\includegraphics[angle=0, width=1.0\\hsize]{fig5.ps} \\caption{$E$--mode angular power spectrum $C^E_{\\rm Eff}$ of the 1.4-GHz data extrapolated to 32-GHz. A frequency spectral slope of $\\alpha = -3.1$ is assumed. The extrapolation of the 2.3-GHz spectrum obtained in the southern area is also reported ($C^E_{\\rm Pks}$). The power spectrum expected for the CMB is shown for comparison (solid). The cosmological parameters used are those of the concordance model as determined in \\citet{spergel03}. \\label{specCMB_E_Fig}} \\end{figure} \\begin{figure} \\centering \\includegraphics[angle=0, width=1.0\\hsize]{fig6.ps} \\caption{$B$--mode angular power spectrum $C^B_{\\rm Eff}$ of the 1.4-GHz data extrapolated to 90-GHz. A frequency spectral slope of $\\alpha = -3.1$ is assumed. The extrapolation of the 2.3-GHz spectrum obtained in the southern area is also reported ($C^B_{\\rm Pks}$). Spectra expected for the CMBP are shown for comparison (solid) by using tensor-to-scalar power ratios of $T/S = 0.01$ and $T/S = 0.001$, respectively. The other cosmological parameters used are those of the concordance model as determined in \\citet{spergel03}. \\label{specCMB_B_Fig}} \\end{figure}  Figure~\\ref{specCMB_B_Fig} reports the $B$--mode spectrum extrapolated to 90-GHz, that is in the frequency range expected to be the best trade-off between Galactic synchrotron and dust contamination. Even though the peak near $\\ell$~$\\sim$~$90$ (angular scale $\\theta$~$\\sim$~$2\\degr$) is marginally accessible by our data, the result is very promising. The comparison with the expected CMBP signal for different $T/S$ values suggests that the properties of Inflation models with a GW amount as low as $T/S = 0.001$ can be explored here, when only the synchrotron contribution is considered. This result is even better than the one obtained in the southern region, where $T/S$ values down to $\\sim$~0.01 were estimated to be accessible.  As already pointed out by \\citet{carretti05b}, this low level of the polarized synchrotron emission can make the ISM dust the major contaminant of CMBP at 90--100~GHz. No estimate of the polarized dust emission is available in our region, and a real comparison is not possible. However, we can attempt a comparison using the existing data. The most proper one for our area would be the upper limit obtained by \\citet{ponthieu05} with the data of the ARCHEOPS experiment, our area being included in the sky region they surveyed. However, their result looks too pessimistic to be applied to our low emission area. In fact, besides being an upper limit, it refers to the whole $|b| > 10\\degr$ portion of the region mapped by that experiment. Moreover, the value they provided (0.2-$\\mu$K$^2$ at 100-GHz) is much higher than the upper limit deduced by \\citet{carretti05b} relative to a large high Galactic latitude area. Instead, we prefer to use the dust total intensity measurements recently obtained at high Galactic latitudes in the area surveyed with the 2003 BOOMERanG experiment \\citep{masi05}. They also find that the dust emission in that area is representative of a large fraction of the high latitude sky (40 per cent of the sky). As polarization fraction we consider both 5 and 20~per~cent, which brackets the 10~per~cent deduced by \\citet{benoit04} for the high Galactic latitudes.  The frequency behaviours are plotted in Figure~\\ref{foregFig}, where the quantity $\\sqrt{\\ell(\\ell+1)\\,C^B_\\ell / (2\\pi)}$ at $\\ell = 90$ is used as good indicator of the emission on the $2\\degr$ scale of the $B$--mode peak. The dust actually looks to be the leading contaminant at 90--100~GHz, even assuming the case of 5~per~cent polarization. As a consequence, the best frequency window where the total foreground contamination is minimum can be {\\it red}-shifted toward $\\sim70$~GHz, as it happens for the CMB anisotropy \\citep{bennett03}. To give an idea of the situation at 70-GHz, Figure~\\ref{specCMB_B_70_Fig} shows the synchrotron angular behaviour at such a frequency.  Besides implications for the design of experiments devoted to the CMBP $B$-mode,  such a red-shift implies a reduction of the dust contamination, because of its frequency behaviour. As a result, even considering the dust contribution, it should be possible to study models with $T/S = 0.01$, and even lower in case of 5~per~cent dust polarization. As mentioned above, this result is even better than that obtained in the other high Galactic latitude area, also reported in Figure~\\ref{foregFig} for comparison. Although observations in other sky regions are needed, this suggests that the southern area is not a special case, and that the possibility to reach the $T/S = 0.01$ value estimated to be achievable in the southern area could be extended to large low emission regions at high Galactic latitudes. Foreground cleaning techniques can allow further improvements of the measurable $T/S$ (e.g. \\citealt{tegmark00,tucci05,verde05}) opening wide possibilities in the explorable Inflation models, especially in view of the recent results of \\citet{boyle05}. In fact, these authors find that the interesting class of Inflation models with minimal fine-tuning have $T/S > 0.01$. Only models with an high degree of fine-tuning can have $T/S$ values less than 0.001. Thus, perspectives to detect the $B$-mode in large low emission areas at high Galactic latitudes become realistic. However, a better assessment of the scenario requires direct measures of the polarized dust emission, possibly at higher frequency (hundreds of GHz) where this Galactic component is stronger."
    },
    "0512/physics0512102_arXiv.txt": {
        "abstract": "\\noindent  According to modern developments in turbulence theory,  the \"dissipation\" scales (u.v. cut-offs) $\\eta$ form a random field related to  velocity increments  $\\delta_{\\eta}u$. In this work  we, using Mellin's transform combined with the Gaussain  large -scale boundary  condition, calculate  probability densities  (PDFs) of velocity increments $P(\\delta_{r}u,r)$ and the PDF  of the  dissipation scales $Q(\\eta, Re)$, where $Re$ is the large-scale Reynolds number.    The resulting expressions  strongly  deviate from the Log-normal PDF  $P_{L}(\\delta_{r}u,r)$ often quoted in the literature.  It is shown that  the   probability density of the small-scale velocity fluctuations  includes  information about the large (integral) scale dynamics  which is responsible for deviation of $P(\\delta_{r}u,r)$ from $P_{L}(\\delta_{r}u,r)$.   A  framework for evaluation of the PDFs of various turbulence characteristics involving spatial derivatives is developed.  The exact relation, free of spurious Logarithms  recently discussed in Frisch et al  (J. Fluid Mech. {\\bf 542}, 97 (2005)),   for the multifractal probability density of velocity increments, not based on the steepest descent evaluation of the integrals is obtained and the calculated function $D(h)$ is close to experimental data.  A novel   derivation  (Polyakov,  2005),   of a well-known result of the multi-fractal theory [Frisch,  \"Turbulence. {\\it  Legacy of A.N.Kolmogorov}\", Cambridge University Press, 1995)) ,   based on  the concepts  described  in this paper,  is also presented. ",
        "introduction": "\\noindent    A reasonably  well  experimentally established anomalous (multi) scaling of  the structure functions  $S_{n,m}=\\overline{(u(x+r)-u(x))^{n}(v(x+r)-v(x))^{n}}\\equiv \\overline{ (\\delta_{r}u)^{m}(\\delta_{r}v)^{m}}$ , where $u$ and $v$ are components of the velocity field parallel and perpendicular to the displacement vector ${\\bf r}=r{\\bf i}$, respectively,   is one of  the properties  of  strong turbulence which makes it \"the last unsolved problem of continuum mechanics\".  [1]-[4].   The  anomalous dimension,  a property  of strongly interacting systems first introduced to  the quantum field theory by Gribov and Migdal [5] and Polyakov [6],    is a notoriously difficult from theoretical viewpoint  concept and  only recently, after many years of trying,  the theory of anomalous scaling of the structure functions of a  passive scalar advected by  a white-in-time random velocity field, has been developed [7]-[8].  Almost simultaneously,    the theory of bi-scaling in  turbulence generated by  the forced  Burgers equation [9]-[10]  was formulated. \\noindent The attempts to explain anomalous dimensions in three dimensional turbulence were made in Refs.[11]-[12].  It was shown that the  Navier-Stokes equations combined with a simple model for the pressure -velocity correlations , lead to  homogeneous differential equations for the structure functions and, as a result, to anomalous scaling exponents which cannot be obtained on dimensional grounds.  \\noindent  During last forty-fifty years, a  substantial effort was devoted to derivation or at least modeling of various  probability densities in turbulence. The first attempts  resulting in the Log-normal PDF  of the dissipation rate fluctuations, consistent with  the multi-scaling, were made  by Kolmogorov [13] and Yaglom [14], using a simple cascade model. This   result  was later criticized by Orszag [15] as,  in general, not   realizable.  Similar Log-normal PDF  was obtained for the  not too large magnitudes of  velocity increment $\\delta_{r}u=u(x+r)-u(x)$ in Ref. [11] , which was later experimentally tested by Kurien and Sreenivasan [3].  Some other attempts based on  analysis of experimental data led to various fits ranging from Log-normal and  Log-Poison expressions to exponential and \"stretched exponential \"  PDFs  $P(\\delta_{r}u)$.  In this paper we, addressing this problem, will restrict ourselves by  considerations based on equations of motion.  \\noindent This paper is organized in a following way. In the Section 2 we introduce the Mellin transform for the probability density of velocity increments  and define the large- scale  Gaussian boundary condition.  We  show that the \"normal  (linear ) scaling\" corresponds  to the Gaussaian PDF of the small-scale fluctuations.  The PDF accounting for small deviations from the linear scaling is calculated   and compared with traditional Log-normal expressions.  In Section 3, we based on  theoretical analysis of the Navier-Stokes equations and numerical simulations by Gotoh and Nakano [16],  justify the boundary conditions used for the derivation of Sec 2.  The random field of the  dissipation scales $\\eta$ (linear dimensions of the dissipative structures) , derived  from  the expression for the dissipation anomaly,  is introduced in Sec 4 and the PDF $Q(\\eta, Re)$ is computed in Section 5.  Discussions and conclusions are presented in Section 6. There, the exact  expression for the probability density of velocity increments  following the multifractal formalism is presented and one of the results of the multi-fractal theory is derived (Polyakov 2005) without the multifractal input.   Some of the concepts and relations  presented  below have been reported in the previous publications [12],  [17]-[18].  Wherever needed, we  repeat them here  for the sake of continuity and clarity. \\\\ ",
        "conclusions": "the small-scale dynamics are strongly coupled to the large-scale phenomena. This  may be  a reason for   a serious reexamination of  the very concept of the turbulence energy cascade which, within the framework of the present development,    seem  neither possible nor needed. An accurate experimental and numerical comparison of the measured and Log-normal PDFs of velocity increments may be extremely important.  \\noindent It has been shown  [18] ( in a different way this is also  an element of the multi-fractal theory)   that  the  scales $\\eta$ form  a random field not necessarily related to  the energy dissipation scales but rather to the linear dimensions of  various dissipation structures defined by the local value of the Reynolds number $Re_{\\eta}=O(1)$. Some of these  structures are responsible for the energy and  the second-order moment $S_{2,0}$ dissipation, while others, more powerful,  ~ for the dynamics of the higher -order moments. Thus, the scale $\\eta$ must be perceived as a dynamic cut-off separating analytic and singular components  of the velocity field. \\noindent The probability density $Q(\\eta,Re)$, calculated in Section 5 is an interesting and easily  measurable quantity. A note of caution is in order:  to make reliable calculations or measurements  of the moments involving spatial derivatives of velocity field,  one has to have a field resolved well enough to exhibit at least a fraction of the analytic range of the corresponding moments of velocity increment. For example, according to (18), the second order  moment of Lagrangian acceleration is proportional to the sixth order structure function calculated on the scale $\\eta_{6}$, while the fourth order one is expressed  through $S_{12}(\\eta_{12})$. In the high Reynolds number flows, the measurements of the twelveth  order structure function including analytic range where $S_{12}\\propto r^{12}$ do not exist. The situation with the moments of dissipation rate is even worse: fourth- order moment  $\\overline{{\\cal E}^{4}}$ is  related to $S_{16}(\\eta_{16})$.  An interesting possibility is being explored by Schumacher [28 ]  running very large ($1024^{3}$) numerical simulations at reasonably low Reynolds numbers $R_{\\lambda}\\approx 10-60$. Analyzing the probability density of the dissipation scales Schumacher and Sreenivasan [29] obtained $Q(\\eta,Re)$ very similar to one shown on Fig. 3.  \\noindent The results presented here were obtained from analysis,  both theoretical  and numerical, of the dynamic equations. No multi-fractal assumptions have been made.   Still, it is interesting to compare the two approaches. In its present form, the multi-fractal (MF) theory consists of two parts [1]. The first one based on an idea of fractal dimension,  attempts to explain the origin of anomalous scaling by assuming  \\begin{equation} S_{p}(r)=(\\frac{r}{L})^{\\xi_{p}}=\\int d\\mu(h)(\\frac{r}{L})^{ph+3-D(h)} \\end{equation} \\noindent The normalised structure functions $S_{p}=(\\frac{\\delta_{r}u}{u_{rms}})^{p}$. The $O((r/L)^{ph})$ term comes from the multi-fractal assumption \\begin{equation} \\delta_{r}u=(r/L)^{h} \\end{equation}  \\noindent defined on a set of fractal dimension $D(h)$ where $h$ is a value of the scaling exponent from the interval  $h_{min}\\geq h\\geq h_{max}$ and $(r/L)^{3-D(h)}$ is the probability of being in the interval $r$ in a volume of dimension $3-D(h)$.   Neglecting the Logs in the steepest descent evaluation of the integral (26),  gives  a relation between the fractal dimension and the exponents $\\xi_{n}$ (For the review see Frisch [1]).  \\noindent  {\\bf  The multifractal PDF  and Mellin transform.} Let us establish  possible relations between the two theories.  Using the gaussian expression for the amplitudes, we have:   \\begin{equation} S_{2p}(r)=(2p-1)!!(\\frac{r}{L})^{\\xi_{2p}}=(2p-1)!!\\int d\\mu(h)(\\frac{r}{L})^{2ph+3-D(h)} \\end{equation}  \\noindent Substituting this into (1) and repeating the calculations of the section 2 gives ($L=1$):  $$P(u,r)=\\frac{1}{2}\\int e^{-x^{2}}dx \\int_{-i\\infty}^{i\\infty}dn \\int \\frac{d\\mu(h)}{dh}dhe^{-n(\\ln\\frac{u}{\\sqrt{2}x}-h\\ln r)}e^{(3-D(h))\\ln r}$$  \\noindent Integration over $n$ gives the  delta -function   $\\delta(h\\ln r-\\ln \\frac{u}{\\sqrt{2}x})$ ($u/x\\propto r^{h}$ )  with the finial result:  $$P(u,r)=\\frac{1}{2u \\ln r }\\int  dx e^{-x^{2}}exp (3-D(\\frac{\\ln \\frac{u}{\\sqrt{2}{x}}}{|\\ln r|})) \\frac{d\\mu(h_{*})}{dh}$$  \\noindent where $h_{*}=\\frac{\\ln \\frac{u}{\\sqrt{2}{x}}}{|\\ln r|}$. Restricitng ourselves by the relation (3) for the exponents and comparing this formula  with the PDF  (5) gives (neglecting the $\\mu$-factor ): \\begin{equation} [3-D(\\frac{\\ln \\frac{u}{\\sqrt{2}{x}}}{|\\ln r|})]\\ln r = -\\frac{ (\\ln \\frac{u}{r^{a}\\sqrt{2}x})^{2}}{4b|\\ln r|} \\end{equation}  \\noindent  No  steepest descent approximation has been used  in deriving this relation. If in accord with the MF theory,  we set  the amplitudes $A(n)=1$ and $\\sqrt{2}x=1$  in (29) and taking into account that $h=\\ln u/\\ln r$, the expression (29) gives:  $$3-D(h)=\\frac{ (\\ln \\frac{u}{r^{a}})^{2}}{4b(\\ln r)^{2}} =\\frac{(h-a)^{2}}{4b}$$  \\noindent  We remind the reader that in accord with Ref. [11],[12], $a\\approx 0.383$ and $b\\approx 0.0166$. This derivation which did not involve the steepest descent evaluation of the integral does not have  a  \"Log-problem\",  discussed in Ref. [30].   The experimental measurements of Cramer function $f(\\alpha)=D(h)+2$ for $h=\\alpha/3$ by Meneveau and Sreenivasan [31] (See also Ref.[1]) are in an extremely close agreement with this expression. The quantitative  differences are: the maximum of the calculated  curve  (with $\\alpha=3h$) is at $\\alpha=1.15$ (instead of $\\alpha=1$ of Ref. [31]) and $f(\\alpha)=0$ at $\\alpha_{1}=0.369$ and $\\alpha_{2}=1.92$ compared with $\\alpha_{1}\\approx 0.5$ and $\\alpha_{2}\\approx 1.8$ of Ref.[31]. The small deviations  come  from  the difference between  $h=\\alpha/3$ used for analysis of experimental data and  the theoretically obtained $h=a=0.383$. This difference decreases if the coordinates are rescaled by factor $0.383/0.333$.   \\noindent  The second part of the MF theory,  dealing  with the small-scale properties of turbulence, is based on the relation (Paladin and Vulpiani [32] ) \\begin{equation} \\frac{\\eta}{L}\\approx Re^{-\\frac{1}{1+h}} \\end{equation}  \\noindent  obtained by combining the MF assumption (27) and the outcome of the balance of the advective and viscous contributions to the Navier-Stokes equations. All  small-scale results derived using the MF formalism are numerically indistinguishable from the ones obtained both above and in the Ref. [18]. To illustrate this point,  we present an  alternative derivation of  the moments of velocity derivative  due to Polyakov [33].  \\noindent {\\bf Polyakov's derivation}. The probability density of velocity difference in the inertial range is:  \\begin{equation} P(\\delta_{r}u,r)=<\\delta(\\delta_{r}u-[u(x+r)-u(x)])> \\end{equation} \\noindent and in the dissipation (analytic)  range \\begin{equation} P_{D}(\\delta_{r}u,r)=<\\delta(\\delta_{r}u-u'r)> \\end{equation} \\noindent Assuming  the two PDFs match at the scale $r=\\eta=\\nu/\\delta_{\\eta}u\\equiv \\nu/u$, introduced   by (15), we have:  \\begin{equation} P_{D}(u,\\eta)=P(u,\\eta)=<\\delta(u-u'(x)\\nu/u)>=\\int dn A(n)\\nu^{\\xi_{n}}u^{-\\xi_{n}-n-1} \\end{equation}  \\noindent Multiplying (33) by $u^{2k}$ and integrating  over $u$  gives:  \\begin{equation} \\frac{1}{2}\\overline{(\\nu u')^{k}}=\\int dn A(n)\\nu^{\\xi_{n}}\\frac{1}{2k-\\xi_{n}-n} \\end{equation}  \\noindent  The integral is evaluated at a pole where $\\xi_{n(k)}+n=2k$ giving  \\begin{equation} \\overline{(u'(x))^{k}}\\propto A(n(k))\\nu^{\\rho(k)} \\end{equation}  \\noindent with  \\begin{equation} \\rho(k)=\\xi_{n(k)}-k=k-n(k) \\end{equation} \\noindent which is identical to  the formula (8.76) of Frisch's book [1] obtained using multi-fractal theory. Comparing the relations (35) and (22) we find that , on the accepted magnitudes of  exponents $\\xi_{n}$,   numerically they  are basically identical for not too large moment numbers $n$.  It is also easy to see that if  in the limit $n\\rightarrow\\infty$, $\\xi_{n}\\propto n^{\\alpha}$ with $0\\leq \\alpha\\leq 1$, the two relations have the same asymptotics.  \\noindent The only approximation involved in derivations of  both relations  (22) and (35), presented in this paper, is the choice of the cut-off $\\eta=\\nu/u$ instead of $\\eta=O(\\nu/u)$. In reality, there exist a random field of the dissipation scales described by the probability density $Q(\\eta,Re)$ given by (25). Thus, a more accurate calculation of both (22) and (34) must involve averaging over the fluctuating cut- off $\\eta$. However, we do not expect this procedure to introduce substantial modifications of the obtained results.  \\noindent  The  theory presented here does not involve any multi-fractal assumptions.  Still, the quantitative (numerical)  agreements between the two approaches hints on the possibility of some qualitative connection.  The essentially dynamic theory developed here couples  the velocity fluctuations at the largest  and  smallest scales. One may speculate that  if a typical structure is basically a strongly convoluted sheet with two $O(L)$  linear dimensions  and  $O(\\eta)$ the third one,  then these  structures can loosely be identified with the multi-ifractal sets of the MF theory."
    },
    "0512/astro-ph0512415_arXiv.txt": {
        "abstract": "We seek to understand whether the stellar populations of galactic bulges show evidence of secular evolution triggered by the presence of the disc. For this purpose we re-analyse the sample of Proctor and Sansom, deriving stellar population ages and element abundances from absorption line indices as functions of central velocity dispersion and Hubble type. We obtain consistent constraints on ages from the three Balmer line indices \\Hb, \\Hg, and \\Hd, based on stellar population models that take the abundance ratio effects on these indices into account. Emission line contamination turns out to be a critical aspect, which favours the use of the higher-order Balmer line indices. Our derived ages are consistent with those of Proctor and Sansom based on a completely different method.  In agreement with other studies in the literature, we find that bulges have relatively low luminosity weighted ages, the lowest age derived being $1.3\\;$Gyr. Hence bulges are not generally old, but actually rejuvenated systems. We discuss evidence that this might be true also for the bulge of the Milky Way. The data reveal clear correlations of all three parameters luminosity weighted age, total metallicity, and \\aFe\\ ratio with central velocity dispersion. The smallest bulges are the youngest with the lowest \\aFe\\ ratios owing to late Fe enrichment from Type Ia supernovae. Using models combining recent minor star formation with a base old population, we show that the smallest bulges must have experienced significant star formation events involving $10-30\\;$ per cent of their total mass in the past $1-2\\;$Gyr. No significant correlations of the stellar population parameters with Hubble Type are found.  We show that the above relationships with $\\sigma$ coincide perfectly with those of early-type galaxies.  In other words, bulges are typically younger, metal-poorer and less \\aFe\\ enhanced than early-type galaxies, because of their smaller masses. At a given velocity dispersion, bulges and elliptical galaxies are indistinguishable as far as their stellar populations are concerned. These results favour an inside-out formation scenario and indicate that the discs in spiral galaxies of Hubble types Sbc and earlier cannot have a significant influence on the evolution of the stellar populations in the bulge component. The phenomenon of pseudobulge formation must be restricted to spirals of types later than Sbc. ",
        "introduction": "\\label{sec:intro} Pseudobulges are bulges formed out of disc material in secular processes \\citep{Kormendy82}. As laid out in detail in the recent review by \\citet{KK04}, they are 'are not just dust features or the outer disk extending inside a classical bulge all the way to the centre', but appear to be built by nuclear star formation. In other words, the presence of pseudobulges in spiral galaxies should be detectable through fingerprints of relatively recent star formation in their stellar populations. This seems to stand in clear contrast to the commonly accepted perception that bulges are generally old \\citep{Renzini99}. There is strong and compelling evidence that the bulk of stellar populations in the Milky Way Bulge are old without significant amounts of recent star formation \\citep{Ortetal95,Renzini99,FWS03,Zocetal03}. On the other hand, blue colours, patchy dust features, and low surface brightnesses are found predominantly in bulges of later type spirals \\citep{deJong96b,Peletal99a}. Moreover, bulge colour and disc colour appear to be correlated pointing toward the presence of secular evolution processes \\citep{PB96,WGF97,Peletal99a,GdA01}. Based on these indications, \\citet{KK04} conclude that stellar populations are at least consistent with the expectation that the latest type galaxies must have pseudobulges.  The aim of this paper is to look into this in more detail, and to search for fingerprints of recent star formation in bulges by deriving average ages and abundance ratios of bulges along the Hubble sequence. We the use of the \\aFe\\ ratio as a measure for Type~Ia supernova enrichment, and hence late star formation.  We study the sample of \\citet[][hereafter PS02]{PS02} comprising 32 spheroids in a relatively large range of Hubble types from E to Sbc.  By means of abundance ratio-sensitive stellar population modelling \\citep*{TMB03a,TMK04}, we derive luminosity weighted ages, metallicities, and \\aFe\\ ratios of the central stellar populations (inner $\\sim 250\\;$pc) from a combination of metal indices (\\Mgb, Fe5270, Fe5335) and Balmer line indices (\\HdA, \\HgA, \\Hb). The resulting stellar population parameters are compared with the results obtained by PS02, and then analysed with two-component models with the aim to quantify the possible contribution of recent star formation on the basis of a generally old population.  An additional constraint will be set by the direct comparison with recent findings on early-type galaxies \\citep{Thoetal05} under the premise that a deviation of bulge properties from those of early-type galaxies may provide further hints on the possible presence of secular evolution and pseudobulge components.  The paper is organised as follows. After a brief summary of previous work in the literature on the stellar populations in bulges (Section~\\ref{sec:previous}), we will first present stellar population parameters ages, metallicities, and \\aFe\\ ratios derived from Mg-, Fe- and the three Balmer line indices and compare them with the results of PS02 (Section~\\ref{sec:parameters}). In Sections~\\ref{sec:comparison} and~\\ref{sec:sfhs} we will confront bulges with early-type galaxies and derive star formation histories to quantify the amount and epoch of possible rejuvenation events. The results are discussed in Sections~\\ref{sec:discussion} and~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} The main aim of this paper is to investigate whether the evolution of the stellar population in bulges is modified by the presence of the disc. We seek to understand whether secular evolution, and maybe the formation of pseudobulges, play an important role in the evolution of spiral galaxies. Our approach is to compare the stellar population properties in bulges with those in elliptical galaxies at a given central velocity dispersion, hence spheroid mass.  For bulges, a suitable sample is PS02 comprising 16 bulges in spirals with types Sa to Sbc, 6 lenticular and 11 elliptical galaxies. We derive luminosity weighted ages, metallicities, and \\aFe\\ ratios from one Balmer line index (\\HdA, \\HgA, or \\Hb\\ considered separately), and the metallic indices \\Mgb, Fe5270, and Fe5335 using the element ratio sensitive stellar population models of \\citet{TMB03a,TMK04}.  As there is relatively little overlap in $\\sigma$ between the bulges and the early-type galaxies in the PS02 sample, we compare the results with the sample of \\citet{Thoetal05}, which contains elliptical galaxies with velocity dispersion as low as $50\\;$km/s. For both samples we obtain very clear relationships between all three stellar population parameters and $\\sigma$.  \\citet{TMK04} demonstrate that the higher-order Balmer line indices are very sensitive to element ratios effects, and that consistent age estimates from \\Hb\\ and \\HgA\\ are obtained only when these effects are taken into account in the models. Here we extend this exercise to \\HdA, as the PS02 sample includes also this index.  We obtain very consistent estimates of ages, metallicities, and \\aFe\\ ratios from the three Balmer line indices. Emission line filling plays a critical role, impacting crucially on the scatter of the derived age-$\\sigma$ relation. The latter is smallest for \\HdA. Importantly, the ages and metallicities derived here using \\HdA\\ are extraordinarily consistent with those given by PS02.  Note that PS02 do not use a specific Balmer line, but perform a minimum $\\chi^2$ fit to all 25 Lick indices, an approach fully complementary to ours. While we use only the few line indices that we understand and model very well, PS02 average out the ignorance of the detail by using the maximum possible information available.  The excellent consistency found here is reassuring and suggests that these two rather orthogonal methods yield correct results. It should be emphasized that the $\\chi^2$ method is not dominated by the particular line indices used in the present study.  In agreement with other studies, we find that bulges have relatively low luminosity weighted ages, the lowest age derived being $1.3\\;$Gyr. Hence bulges are not overall old, but are actually rejuvenated systems. Interestingly, there is evidence that the bulge of the Milky Way also fits into this picture.  We find clear correlations of all three parameters luminosity weighted age, total metallicity, and \\aFe\\ ratio with central velocity dispersion, the smallest bulges being the youngest with the lowest \\aFe\\ ratios owing to late Fe enrichment from Type Ia supernovae.  We construct composite models in which a young subcomponent is superimposed over an underlying old population, in order to constrain the epoch and mass fraction of the rejuvenation event. We show that the smallest bulges must have experienced significant star formation events involving $10-30\\;$ per cent of their total mass in the past $1-2\\;$Gyr. Curiously, these results are not consistent with the age estimates for the Bulge of the Milky Way in the literature, which appears to have an overall old stellar population and no traces of recent star formation. We discuss new evidence that at least the central $\\sim 500\\;$pc of the Milky Way bulge contains a significant fraction of young stellar populations.  The comparison with the \\citet{Thoetal05} sample reveals that the above relationships with $\\sigma$ coincide perfectly with those of early-type galaxies.  In other words, bulges are typically younger, metal-poorer and less \\aFe\\ enhanced than early-type galaxies, only because of their smaller masses. At a given velocity dispersion, bulges and elliptical galaxies are indistinguishable as far as the basic properties of their stellar populations are concerned.  No significant correlations of the stellar population parameters with Hubble Type as late as Sbc are found, instead. In other words, the stars in bulges do not originate in the discs.  This result also agrees with the finding that structural parameters like disc and bulge scale lengths, as well as bulge-to-disc ratios, are correlated with bulge luminosity rather than with Hubble type.  If central spheroids have the same properties in galaxies with and without discs, this clearly favours inside-out galaxy formation \\citep{vdBosch98} according to which the disc forms after the bulge. Models that aim to explain the formation of bulges through disc fragmentation processes need to push the formation epoch to relatively high redshifts assuming high dissipation efficiencies \\citep{Imeetal04}.  Only in spiral galaxies of Hubble types later than Sbc discs can have a significant influence on the evolution of the stellar populations in the bulge component. This fits with the fact that Sersic index drops significantly in the transition between Hubble types Sbc and Sd \\citep{BGP05}. Secular evolution through the disc and the phenomenon of pseudobulge formation is most likely restricted to spirals of types Sc and later."
    },
    "0512/astro-ph0512309_arXiv.txt": {
        "abstract": "The collapse of a non-collisional dark matter and the formation of pancake structures in the universe are investigated approximately. Collapse is described by a system of ordinary differential equations, in the model of a uniformly rotating, three-axis, uniform density ellipsoid. Violent relaxation, mass, and angular momentum losses are taken into account phenomenologically. The formation of the equilibrium configuration, secular instability and the transition from a spheroid to a three-axis ellipsoid are investigated numerically and analytically in this dynamical model. ",
        "introduction": "According to modern cosmological ideas, most of the matter in the Universe is in the form of so-called cold dark matter, consisting of non-relativistic particles. The study of the formation of dark matter objects in the Universe is based on N-body si\\-mulations, which are very time consuming. In this situation, a simplified approach may become useful, because it allows one to investigate many different variants and to obtain some new principally important features of the problem, which could be lost or not visible during long numerical work.  The systematic analysis of ellipsoidal figures of equilibrium is presented in the classical book of \\citet{cha}. Different types of incompressible ellipsoids are investigated there on the basis of virial equations, as well as their dynamical and secular stability. In particular, the points of onset of the bar-mode dynamical and secular instabilities of the Maclaurin spheroid, for its transition to the Jacobi ellipsoid, are found. \\citet{LB64, LB65} applied the Chandrasechar's virial tensor method for a rotating, self-gravitating spheroid of pressure-free gas and showed the growth of non-axisymmetric perturbations during collapse. It was also discussed that the slowly shrinking Maclaurin spheroid will enter the Jacobi series if it shrinks slowly enough for the dissipative mechanisms to be operative. There have been numerical investigations of collapsing pressureless spheroids in the papers of \\citet{LB64} and \\citet{LMS}. The dynamics of a non-rotating sphere in two dimensions was considered in \\citet{LB79}, and it was shown that the pressure prevents the development of large-scale shape instability if initially the gravity is more than three-fifths pressure resisted. There is a wide-ranging review of this topic in the paper of \\citet{LB96} in memory of Chandrasekhar.  Subsequent investigations of rapidly rotating figures are connected with stars, and with large-scale structure of the Universe. In \\citet{RH91} the virial equations for rotating Riemann ellipsoids of incompressible fluid are demonstrated to form a Hamiltonian dynamical system. There is a detailed description of the ellipsoid model of rotating stars in the papers of \\citet*{lrs1, lrs4, lrs5}. Using a variational principle, they derived and investigated the equations for the evolution of a compressible Riemann-S ellipsoid, incorporating viscous dissipation and gravitational radiation. The solutions of these approximate equations permitted them to obtain equilibrium models, and to investigate their dynamical and secular stability. In a recent paper \\citep{shap04} secular bar-mode instability driven by viscous dissipation was investigated, using these equations, and the point of instability of compressible Maclaurin spheroids was found.  The modern theory of a large-scale structure is based on the ideas of \\citet{z70}, concerning the formation of strongly non-spherical structures during non-linear stages of the development of gravitational instability, known as `Zel'dovich's pancakes'. Numerical simulations for these objects have been performed subsequently by many groups. In the case of structures in dark matter, we are dealing with non-collisional non-relativistic particles, interacting only by gravitation. The development of gravitational instability and collapse in dark matter do not lead to any shock formation or radiation losses, but are characterized by non-collisional relaxation. This relaxation is based on the idea of a `violent relaxation' of \\citet{ref2}.  Here we derive and solve the equations for the dynamical behaviour in a simple model of a compressible uniformly rotating ellipsoid. We derive the equations for axes with the help of variation of the Lagrange function of an ellipsoid. Correct description of pressure effects, attained by such an approach, and the addition of relaxation permit us to obtain the dynamics of motion without any numerical singularities. In our model, motion along three axes takes place in the gravitational field of a uniform density ellipsoid, with account of the isotropic pressure, represented in an approximate non-differential way. Relaxation leads to a transformation of the kinetic energy of ordered motion into kinetic energy of chaotic motion, and to increases in the effective pressure and thermal energy. All losses are connected with runaway particles. The collapse of the rotating three-axis ellipsoid is approximated by a system of ordinary differential equations, where relaxation and losses of energy, mass and angular momentum are taken into account phenomenologically. The system is solved numerically for several parameters, characte\\-rizing the configuration. The approach in this work is similar to that used in \\citet{ref1}, where only dark matter spheroids ($a = b \\neq c$) were considered. In this case there are analytical formulae for the gravitational potential and forces. In the description of violent relaxation and different kinds of losses, we use the same approach as \\citet{ref1}. In the present paper we also find the point of the onset of secular instability of a compressible Maclaurin spheroid, using the derived dynamical equations, and also from the analysis of the sequence of equilibrium configurations. A simple analytical formula for the point of onset of the bar-mode secular instability of the Maclaurin spheroid is found. ",
        "conclusions": "We have investigated the dynamics of a three-axis dark matter ellipsoid. The equations of motion for the axes of a uniform compressible ellipsoid have been obtained by variation of the Lagrange function, in which violent relaxation and losses of matter, energy and angular momentum have been included phenomenologically.  The system was solved numerically, until the formation of stationary rotating figures in the presence of relaxation. For lower angular momentum $M$ we have the formation of a compressed spheroid, while at larger $M$ we follow the development of a three-axial instability and the formation of a three-axial ellipsoid. The instability in this approximation happens at the bifurcation point of the sequence of Maclaurin spheroids, where the Jacobi ellipsoidal system starts.  The bifurcation point coinciding with the point of loss of stability is found analytically in the form of a simple formula, by static and dynamic approaches. Numerical and analytical considerations give identical results.  The development of instability, connected with radial orbits, is obtained for slowly rotating collapsing bodies."
    },
    "0512/astro-ph0512623_arXiv.txt": {
        "abstract": "We present a spectroscopic survey using the MMT/Hectospec fiber spectrograph of 24~\\micron\\ sources selected with the \\spitzer\\ Space Telescope in the \\spitzer\\ First Look Survey.  We report 1296 new redshifts for 24~\\micron\\ sources, including 599 with $f_\\nu(24\\micron) \\geq 1$~mJy.   Combined with 291 additional redshifts for sources from the Sloan Digital Sky Survey (SDSS), our observing program was highly efficient and is $\\sim$90\\% complete for $\\isdss \\leq 21$~mag and $f_\\nu(24\\micron) \\geq 1$~mJy, and is 35\\% complete for $\\isdss \\leq 20.5$~mag and 0.3~mJy $\\leq f_\\nu(24\\micron) <$1.0~mJy.  Our Hectospec survey includes 1078 and 168 objects spectroscopically classified as galaxies and QSOs, respectively. Combining the Hectospec and SDSS samples, we find 24~\\micron--selected galaxies to $z_\\mathrm{gal}\\leq 0.98$ and QSOs to $z_\\mathrm{QSO}\\leq 3.6$, with mean redshifts of $\\langle z_\\mathrm{gal} \\rangle$=0.27 and $\\langle z_\\mathrm{QSO} \\rangle$=1.1. As part of this publication, we include the redshift catalogs and the reduced spectra; these are also available through the NASA/IPAC Infrared Science Archive.\\footnote{http://irsa.ipac.caltech.edu/\\label{footnote:url}} ",
        "introduction": "\\label{section:intro}  Observations with the Infrared (IR) Astronomical Satellite (\\textit{IRAS}) discovered that much of the bolometric emission associated with star--formation and active--galactic nuclei occurs in the thermal infrared.  The analysis of \\textit{IRAS} sources indicated that most ($\\simeq$70\\%) of the light emitted from local, normal galaxies comes at UV and optical wavelengths \\citep[\\eg][]{soi91}. However, measurements of the IR background found that the total far--IR emission ($\\lambda = 8-1000$~\\micron) of galaxies is comparable to that measured at UV and optical wavelengths \\citep[\\eg,][]{hau98}.  Therefore, over the history of the Universe, roughly half of the photons from star formation or black--hole accretion processes are emitted at IR wavelengths \\citep{elb02,dol06}.  Subsequent studies of IR number counts from the Infrared Space Observatory (\\textit{ISO}; Elbaz et al.\\ 1999) and more recently from the \\spitzer\\ Space Telescope \\citep{mar04,pap04} showed that IR--luminous galaxies have evolved rapidly, implying that they are a much more common phenomenon at high redshifts.  Studying the increase in the IR--active phases of galaxies requires measuring the properties of these objects as a function of redshift. Observations at 24~\\micron\\ from the multiband imaging photometer for \\spitzer\\ \\citep[MIPS,][]{rie04} are particularly well suited for such studies.  \\citet{soi87} concluded that starburst galaxies radiate as much as $\\sim$40\\% of their luminosity in the mid--IR (8--40~\\micron). The mid--IR emission from starforming galaxies correlates almost linearly with total IR luminosity over a range of galaxy type \\citep[\\eg,][]{spi95,rou01,pap02,cal05}.   The angular resolution of \\spitzer\\ at 24~\\micron\\ is roughly a factor of 3 and 7 better than that at 70 and 160~\\micron, respectively, allowing unambiguous source identification and probing the IR emission from many more sources than at the longer wavelengths.   Already, early studies of \\spitzer\\ 24~\\micron\\ sources with photometric redshifts over relatively small fields ($\\lsim 0.5$~sq.\\ deg) indicate that the bright end of the IR luminosity function evolves strongly from $z\\sim0$ to 1 \\citep{lef05,per05}.  To improve our understanding of the nature and evolution of IR-luminous phases of galaxies, we first need to construct large samples of objects with spectroscopic redshifts. In this paper, we publish the results of our survey with the MMT/Hectospec multi-fiber spectrograph in the \\spitzer\\ First Look Survey (FLS).  We report new redshifts for 1296 objects selected in the \\isdss--band and at 24~\\micron\\ over 3.3~deg$^2$.   Here we publish the catalogs and reduced, flux--calibrated spectra.   We are currently using these spectroscopic data in conjunction with surveys in other fields to study the evolution of the IR--luminous galaxy population.  We organize this paper as follows. In \\S~\\ref{section:data}, we discuss the \\spitzer\\ FLS dataset, our 24~\\micron\\ catalog, and our spectroscopic target selection.  In \\S~3 we describe the spectroscopic observations and data reduction.   In \\S~4 we discuss the spectroscopic completeness of the catalog.  In \\S~5  we present the Hectospec spectra and redshift catalog.  In \\S~6 we summarize our results and discuss the redshift distribution of  our sample of 24~\\micron\\ sources.  All magnitudes in this paper correspond to the AB system \\citep{oke83}, where $m_\\mathrm{AB} = 23.9 - 2.5 \\log ( f_\\nu / 1\\;\\mu\\mathrm{Jy} )$. ",
        "conclusions": "We have obtained 1296 redshifts for \\spitzer\\ 24~\\micron--selected sources in the \\spitzer\\ FLS using the Hectospec fiber spectrograph on the MMT.  Our observing program was highly efficient ($\\simeq$98--99\\% redshift success rate).  It is 90\\% complete for $f_\\nu(24\\micron) \\geq 1$~mJy and $\\isdss \\leq 21$~mag, and is 37\\% complete for 0.3~mJy $< f_\\nu(24~\\micron) <$ 1~mJy and $\\isdss \\leq 20.5$~mag. As part of this publication we provide catalogs for the full parent sample, and the SDSS and  MMT redshift catalogs.  We also publish our reduced, flux--calibrated spectroscopic Hectospec data.  Our spectroscopic survey of the \\spitzer\\ FLS identifies galaxies and QSOs over a large redshift range.  Figure~\\ref{fig:zhist} shows the redshift distribution for the 24~\\micron\\ sources from the primary, secondary, and tertiary samples in the FLS, including 1246 sources identified as galaxies and QSOs from our MMT/Hectospec survey, and 280 additional galaxies and QSOs with redshifts from SDSS.  Our Hectospec survey of 24~\\micron\\ sources identifies galaxies to $z\\leq 0.98$ and QSOs to $z\\leq 3.6$.   In figure~\\ref{fig:maxz}, we show the MMT/Hectospec spectra of the highest redshift objects spectroscopically classified as a galaxy and QSO.  \\begin{figure} \\epsscale{0.9} \\plotone{f6.eps} \\epsscale{1.0} \\caption\\figcapzhist \\end{figure}  \\begin{figure} \\epsscale{0.75} \\vbox{\\plotone{f7a.eps}} \\vspace{0.15in} \\vbox{\\plotone{f7b.eps}} \\caption\\figcapmaxz \\end{figure}  The redshift distribution for galaxies spans  $0.01 \\lsim z_\\mathrm{gal} \\lsim 1$, and the distribution for QSOs extends to higher redshift $0.1 \\lsim z_\\mathrm{QSO} \\lsim 5$.   The mean redshifts for galaxies and QSOs (including the Hectospec and SDSS data) are $\\langle z_\\mathrm{gal} \\rangle = 0.28$ and $\\langle z_\\mathrm{QSO} \\rangle = 1.1$, with standard deviations of $\\sigma_\\mathrm{gal} = 0.17$ and $\\sigma_\\mathrm{QSO} = 0.84$.   For the primary sample, the mean redshifts for galaxies and QSOs are $\\langle z_\\mathrm{gal} \\rangle = 0.24$ and $\\langle z_\\mathrm{QSO} \\rangle = 1.0$, with standard deviations of $\\sigma_\\mathrm{gal} = 0.17$ and $\\sigma_\\mathrm{QSO} = 0.77$.   Similarly, for the secondary sample, the mean redshifts for galaxies and QSOs are $\\langle z_\\mathrm{gal} \\rangle = 0.31$ and $\\langle z_\\mathrm{QSO} \\rangle = 1.5$, with standard deviations of $\\sigma_\\mathrm{gal} = 0.16$ and $\\sigma_\\mathrm{QSO} = 0.94$.  The photometric and redshift catalogs, and reduced spectra are available with the electronic edition of this publication.  They are also available through the NASA/IPAC Infrared Science Archive (IRSA; see footnote~\\ref{footnote:url}). We are currently pursuing further redshift surveys in \\spitzer\\ fields, building the datasets needed to understand the IR--active galaxy population."
    },
    "0512/astro-ph0512637_arXiv.txt": {
        "abstract": "The location of the solar dynamo is discussed in the context of new insights into the theory of nonlinear turbulent dynamos. It is argued that, from a dynamo-theoretic point of view, the bottom of the convection zone is not a likely location and that the solar dynamo may be distributed over the convection zone. The near surface shear layer produces not only east-west field alignment, but it also helps the dynamo disposing of its excess small scale magnetic helicity. ",
        "introduction": "It is commonly taken for granted that the solar dynamo has to work at the bottom of the convection zone, or that at least the toroidal field is generated or stored down there (Spiegel \\& Weiss 1980, Golub et al.\\ 1981, Galloway \\& Weiss 1981, Choudhuri 1990). This expectation results mostly from the fact that only at the bottom of the convection zone the dynamical time scales associated with convection and magnetic buoyancy are long enough to be comparable with the rotational period. There is also the notion that the magnetic field needs to be `stored' over a significant fraction of the solar cycle period and that this is only conceivable at or below the base of the convection zone. There are several other aspects in favor of placing the dynamo at the bottom of the convection zone. One is the large extent of the active regions (up to $100\\Mm$) that is only compatible with length scales typical of the deep convection zone (Galloway \\& Weiss 1981). Another argument is that it is at the bottom of the convection zone that we have a strong radial shear layer where $r\\partial\\Omega/\\partial r\\neq0$. However, there is of course also latitudinal differential rotation ($\\partial\\Omega/\\partial\\theta\\neq0$) that is actually stronger, and there is still extremely strong radial shear just beneath the surface in the uppermost $30\\Mm$ of the Sun (see \\Fig{bene99}). So, we see that the shear argument is problematic. In addition, at the bottom of the convection zone the sign of the radial shear is such that standard dynamo theory would predict an equatorward migration only when the $\\alpha$ effect is negative. Very near the bottom of the convection zone the $\\alpha$ effect is indeed predicted to have the opposite sign according to the standard formalism (Krivodubskii 1984). However, there is a whole host of other problems. First of all, the radial shear seen at the bottom of the convection zone is strongest at the poles and this is also where $\\alpha$ is strongest. So, in spite of spherical geometry factors the magnetic activity predicted by overshoot layer dynamos is far too strong at the poles and needs to be artificially suppressed if this approach is to be viable (R\\\"udiger \\& Brandenburg 1995, Markiel \\& Thomas 1999). Secondly, such overshoot layer dynamos (also sometimes called tachocline dynamos) have the well-known problem of producing too many toroidal field belts in the meridional plane (Moss et al.\\ 1990). Furthermore, the negative radial angular velocity gradient in the bulk of the convection zone and especially at the bottom tends to produce the wrong migration direction of the magnetic activity, i.e.\\ poleward rather than equatorward (Parker 1987). Although this problem could be fixed by invoking a strong negative value of $\\alpha$ at the bottom of the convection zone, there remains always the problem with the phase relation between radial and azimuthal fields, i.e.\\ $B_rB_\\phi$ is observed to be negative, but it would be positive with positive radial shear (Yoshimura 1976, Stix 1976).  \\begin{figure}[t!]\\centering \\includegraphics[width=0.70\\textwidth]{brandenburg_fig1}\\caption{ Radial profiles of the internal solar rotation rate, as inferred from helioseismology. The rotation rate of active zones at the beginning of the cycle (at $\\approx30^\\circ$ latitude) and near the end (at $\\approx4^\\circ$) is indicated by horizontal bars, which intersect the profiles of rotation rate at $r/R_\\odot\\approx0.97$. Courtesy of Benevolenskaya et al.\\ (1999). }\\label{bene99}\\end{figure}  Even if one ignored all these problems, there are still a number of difficulties associated with the idea of having a dynamo operating at the bottom of the convection zone. Firstly, in order for the flux tubes to be correctly oriented after their ascent, the field strength of the flux tube has to be very strong ($\\sim100\\kG$) to resist extraordinarily strong distortions and tilt. However, it is hard to imagine that the field strength exceeds the equipartition value ($\\sim1\\kG$) by a factor of a hundred, and this has not yet been demonstrated. Secondly, it is hard to imagine that the flux tubes would not disrupt by expanding too much before forming a neat sunspot pair.  These are problems and difficulties that we have been living with for quite a few years when constructing overshoot layer dynamos. However, there is also the possibility of placing the dynamo right in the middle of the convection zone. This idea may appear rather uncomfortable at first, but to people working in dynamo theory it is a rather natural and appealing scenario. The basic picture is one where dynamo action occurs in the bulk of the convection zone, affected obviously by the near-surface shear layer. Downward pumping will also operate, so as to prevent the magnetic field from floating upwards on too short a time scale (Nordlund et al.\\ 1992, Tobias et al.\\ 1998). However, in this scenario the field that we observe as sunspots at the surface is likely to come from the near surface layers, where sunspots may form as a result of convective collapse of magnetic fibrils (Zwaan 1978, Spruit \\& Zweibel 1979), possibly facilitated by negative turbulent magnetic pressure effects (Kleeorin et al.\\ 1996) or by an instability (Kitchatinov \\& Mazur 2000) causing the vertical flux to concentrate into a tube. The anticipated averaged field strength in the convection zone would be about $300\\G$, i.e.\\ about $10\\%$ of the equipartition value. This field may then get amplified locally near the surface. In that sense, sunspots are not be deeply rooted, but rather a shallow phenomenon rooted at a depth of $20$--$30\\Mm$. ",
        "conclusions": "The main point of this discussion is to stress that the solar dynamo may well work in the bulk of the convection zone. The near surface shear may not only be responsible for east-west alignment and toroidal field production, but it may also play a role in disposing of small scale magnetic and current helicities from the dynamo, e.g.\\ via coronal mass ejections (Blackman \\& Brandenburg 2003)."
    },
    "0512/astro-ph0512401_arXiv.txt": {
        "abstract": "{We present here the results of a deep (130 ks) XMM-Newton observation of the cluster of galaxies 2A~0335+096. The deep exposure allows us to study in detail its temperature structure and its elemental abundances. We fit three different thermal models and find that the multi-temperature {\\it{wdem}} model fits our data best. We find that the abundance structure of the cluster is consistent with a scenario where the relative number of Type~Ia supernovae contributing to the enrichment of the intra-cluster medium is $\\sim25$\\%, while the relative number of core collapse supernovae is $\\sim75$\\%. Comparison of the observed abundances to the supernova yields does not allow us to put any constrains on the contribution of Pop~III stars to the enrichment of the ICM. Radial abundance profiles show a strong central peak of both Type~Ia and core collapse supernova products. Both the temperature and iron abundance maps show an asymmetry in the direction of the elongated morphology of the surface brightness. In particular the temperature map shows a sharp change over a brightness edge on the southern side of the core, which was identified as a cold front in the Chandra data. This suggests that the cluster is in the process of a merger with a subcluster. Moreover, we find that the blobs or filaments discovered in the core of the cluster by Chandra are, contrary to the previous results, colder than the ambient gas and they appear to be in pressure equilibrium with their environment. ",
        "introduction": "\\label{intro}  Clusters of galaxies are the largest known gravitationally bound structures in the universe. According to the standard cosmological scenario they form and grow along the filaments through merging with groups and individual galaxies. Optical and X-ray studies reveal that clusters of galaxies are still forming at the present epoch.  The large effective area and superb spectral resolution of \\emph{XMM-Newton} together with the high spatial resolution of \\emph{Chandra} allow us to study clusters of galaxies with unprecedented detail. Analysis of data obtained by these satellites led to a number of important results in the recent years. In the central parts of many clusters of galaxies the gas density is high enough that the radiative cooling time of the gas is shorter than the age of the cluster. As the gas cools the pressure decreases, which causes a net inflow toward the center of the cluster. Many clusters of galaxies indeed show a temperature drop by a factor of three or more within the central 100 kpc radius \\citep[for a review on cooling flows see][]{fabian1994}. However, the spectra obtained by the XMM-Newton Reflection Grating Spectrometer (RGS) show no evidence for strong cooling rates of gas below $30-50$\\% of the maximum temperature of the ambient gas, which forces us to look for additional heating mechanisms in the cores of clusters \\citep[]{peterson2001,tamura2001a,kaastra2001}. The high resolution images from Chandra led recently to the discovery of {\\it{cold fronts}}, associated with motion of the cluster cores and to the identification of filamentary structure in the cores of a number of clusters \\citep{Markevitch2000}. Spatially resolved spectroscopy of many clusters of galaxies shows a strongly centrally peaked distribution of metal abundances \\citep{tamura2004}.  Because of their large potential wells, clusters of galaxies retain all the enriched material produced in the member galaxies. This makes them a unique environment for elemental-abundance measurements and for the study of the chemical enrichment history of the universe.  In this paper we study the X-ray bright cluster 2A~0335+096 using spatially-resolved and high-resolution spectra obtained during a 130 ks observation with the European Photon Imaging Camera \\citep[EPIC,][]{turner2001,struder2001} and the Reflection Grating Spectrometer \\citep[RGS,][]{herder2001} aboard \\emph{XMM-Newton} \\citep{jansen2001}. The properties of 2A~0335+096 allow us to address here some of the above mentioned issues.  2A~0335+096 was first detected as an X-ray source by \\emph{Ariel V} \\citep{Cooke1978}, and was found to be associated with a medium compact Zwicky cluster \\cite[]{Zwicky1965,Schwartz1980}. The presence of a cooling flow was first noted by \\citet{Schwartz1980} in the data obtained by \\emph{HEAO 1}. X-ray observations with \\emph{EXOSAT} \\cite[]{Singh1986,Singh1988} and \\emph{Einstein} \\citep{White1991} confirmed the presence of the cooling flow. Observations with \\emph{ROSAT} \\cite[]{Sarazin1992,Irwin1995} show a filamentary structure in the central region of the cooling flow. Observations with \\emph{ASCA} \\citep{Kikuchi1999} show a hint of a centrally peaked metallicity distribution. Using data obtained by \\emph{BeppoSAX}, \\citet{degrandi2001,degrandi2002} analyzed the metallicity and temperature profile of the cluster and found a centrally peaked metallicity gradient. Using the same dataset, \\citet{Ettori2002} estimated the total mass of the cluster within the region with overdensity of 2500 times the critical density to be $\\sim1.6\\times10^{14}$ M$_{\\odot}$,  while the mass of the gas was found to be $\\sim2.0\\times10^{13}$ M$_{\\odot}$.   Recent \\emph{Chandra} observation shows a complex structure in the core of the cluster: a cold front south of the center, unusual X-ray morphology consisting of a number of X-ray blobs and/or filaments on scales $\\gtrsim3$ kpc, along with two prominent X-ray cavities \\citep{Mazzotta2003}. Moreover, the \\emph{Chandra} observation shows that the cluster has a cool dense core and its radial temperature gradient varies with position angle. The radial metallicity profile has a pronounced central drop and an off-center peak \\citep{Mazzotta2003}. A previous shorter observation with \\emph{XMM-Newton} shows an increase of the Fe abundance toward the center with a strong central peak \\citep{tamura2004}. The central galaxy of 2A~0335+096 is a cD galaxy with a very extended optical emission line region (H$\\alpha+$[NII]) to the northeast of the galaxy. Moreover, the central region of the galaxy is anomalously blue, indicating recent star formation \\cite[]{Romanishin1988}. \\citet[]{edge2001} reports a detection of CO emission (implying $2\\times10^{9}$ M$_{\\odot}$ of molecular gas) and IRAS 60 $\\mu$m continuum. These observations indicate a mass deposition rate of a cooling flow of $<5$ M$_{\\odot}$ yr$^{-1}$. A radio study of 2A~0335+096 shows a radio source coincident with the central galaxy, which is surrounded by a mini-halo \\citep{sarazin1995}.  Throughout the paper we use $H_{0}=70$ km$\\, $s$^{-1}\\, $Mpc$^{-1}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$, which imply a linear scale of 42~kpc\\, arcmin$^{-1}$ at the cluster redshift of $z=0.0349$. Unless specified otherwise, all errors are at the $1\\sigma$ confidence level. ",
        "conclusions": "We have analyzed spatially resolved and high resolution spectra of the cluster of galaxies 2A~0335+096 obtained during a deep XMM-Newton observation. We found that: \\begin{itemize} \\item We unambiguously detect multi-temperature structure in the core of 2A~0335+096: for the determination of elemental abundances, a multi-temperature model is mandatory. The {\\it{wdem}} model is a good description of the temperature structure of the gas in 2A~0335+096. \\item The blobs/filaments found in the core of 2A~0335+096 are significantly colder than the ambient gas and they appear to be in pressure equilibrium with their environment. \\item We detect a strong central peak in the abundance distributions of both SN~Ia and SN~II products. \\item The relative number of SN~Ia contributing to the enrichment of the intra-cluster medium in the central 130~kpc of 2A~0335+096 is $\\sim25$\\%, while the relative number of SN~II is $\\sim75$\\%. Comparison of the observed abundances to the supernova yields does not allow us to put any constrains on the contribution of Pop~III stars to the enrichment of the ICM. We observe significantly higher calcium abundance than predicted by supernova models. \\item The detected asymmetry in the temperature and iron abundance distribution further supports the merger scenario, in which the subcluster had a higher metallicity.    \\end{itemize}"
    },
    "0512/astro-ph0512292_arXiv.txt": {
        "abstract": "We analyze the gas accretion flow through a planet-produced gap in a protoplanetary disk. We adopt the alpha disk model and ignore effects of planetary migration. We develop a semi-analytic, one-dimensional model that accounts for the effects of the planet as a mass sink and also carry out two-dimensional hydrodynamical simulations of a planet embedded in a disk.  The predictions of the mass flow rate through the gap based on the semi-analytic model generally agree with the hydrodynamical simulations at the 25\\% level. Through these models, we are able to explore steady state disk structures and over large spatial ranges. The presence of an accreting $\\sim1 M_J$ planet significantly lowers the density of the disk within a region of several times the planet's orbital radius. The mass flow rate across the gap (and onto the central star) is typically 10\\% to 25\\% of the mass accretion rate outside the orbit of the planet, for planet-to-star mass ratios that range from $5\\times10^{-5}$ to $1\\times10^{-3}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512547_arXiv.txt": {
        "abstract": "We have performed a pilot $Chandra$ survey of an off-center region of the Coma cluster to explore the X-ray properties and Luminosity Function of normal galaxies.   We present results on 13 $Chandra$-detected galaxies with optical photometric matches, including four spectroscopically-confirmed Coma-member galaxies. All seven spectroscopically confirmed giant Coma galaxies in this field have detections or limits consistent with low X-ray to optical flux ratios (${f_{X}\\over{f_{R}}}<10^{-3}$). We do not have sufficient numbers of X-ray detected galaxies to directly measure the galaxy X-ray Luminosity Function (XLF).  However, since we have a well-measured {\\it optical} LF, we take this low X-ray to optical flux ratio for the 7 spectroscopically confirmed galaxies to translate the optical LF to an XLF.  We find good agreement with Finoguenov et al. (2004), indicating that the X-ray emission per unit optical flux per galaxy is suppressed in clusters of galaxies, but extends this work to a specific off-center environment in the Coma cluster. Finally, we report the discovery of a region of diffuse X-ray flux which might correspond to a small group interacting with the Coma Intra-Cluster Medium (ICM). ",
        "introduction": "}  The launch of new X-ray observatories over the last few years has extended the study of X-ray emission from galaxies due to accreting binaries and hot interstellar medium (ISM) to cosmologically interesting distances. \\citep[e.g.,][]{Horn01}.  Recently, the first X-ray Luminosity Function (XLF) for normal galaxies (non-AGNs) was constructed at $z\\approx0.3$ and $z\\approx0.7$ \\citep[][]{Norman2004,Ptak2006} using data from the two deepest extragalactic X-ray surveys \\citep[e.g.,][]{davocatalog}. Moreover, a number of major studies from {\\it observationally complete} samples (i.e., observed to some fixed luminosity sensitivity) of X-ray detected/selected normal/star-forming galaxies have revealed that X-ray emission closely traces star formation rate in galaxies \\citep[e.g.,][]{BauerXII,Ranalli03,GrimmSFR,Georgakakis04}.  We have yet to explore galaxies in their most typical environments in the X-ray band -- this is the cluster and group environment, where most of the galaxies in the current universe are found \\citep[][]{Mulchaey03}. Also, in order to understand the X-ray properties of galaxies at relatively high redshifts, one needs a nearby control sample, with a well-known selection function. Until very recently, this was not possible as there have been relatively few complete samples of X-ray selected galaxies assembled at $z\\simlt0.1$. Recent work using wide-field optical surveys combined with $Chandra$ and $XMM$-$Newton$ archival data \\citep{Georgakakis04,HornSDSS} have reached $z\\approx0.1$ while the $ROSAT$ All-Sky Survey data are expected to provide an estimate of statistical properties of X-ray sources and their Luminosity Function (LF) in the local Universe\\citep[$\\simlt100$~Mpc; ][]{Tajer05}.  These surveys necessarily cover very large solid angles on the sky (several square degrees and larger) which require many pointings by the relevant X-ray missions (whose fields of view are typically 0.1-0.2 square degrees).  To understand X-ray properties of local galaxies, we employ a different strategy than the wider-field X-ray galaxy studies \\citep{Georgakakis04, HornSDSS,Tajer05} by performing a pilot $Chandra$ survey in the outskirts of the nearest rich cluster of galaxies, the Coma cluster \\citep[$z=0.023$;][] {Colless1996}.   The field is located $\\approx41$~arcmin ($\\approx1.2$~Mpc,$\\approx0.4$ virial radii) from the cluster center \\citep[see Figure 1; the virial radius is 2.9 Mpc, assuming $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$][]{Lokas2003}.  We target a cluster due to the large number of galaxies, making the X-ray observations more efficient. The reason for targeting a region away from the center of Coma is the lower X-ray intracluster medium (ICM) surface brightness there. One of our main goals is to assemble the XLF for normal galaxies down to the faintest X-ray limits possible. Thus far,  X-ray luminosity functions have been assembled in the field at high redshift \\citep[e.g,][]{Norman2004,Ptak2006}, for elliptical galaxies in the nearby Coma cluster of galaxies \\citep{Fino2004} and in an XMM-Newton survey of 2dF/SDSS fields at $z\\approx0.1$ \\citep{Georg2005}. $Chandra$ is ideal for this purpose, as it allows us to unambiguously resolve the galaxies from any residual ICM background.  The $Chandra$ observations cover approximately 20\\% of the full $\\approx 30^{\\prime} \\times 50^{\\prime}$ area in this part of Coma that has been under intense study. This region, and a corresponding region at the center of Coma, has extensive optical photometric and medium resolution (6--9\\AA) spectroscopic data for a well-defined sample of giant and dwarf galaxies \\citep{Mobasher2001spectra,Pogg2001spectra}. As a result, the {\\it optical} LFs for galaxies at both the core and outskirts of the Coma cluster was constructed. This allows for a direct comparison between the optical and X-ray LFs established for the same field.  Recently, using the XLF for bright ellipticals (over a large area) in the Coma cluster it was shown that the X-ray emission from Coma galaxies has undergone adiabatic compression by the surrounding ICM \\citep{Fino2004ICM}.  Furthermore, it was demonstrated that the X-ray activity of Coma-member galaxies is suppressed with respect to the field by a factor of 5.6 \\citep{Fino2004}.  However, detailed analyses of these elliptical galaxies, which were mainly located interior to our field (within 1~Mpc of the center of Coma), showed that their X-ray emission, although showing evidence of some compression by the surrounding ICM, was not largely different from field galaxies \\citep{Fino2004ICM}.    \\cite{Fino2004} suggested that it was not so much that the X-ray luminosity was suppressed but that the X-ray/optical flux {\\it ratio} was suppressed.   The reasons for this might include past gas stripping and/or past suppression of star formation in these Coma galaxies.  Throughout this paper we assume $H_{0}=70$~km s$^{-1}$ Mpc$^{-1}$. The Galactic neutral hydrogen column density is low in the direction of Coma \\citep[$N_{\\rm H} = 9.2 \\times10^{19}$~cm$^{-2}$; ][]{Stark92}. Finally, throughout this paper, we assume a distance modulus of 35.13 magnitudes \\citep{Mobasher2003LF}. ",
        "conclusions": "Through a 60~ks observation in the outskirts of the Coma cluster,  we have identified nine X-ray-detected galaxies whose optical colors indicate likely Coma membership (four with spectroscopic confirmation).  All seven spectroscopically-confirmed Coma members in this field have detections or limits consistent with very low X-ray/optical flux ratios.  We have thus confirmed the suppression of X-ray emission from galaxies in clusters (with respect to their optical emission) found previously in an XMM-Newton study.  The notable aspect of our result is that it is for an off-center region in the cluster. This may indicate that all of the impact of the cluster ICM on X-ray/optical flux ratios occurs at relatively low ram pressures at 1.1~Mpc from the cluster core.  This particular field would benefit from complete optical spectroscopic coverage to establish whether any of the $Chandra$ sources with higher X-ray/optical flux ratios might be members (and thus would run counter to the apparently lower X-ray/optical flux ratios for the current spectroscopic sample).  These sources {\\it could} be moderate luminosity AGN, which would be interesting due to the relatively high AGN fraction.  Our results indicate that an X-ray survey sensitive to $\\approx2\\times10^{-15}$~\\flux (0.5--2~keV) can begin to probe normal galaxies at the distance of Coma.  However, to obtain sufficient numbers of galaxies to directly constrain the XLF one would need to observe a larger area (3--4 times the area) or observe to fainter X-ray fluxes.   Such a wider-field survey would also build upon the existing and planned excellent multi-wavelength data in the field, including the recent  Spitzer IRAC (P.I. Hornschemeier) observations which cover the entire field, as well as the upcoming approved GALEX observations (P.I. Hornschemeier). Additionally, the Sloan Digital Sky Survey has recently surveyed parts of Coma; the first portions of these data will be available as part of SDSS DR5."
    },
    "0512/astro-ph0512221_arXiv.txt": {
        "abstract": "{We present the [$X$/H] trends as function of the elemental condensation temperature T$_C$ in 88 planet host stars and in a volume-limited comparison sample of 33 dwarfs without detected planetary companions. We gathered homogeneous abundance results for many volatile and refractory elements spanning a wide range of T$_C$, from a few dozens to several hundreds kelvin. We investigate possible anomalous trends of planet hosts with respect to comparison sample stars in order to detect evidence of possible pollution events. No significant differences are found in the behaviour of stars with and without planets. This result is in agreement with a ``primordial'' origin of the metal excess in planet host stars. However, a subgroup of 5 planet host and 1 comparison sample stars stands out for having particularly high [$X$/H] vs.\\ T$_C$ slopes. ",
        "introduction": "The announcement of the first planet orbiting around the solar type star \\object{51Peg} by Mayor \\& Queloz (\\cite{May95}) marked the beginning of a steadily growing series of extrasolar planet discoveries. Radial velocity programmes have now found more than 150 planetary systems in the solar neighborhood. The study of all these planets and their stellar parents opens new opportunities to understand the mechanisms involved in planetary formation and evolution. Current analyses are giving the first statistically significant results about the properties of the new systems (e.g.\\ Jorissen et al.\\ \\cite{Jor01}; Zucker \\& Mazeh \\cite{Zuc02}; Udry et al.\\ \\cite{Udr03}; Eggenberger et al.\\ \\cite{Egg04}).  The first strong link established between the planetary companions and their host stars is the fact that planet-harbouring stars are on average more metal-rich with respect to field stars. This idea has been put forward by Gonzalez (\\cite{Gon97}), while the first clear evidence has been published by Santos, Israelian \\& Mayor (\\cite{San01}). Further studies have been confirming this result as new planet host candidates have been discovered (e.g. Gonzalez et al.\\ \\cite{Gon01}; Laws et al.\\ \\cite{Law03}; Santos et al.\\ \\cite{San03}, \\cite{San04b}, \\cite{San05}; Fischer \\& Valenti \\cite{Fis05}; for a review see Santos et al.\\ \\cite{San04a}). This characteristic led to the suggestion that gas giant planet formation is favored by high stellar metallicity (Santos, Israelian \\& Mayor \\cite{San00}, \\cite{San01}), so that planetary systems would be more likely to form out of metal-enriched primordial clouds. Alternatively, the metallic excess in these stars may be attributed to the pollution by the late ingestion of planetary material (Laughlin \\cite{Lau00}, Gonzalez et al.\\ \\cite{Gon01}).  Several results supporting the ``self-pollution\" scenario were published. Murray \\& Chaboyer (\\cite{Mur02}) concluded that stochastic pollution by $\\sim$5$M_{\\oplus}$ of iron-rich material, together with selection effects and high intrinsic metallicity, may explain the observed metallicities in planet-harbouring stars. Israelian et al.\\ (\\cite{Isr01}, \\cite{Isr03}) found evidence for a planet (planets or planetary material) having been engulfed by the discovery of a significant amount of $^6$Li in the stellar atmosphere of the parent star \\object{HD\\,82943}. However, the amount of accreted matter was not enough to explain the global stellar metallicity. Ingestion of planetary material may also explain the lithium and iron enhancement found by Gonzalez (\\cite{Gon98}) and Laws \\& Gonzalez (\\cite{Law01}) in the primary component of the binary system 16 Cyg. However, most studies today suggest that a primordial origin is much likelier to explain the metallicity excess in planet host stars. Pinsonneault et al.\\ (\\cite{Pin01}) ruled out the ``self-pollution'' hypothesis since the iron excess did not show the expected T$_{\\rm eff}$ dependence. However, some different premises proposed by Vauclair (\\cite{Vau04}) might invalidate their argument. An additional point in favour of the primodial scenario is the fact that the frequency of planets is a rising function of [Fe/H] (Santos et al.\\ \\cite{San01}, \\cite{San04b}; Reid \\cite{Rei02}). Moreover, despite of their huge convective envelopes, giants with planets do not present [Fe/H] lower than other planet hosts. We cannot thus completely rule out any of the different hypotheses proposed to explain the metallicity excess in planet host stars. Note also the strange behaviour reported for [$\\alpha$/Fe] at supersolar [Fe/H], hard to be explained by Galactic chemical evolution (e.g. Bodaghee et al.\\ \\cite{Bod03}; Gilli et al.\\ \\cite{Gil05}; for a review see Israelian \\cite{Isr04a}, \\cite{Isr05}).  The abundance analyses of elements other than iron may give clues to this open question. Light elements are very important tracers of the internal structure and history of solar-type stars and therefore they can help to distinguish between different planet formation theories (Sandquist et al.\\ \\cite{Sand02}; Santos et al.\\ \\cite{San02}, \\cite{San04c}; Israelian et al.\\ \\cite{Isr04b}). Volatile (with lower condensation temperatures T$_C$) and refractory (with higher condensation temperatures T$_C$) elements can also give information about the role played by pollution events on the global stellar chemical composition. In fact, the elements of the former group are expected to be deficient in accreted materials relative to the latter. If the infall of large amounts of rocky planetary material was the main cause of the metallicity excess in planet host stars, as the ``self-pollution'' scenario claims, an overabundance of refractory elements with respect to volatiles should be observed. This would imply an increasing trend of abundance ratios [X/H] with the elemental condensation temperature T$_C$. However, the engulfment of a whole planet (or the rapid infall of planetary material) may avoid the evaporation of volatile elements before being inside the star, leading to no peculiar trends of the stellar abundances with T$_C$.  Smith, Cunha, \\& Lazzaro (\\cite{Smi01}) reported that a small subset of stars with planets exhibited an increasing [X/H] trend with T$_C$ and concluded that this trend pointed to the accretion of chemically fractionated solid material into the outer convective layers of these solar-type stars. They made use of the abundance results of 30 stars with planets reported by Gonzalez et al.\\ (\\cite{Gon01}) and Santos et al.\\ (\\cite{San01}), and compared them to those of 102 field stars from Edvardsson et al.\\ (\\cite{Edv93}) and Feltzing \\& Gustafsson (\\cite{Fel98}). Takeda et al.\\ (\\cite{Tak01})  also searched for a correlation between chemical abundances and T$_C$. They found that all volatile and refractories elements behave quite similarly in an homogeneously analysed set of 14 planet-harbouring stars and 4 field stars. Sadakane et al.\\ (\\cite{Sad02}) confirmed this result in 12 planet host stars, supporting a likelier primordial origin for the metal enhancement. Unfortunately, more results for planet-harbouring stars and an homogeneous comparison with field stars were needed to perform a more convincing test.  In this paper, we study the T$_C$ dependence of abundance ratios [$X$/H] uniformly derived in a large set of 105 planet host stars and in a volume-limited comparison sample of 88 stars without known planets. Some preliminary results were reported by Ecuvillon et al.\\ (\\cite{Ecu05b}). The large range of different T$_C$ covered by the analysed elements, which spans from 75 to 1600\\,K, permits us to investigate possible anomalies in targets with planets beside comparison sample stars, and to detect hints of pollution. Our results offer new clues to understand the relative contribution of fractionated accretion to the metallicity excess observed in planet host stars. ",
        "conclusions": "After gathering detailed and uniform abundance ratios of volatile and refractory elements in a large set of planet host stars and a volume-limited comparison sample of stars without any known planets, we derived the [$X$/H] vs.\\ T$_C$ trend and the slope value corresponding to the linear fit for each target. Planet host stars present an average slope higher than the comparison sample. However, this characteristic is mainly due to chemical evolutionary effects, more evident in planet host stars because of their metal-rich nature. The obtained increasing trend of T$_C$ with metallicity is a consequence of the Galactic chemical evolution. There does not seem to appear any remarkable difference in the behaviour of stars with and without planetary companions. However, a subset of 5 planet host stars and 1 comparison sample target with slopes falling out the trends are proposed as possible candidates to exhibit the signature of fractionated accretion. The larger number of planet host stars with respect to the comparison sample and the statistical uncertainties affecting the slope values do not allow a conclusive interpretation. \\\\ Further evidence of pollution are investigated by looking for any possible dependence on T$_{\\rm eff}$ and planetary parameters. No clear trends emerge with the stellar T$_{\\rm eff}$, contrary to the results published by Gonzalez (\\cite{Gon03}). We did not observe any significant behaviour with the planetary mass, orbital period, separation and eccentricity. Similar results were obtained by Santos et al.\\ (\\cite{San03}) when looking for possible correlations between the physical parameters of the exoplanets and the metallicity excess of their parent stars. Some of the targets we reported as candidates for possible selective accretion show planetary parameters compatible with a ``self-pollution'' scenario.\\\\ In conclusion, these possible candidates for self-pollution have to be carefully analyzed and submitted to further tests, in order to confirm or reject this suggestion. On the whole, our results do not point to a solely ``primordial'' or ``self-pollution'' scenario to explain the observed trends. Although in most cases a mainly primordial origin of the metallic excess in planet host stars seems likelier, probably a more complex mechanism combining both scenarios may underlie the observations."
    },
    "0512/astro-ph0512017_arXiv.txt": {
        "abstract": "Several analyses of the microwave sky maps from the Wilkinson Microwave Anisotropy Probe (WMAP) have drawn attention to alignments amongst the low-order multipoles. Amongst the various possible explanations, an effect of cosmic topology has been invoked by several authors. We focus on an alignment of the first four multipoles ($\\ell = 2$ to $5$) found by Land and Magueijo (2005), and investigate the distribution of their alignment statistic for a set of simulated cosmic microwave background maps for cosmologies with slab-like topology. We find that this topology does offer a modest increase in the probability of the observed value, but that even for the smallest topology considered the probability of the observed value remains below one percent. ",
        "introduction": "Several recent analyses of the WMAP satellite maps have pointed out an unexpected degree of alignment between the low-order multipoles of the cosmic microwave background (CMB) anisotropy \\cite{TDH,DTZH,Eriksen2003,Eriksen2004a,Hansen,Biel,LM,CHSS,BMRT}. Various explanations have been put forward for these alignments, ranging from statistical fluke or foreground contamination through to a genuinely cosmological interpretation in terms of breakdown of statistical isotropy. Such a breakdown would be a natural consequence of the Universe possessing a non-trivial topology of characteristic scale comparable to the observable Universe (for a selection of cosmic topology review papers see Ref.~\\cite{reviews}).  In this paper, we do not seek to address the interpretation of the observational data, but rather aim to test whether or not slab-space cosmic topologies give rise to the kind of alignments that are tentatively reported to have been observed in the first-year WMAP data. The observational indication is that there exists a preferred direction for the low multipoles. For instance, Tegmark \\textit{et al.}~\\cite{TDH} and de Oliveira-Costa \\textit{et al.}~\\cite{DTZH} noted that the quadrupole and octupole were closely aligned with one another, and approximately planar. Land and Magueijo \\cite{LM} (hereafter LM) sought the alignment for each multipole $\\ell$ that maximized the proportion of power contributed by a single $m$ mode, and noted that the alignments of the first four multipoles were much closer than would be expected under statistical isotropy. These authors have all suggested that such alignments may be an indication of a slab topology where only one dimension is compact (finite and unbounded).  The principal aim of this paper is to determine whether the LM alignment is a prediction of slab-space cosmic topology. We simulate CMB maps for spatially-flat slab topologies, for different sizes of the compact dimension, and derive the statistics of the alignments as defined by LM. We find that the degree of alignment in the observed data remain anomalous even in slab-space topologies. ",
        "conclusions": "Our main results are as follows. We have confirmed the observed value of $\\hat{\\theta}$ found by Land and Magueijo, while noting that it is quite dependent on the choice of maps used. We have also confirmed their result that Gaussian skies have only about 0.1\\% chance of finding a value as low as that observed in the TOH maps. By analyzing a set of slab-topology maps, we have found that there is a slightly-enhanced probability of such a low value being obtained, but in absolute terms it remains extremely unlikely. We conclude that slab topology is not the explanation for the multipole alignment found by Land and Magueijo. The resolution must lie elsewhere, perhaps in other topologies, or instead in other cosmological assumptions, or in foreground or instrumental noise."
    },
    "0512/astro-ph0512367_arXiv.txt": {
        "abstract": "The effects of mass-varying neutrinos on cosmic microwave background (CMB) anisotropies and large scale structures (LSS) are studied. In these models, dark energy and neutrinos are coupled such that the neutrino masses are functions of the scalar field playing the role of dark energy. We begin by describing the cosmological background evolution of such a system. It is pointed out that, similar to models with a dark matter/dark energy interaction, the apparent equation of state measured with SNIa can be smaller than -1. We then discuss the effect of mass-varying neutrinos on the CMB anisotropies and the matter power spectrum. A suppression of power in the CMB power spectrum at large angular scales is usually observed. We give an explanation for this behaviour and discuss different couplings and quintessence potentials to show the generality of the results obtained.  We perform a likelihood analysis using wide-ranging SNIa, CMB and LSS observations to assess whether such theories are viable. Treating the neutrino mass as a free parameter we find that the constraints on the coupling are weak, since CMB and LSS surveys give only upper bounds on the neutrino mass. However, fixing a priori the neutrino masses, we find that there is some evidence that the existence of such a coupling is actually preferred by current cosmological data over the standard $\\Lambda$CDM cosmology. ",
        "introduction": "Recent cosmological observations indicate that the expansion of the universe is accelerating \\cite{observation1}-\\cite{observation3}. It follows from General Relativity that the dominant energy component today must have negative pressure. Many candidates have been proposed over the years, including scalar field models, which are well motivated from the point of view of particle physics theories, see e.g. \\cite{wetterich}-\\cite{wetterich2}. The main prediction of these types of models is that the dark energy equation of state becomes a dynamical quantity, and can vary from the usual value of $w=-1$ for a cosmological constant. Although such models are very attractive, they are plagued with several theoretical difficulties, such as the stability of the potential under quantum corrections \\cite{doran} or why the dark energy scalar field seems not to mediate a force between normal matter particles \\cite{carroll}.  In addition, the energy scale of the scalar field is put in by hand and usually not connected to a more fundamental energy scale. However, attempts have been made to address these problems, such as models with ultralight pseudo-Nambu-Goldstone bosons (see, for example, \\cite{frieman}, \\cite{nilles} and \\cite{kaloperdark}; for a review, see \\cite{pecceireview}).  It is expected that any explanation for dark energy will involve physics beyond the standard model of particle physics. Recently, a new class of models have been proposed, which entertain the idea of a possible connection between neutrinos and dark energy.  Their theoretical and observational consequences have already been studied very extensively \\cite{neutrinobeg}-\\cite{neutrinoend}. The main motivation for a connection between dark energy and neutrinos is that the energy scale of dark energy (${\\cal O}(10^{-3})$~eV) is of the order of the neutrino mass scale. In these models the neutrino mass scale and the dark energy scale are linked to each other, and hence the observed non-zero neutrino masses (see \\cite{massiveneutrinos1}-\\cite{massiveneutrinos3}) cannot be understood without an understanding of dark energy. Also, one may hope that these models might provide an explanation for the coincidence problem \\cite{fardon1}.  In this paper we investigate the cosmology of neutrino models of dark energy.  We take into account the {\\it full} evolution of the neutrinos, i.e. studying the relativistic and non-relativistic regimes and the transition in between.  Armed with a complete numerical model for the evolution of the coupled neutrinos, we compare the background evolution with Supernova data and study how the modified perturbations affect the cosmic microwave background radiation (CMB) temperature anisotropies and large scale structures (LSS) matter power spectrum. We thereby present the details of the results outlined in \\cite{us} and discuss other forms of coupling between dark energy and neutrinos and potentials for the dark energy field.  The paper is organized as follows: In Section II we discuss the background evolution of the coupled dark energy-neutrino system in the context of a typical quintessential potential.  In Section III we derive the evolution equations for cosmological perturbations in neutrino models of dark energy, and present the modified CMB and matter power spectra. In Section IV we discuss other couplings and potentials, such as inverse power-law potentials and field--dependent couplings. In Section V we compare our theory with data, using a public Markov-Chain Monte-Carlo data analysis program.  We conclude in Section VI. ",
        "conclusions": "We have investigated models of dark energy which couple a quintessence scalar field to massive neutrinos.  In these models, dark energy and neutrinos are coupled such that the neutrino masses become functions of the scalar field.  The effects of such models on the cosmological background evolution, on the cosmic microwave background anisotropies, and on the formation of large scale structures were analyzed. Additionally, we have also performed a likelihood analysis on the parameter space of such theories.  We have focused on two specific models: In the first, the coupling between neutrinos and dark energy is constant and the quintessential potential is an exponential. The second model, which is better motivated from the particle physics point of view, has a neutrino--coupling which depends on the quintessence field (hence changes with time), whilst the scalar field has a power-law potential. In spite of some specific differences between these two models (such as the energy density stored in the scalar field at recombination for example), the effects of the coupling on the CMB anisotropies and on the matter power spectrum are nevertheless explainable by the basic mechanisms that we have identified earlier. Namely, the coupling modifies the background history and induces an ISW contribution to the CMB spectrum; the matter power spectrum is modified by the magnitude of the neutrino mass during structure formation. Given the generality of these explanations, the conclusions drawn from this investigation could probably be applied to any similar model with a neutrino-dark energy coupling.  It is important to note that in our models, the dark energy sector is described by a \\emph{light} scalar field, with a mass which is at most of order $H$.  This is in contrast to previous models \\cite{fardon1} in which the mass of the scalar field is much larger than $H$ for most of its history. The latter can have significant effects upon the behaviour of the neutrinos and the growth of their perturbations, and which is difficult to reconcile with current astronomical data \\cite{zalda}.  Solving the collisionless Boltzmann equation for the neutrinos, we have investigated the relativistic and non-relativistic regimes and the transition period in between.  Initially the neutrinos are highly relativistic, and during this period the quintessence field is frozen. The mass of the neutrinos therefore remains constant. As the neutrinos become non-relativistic they begin to exchange energy with the quintessence field via the coupling term.  At a temperature scale comparable to the neutrino mass, the neutrinos become non-relativistic, whilst the quintessence field is dominated by kinetic energy.  It is at this point that the neutrino mass begins to evolve significantly.  The details of this behaviour and evolution depends on the choice of the coupling $\\beta$ and the potential parameter $\\sigma$.  In fact, the masses of the neutrinos can be heavier or lighter in the past depending on the choice of potential and coupling parameters.  The coupling of neutrinos to dark energy slightly alters the evolution of the cosmological background.  It was found that similarly to models with a dark matter/dark energy interaction, the apparent equation of state measured with Type Ia Supernovae at high redshift can be smaller than $-1$, without introducing phantom fields, and might even cross the boundary $w=-1$.  The most obvious modifications to the CMB anisotropy spectrum occur for large angular scales, with $\\ell<100$, where the dominant contribution to the anisotropies is generated by the Integrated Sachs-Wolfe Effect (IWS).  This arises due to the evolution of the gravitational potentials along the photon path from the surface of last scattering.  The modification to the cosmological background arising from the neutrino coupling can also have a significant effect upon the evolution of the perturbations.  We generally observe an increase in power for $10<\\ell<100$, whilst for $\\ell<10$ we find either an excess or reduction in power depending upon our choice of parameters. For the models where the neutrinos were much heavier in the past than today, we also observe a slight shift in the peaks and a modification in their relative amplitude.  The matter power spectrum exhibits free-streaming damping even in the presence of dark energy--neutrino coupling.  However, since the damping scale is mainly dependent on the value of the neutrino mass at the end of their relativistic stage, our results appear similar to the standard models with CDM and hot dark matter in which the mass is fixed at the relativistic plateau. It is obvious that the mass infered from the damping of the matter power spectrum is, in general, different from the neutrino mass measured with experiments in the laboratory.  We performed a likelihood analysis using SNIa, CMB and LSS datasets. Initially, we used the standard parameterization for our cosmological model, characterized by exponential dependence of the dark energy potential and neutrino mass on the scalar field. For a flat universe we varied all of the matter parameters, the Hubble constant, the initial power spectrum spectral index and amplitude and the instantaneous reionization parameter $z_{\\rm re}$. As expected, the cosmological data did not place strong constraints on our new coupling parameters.  This is no surprise, since it is well known that the current best fit analysis of cosmological data can only place an upper limit on the mass of the neutrino, and a zero neutrino mass is not excluded by most cosmological data sets. An interesting outcome was that couplings of order unity are perfectly acceptable with the actual data.  To proceed, we chose to perform an analysis using two values of the neutrino mass today, $m_\\nu=0.2$~eV and $m_\\nu=0.3$~eV, to investigate whether models of neutrino--dark energy coupling could in principle be constrained if neutrinos were independently confirmed to have a significant mass ($m_\\nu \\gtrsim 0.1$~eV), consistent with current experiments.  For both the $m_\\nu=0.2$~eV and $m_\\nu=0.3$~eV models we found that non-zero values of neutrino coupling strengths of order unity are preferred by the data.  We also saw that for these models a non--zero value for $\\sigma$ is preferred over the usual cosmological constant, although $\\sigma=0$ is not excluded at the $68\\%$ level.  Models with heavier neutrinos allow stronger constraints to be placed upon the strength of the coupling.  Indeed, for the $0.3$~eV neutrinos we found that neutrino--dark energy coupling is preferred at the $1$ sigma confidence level.  One should note that these constraints rely upon the assumption that the neutrino mass is known, and that the neutrinos have a mass $m_\\nu \\gtrsim 0.1$~eV. Although this assumption is consistent with current neutrino experiments, we can only make the statement that should the mass of the neutrino be found to be greater than $0.1$~eV, then current cosmological data can be used to constrain the strength of any neutrino--dark energy coupling."
    },
    "0512/astro-ph0512151_arXiv.txt": {
        "abstract": "We report on preliminary results of a hybrid non-LTE analysis of high-resolution, high-S/N spectra of the helium-rich subdwarf B star Feige\\,49 and the helium-poor sdB HD\\,205805. Non-LTE effects are found to have a notable impact on the stellar parameter and abundance determination. In particular the He\\,{\\sc i} lines show significant deviations from detailed balance, with the computed equivalent widths strengthened by up to $\\sim$35\\%. Non-LTE abundance corrections for the metals (C, N, O, Mg, S) are of the order $\\sim$0.05\\,--\\,0.25\\,dex on the mean, while corrections of up to $\\sim$0.7\\,dex are derived for individual transitions. The non-LTE approach reduces systematic trends and the statistical uncertainties in the abundance determination. Consequently, non-LTE analyses of a larger sample of objects have the potential to put much tighter constraints on the formation history of the different sdB populations than currently discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512198_arXiv.txt": {
        "abstract": "We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. In this paper we present the results for NGC~2158, an intermediate age, populous cluster. We have monitored the cluster for over 260 hours, spread over 59 nights. We have detected one candidate transiting low luminosity object, with eclipse depth of 3.7\\% in the $R$-band. If the host star is a member of the cluster, the eclipse depth is consistent with a 1.7 $R_J$ object. Cluster membership of the host is supported by its location on the cluster main sequence (MS) and its close proximity to the cluster center (2\\arcmin). We have discovered two other stars exhibiting low-amplitude (4-5\\%) transits, V64 and V70, but they are most likely blends or field stars.  Given the photometric precision and temporal coverage of our observations, and current best estimates for the frequency and radii of short-period planets, the expected number of detectable transiting planets in our sample is 0.13. We have observed four outbursts for the candidate cataclysmic variable V57. We have discovered 40 new variable stars in the cluster, bringing the total number of identified variables to 97, and present for them high precision light curves, spanning 13 months. ",
        "introduction": "} We have undertaken a long-term project, Planets in Stellar Clusters Extensive Search (PISCES), to search for transiting planets in open clusters. To date we have published a feasibility study based on one season of data for NGC~6791 (Mochejska et al.\\ 2002, hereafter Paper~I) and a catalog of 57 variable stars for our second target, NGC~2158, based on the data from the first observing season (Mochejska et al.\\ 2004, hereafter Paper~II). We have also published the results of an extensive search for transiting planets in NGC 6791, based on over 300 hours of observations, spread over 84 nights (Mochejska et al.\\ 2005, hereafter Paper III). We have not detected any promising candidates, and derived an estimate of 1.7 expected transiting planets.  In this paper we present the results of a search for transiting planets in the open cluster NGC~2158 $[(\\alpha,\\delta)_{2000}=(6^h7^m, +24^{\\circ}0'); (l,b)=(186\\fdg 63, +1\\fdg 78)]$. It is a very populous, intermediate age ($\\tau$=2-3 Gyr), rather metal poor ([Fe/H]=-0.46) open cluster, located at a distance of 3.6 kpc (Carraro et al.\\ 2002, hereafter Ca02; Christian, Heasley and Janes 1985).  Searching for planets in open clusters eliminates the problem of false detections due to blended eclipsing binary stars, which are a significant contaminant in the Galactic field searches (over 90\\% of all candidates; Konacki et al.\\ 2003; Udalski et al.\\ 2002a, 2002b). Blending causes a large decrease of the depth of the eclipses and mimics the transit of a much smaller object, such as a planet. As opposed to dense star fields in the disk of our Galaxy, open clusters located away from the galactic plane are sparse enough for blending to be negligible.  There are two key elements in a survey for transiting planets. The most commonly emphasized requirement is the high photometric precision, at the 1\\% level. The more often overlooked factor is the need for very extensive temporal coverage.  Extensive temporal coverage is important because even for planets with periods between 1 and 2 days, the fractional transit length is only $\\sim$5\\% of the period, and it drops to $\\sim$2\\% for periods 2-10 days. During the remaining 95-98\\% of the period the system is photometrically indistinguishable from stars without transiting planets. To our best knowledge, PISCES is the most extensive search for transiting planets in open clusters in terms of temporal coverage with a 1 m telescope.  The paper is arranged as follows: \\S 2 describes the observations, \\S 3 summarizes the reduction procedure, \\S 4 outlines the search strategy for transiting planets, \\S 5 gives an estimate of the expected number of transiting planet detections and \\S 6 contains the variable star catalog. Concluding remarks are found in \\S 7. ",
        "conclusions": "} In this paper we have performed an extensive search for transiting planets in the intermediate age, populous cluster NGC~2158. The cluster was monitored for over 260 hours during 59 nights.  We have identified a low-luminosity transiting object candidate, TR1. The $R$-band amplitude of 3.7\\% implies a 1.66 $R_J$ radius for the transiting companion. The location of the host star on the cluster MS and its proximity to the cluster center seem to indicate that it is a member of the cluster. Higher accuracy light curves are required to better costrain the radius and period of TR1, and followup spectroscopy, to estimate the mass of the transiting object, through a measurement, or an upper limit on the central object's radial velocity variations.  Assuming a planet frequency from radial velocity surveys, we estimate that we should have detected 0.13 transiting planets with periods between 1 and 10 days, with our photometric precision and temporal coverage. The main limitation on our detection efficiency was imposed by the photometric precision.  We have discovered 40 new variable stars in NGC 2158: 13 eclipsing binaries, 23 other periodic variables and four non-periodic variables, Together with 57 variables discovered in Paper II, this brings the total number of variables known in this cluster to 97. We have also presented high photometric precision light curves, spanning almost 13 months, for all previously known variables.  Transiting planets have proven to be more challenging to detect than initially expected, as shown by the paucity of detections from the many searches under way in open clusters (i.e.\\ Bruntt et al.\\ 2003; UStAPS: Hood et al.\\ 2005; EXPLORE/OC: von Braun et al.\\ 2004; STEPSS: Marshall et al.\\ 2005) and in the Galactic field (i.e.~EXPLORE: Mall{\\' e}n-Ornelas et al.\\ 2003; OGLE: Udalski et al.\\ 2002a; STARE: Alonso et al.\\ 2003; HAT: Bakos et al.\\ 2004; WASP0: Kane et al.\\ 2005\\footnote{For a more complete list of transiting planet searches, please refer to {\\tt http://star-www.st-and.ac.uk/\\~{}kdh1/transits/table.html} and {\\tt http://www.obspm.fr/encycl/searches.html}}). To date, only six planets have been discovered independently by transit searches, all of them in the field, and five of those were initially identified by OGLE (Udalski et al.\\ 2002a, 2002b, 2002c, 2003; Alonso et al.\\ 2004)."
    },
    "0512/hep-ph0512187_arXiv.txt": {
        "abstract": "I discuss the historical and conceptual roots of reasoning about the parameters of fundamental physics and cosmology based on selection effects.   I argue concretely that such reasoning can and should be combined with arguments based on symmetry and dynamics; it supplements them, but does not replace them. ",
        "introduction": "Our previous Rees-fest ``Anthropic Arguments in Fundamental Physics and Cosmology'' at Cambridge in 2001 had much in common with this one, in terms of the problems discussed and the approach to them.  Then as now the central concerns were apparent conspiracies among fundamental parameters of physics and cosmology that appear necessary to insure the emergence of life.  Then as now the main approach was to consider the possibility that significant observational selection effects are at work, even for the determination of superficially fundamental, universal parameters.  That approach is loosely referred to as anthropic reasoning, which in turn is often loosely phrased as the anthropic principle: the parameters of physics and cosmology have the values they do in order that intelligent life capable of observing those values can emerge.   That formulation upsets many scientists, and rightly so, since it smacks of irrational mysticism.  On the other hand, it is simply a true fact that intelligent observers are located only in a miniscule fraction of the world, and in places with special properties.  As a trivial consequence, probabilities conditioned on the presence of observers will differ grossly from probabilities per unit volume.   Much finer distinctions are possible, and useful; but I trust that this word to the wise is enough to it make clear that we shouldn't turn away from straightforward logic just because it can be made to sound, when stated sloppily, like irrational mysticism.  For all their commonality of content, the spirit pervading the two gatherings seemed quite different, at least to me.   One sign of the change is the different name attached to the present gathering.  This time it's ``Expectations of a Final Theory''.  The previous gathering had a defensive air.  It prominently featured a number of physicists who subsisted on the fringes, voices in the wilderness who had for many years promoted strange arguments about conspiracies among fundamental constants and alternative universes.   Their concerns and approaches seemed totally alien to the consensus vanguard of theoretical physics, which was busy successfully ;-) constructing a unique and mathematically perfect Universe.  Now the vanguard has marched off to join the prophets in the wilderness.    According to the new zeitgeist, the real world of phenomena must be consulted after all, if only to position ourselves within a perfect, but inaccessible, Multiverse.   Estimating selection effects, in practice, requires considerations of quite a different character than what we've become accustomed to in the recent practice of theoretical (i.e., hep-th) physics:  looser and more phenomenological, less precise but more accurate. ",
        "conclusions": ""
    },
    "0512/astro-ph0512616_arXiv.txt": {
        "abstract": "In this work we present a sensitive and systematic single-dish survey of CCS emission (complemented with  ammonia observations) at 1 cm, toward a sample of low- and intermediate-mass young star forming regions known to harbor water maser emission, made with NASA's 70 m antenna at Robledo de Chavela, Spain. Out of the 40 star forming regions surveyed in the CCS(2$_{1}$$-$1$_{0}$) line, only 6 low-mass sources show CCS emission: one transitional object between pre-stellar and protostellar Class 0 phase (GF9-2), three Class 0 protostars (L1448-IRS3, L1448C, and B1-IRS), a Class I source (L1251A), and a young T Tauri star (NGC2071-North). Since CCS is considered an ``early-time'' ($\\la 10^5$ yr) molecule, we explain these results by either proposing a revision of the classification of the age of NGC2071-North and L1251A, or suggesting the possibility that the particular physical conditions and processes of each source affect the destruction/production of the CCS. No statistically significant relationship was found between the presence of CCS and parameters of the molecular outflows and their driving sources. Nevertheless, we found a significant relationship between the detectability of CCS and the ammonia peak intensity (higher in regions with CCS), but  not with its integrated intensity. This tendency found may suggest that the narrower ammonia line widths in the less turbulent medium associated with younger cores may compensate for the differences in ammonia peak intensity, rendering differences in integrated intensity negligible. From the CCS detection rate we derive a lifetime of this molecule of $\\simeq$ (0.7$-$3)$\\times$10$^{4}$ yr in low-mass star forming regions. ",
        "introduction": "The lines of the CCS molecule are a powerful tool for studying the physical conditions and the structure of dark molecular clouds, because they are intense, abundant and not very opaque in these regions. In addition, CCS is very useful for performing dynamical studies, since it is heavier than other high density gas tracers and it has no hyperfine structure \\citep{Sai87,Suz92}.  Moreover, CCS is useful for obtaining information about the age of molecular clouds. Previous single-dish observations show that CCS lines are intense in starless, cold, quiescent cores, while ammonia tends to be abundant in star forming regions \\citep{Suz92}. It has been suggested that CCS is present in the first stages of molecular cloud evolution, but it is soon destroyed (on a timescale of $\\simeq$10$^{5}$ years, \\citealt{Mil90,Suz92}) after the formation of a dense core. This destruction process is induced by the core contraction \\citep{Suz92} that initiates star formation. On the other hand, when the molecular cloud evolves, the physical conditions and the chemical evolution in molecular cores favor the formation of other molecules like NH$_{3}$ \\citep{Suz92}. For this reason the abundance ratio [NH$_{3}$]/[CCS] has been considered as an indicator of the evolution of molecular cores. This time-dependent chemistry is responsible for the pronounced spatial anticorrelation observed in the emission from these two species, where ammonia tends to trace the inner regions and the CCS is found to be located outside, surrounding ammonia cores, in a clumpy distribution \\citep{Hir92, Kui96, Lai03}.  If this evolutionary trend (CCS more abundant in starless cores and NH$_{3}$ in more evolved, active star forming regions) is correct, any chance to use CCS as a tool to study star-formation processes will necessarily happen during the very first stages of the stellar evolution, i.e., Class 0 protostars. Interestingly, in these young protostars some of the key phenomena characteristic of the star formation (infall, disk formation and powerful mass-loss phenomena) are especially prominent and coeval, and have their own kinematical signature, which make kinematical studies of their environment especially interesting. In this early stage of stellar evolution, the presence of water maser emission at 22 GHz is rather common. This emission is considered a good tracer of mass-loss activity in young stellar objects (YSOs) in general \\citep{Rod80,Deb05}, and a good indicator of the age of low-mass YSOs, since they tend to be excited preferentially in the Class 0 stage \\citep{Fur01}.  The critical scales for these star formation processes are of the order of $\\simeq 100-1000$ AU in the case of low mass stars ($\\leq 1\\arcsec-10\\arcsec$ for nearby molecular clouds; \\citealt{Bel01,Che95,Cla98}). In order to study these phenomena with high enough resolution, it is necessary to carry out interferometric observations. This is feasible for strong maser lines, but very difficult for thermal emission (like CCS and ammonia) since it is usually weak and therefore requiring high sensitivity is needed. At wavelengths of $\\sim 1$ cm an instrument like the Very Large Array (VLA) provides high angular resolution, but at the expense of worse atmospheric phases that blur the emission. Weak sources, in which self-calibration is not possible, would not be detected, or the quality of their maps would be poor. Nevertheless, these kinds of problems have been addressed, in the case of continuum emission, by using the cross-calibration technique. This method involves the simultaneous observation of 1 cm continuum and strong water masers, and the use of the self-calibration solutions of the latter to correct phase and amplitude errors in the weaker continuum \\citep{Tor96, Tor97, Tor98}. This technique of cross-calibration using water masers has never been applied to spectral lines. The CCS(2$_{1}$$-$1$_{0}$)  transition at $\\sim$22 GHz is probably the best candidate to perform it, since it is only 110 MHz away from the H$_{2}$O (6$_{16}$$-$5$_{23}$) maser line.   A high-resolution study of the kinematics and the chemical evolution of the environment of the Class 0 source B1-IRS was carried out  by \\citet{deG05}. In that work, CCS, ammonia, and water maser emissions at 1 cm were observed with the VLA, although cross-calibration was not possible there, since the maser was not strong enough. Three different CCS clumps were detected, whose kinematical pattern was interpreted as gas interacting with the molecular outflow that exists in the region. This interaction was not observed using other molecular tracers or with lower-resolution observations of CCS in this source \\citep{Hir02,Lai03}. The evidence of interaction with the outflow led \\citet{deG05} to suggest the possibility that CCS abundance could be enhanced via shock-induced chemistry. Moreover, in that region  a spatial anticorrelation between CCS and ammonia at scales of $\\simeq$5$\\arcsec$ was observed for the first time, which illustrates the importance of time-dependent chemistry on small spatial scales.   In this work we present a sensitive and systematic single-dish survey of CCS emission (complemented with NH$_{3}$ observations) at 1 cm wavelength, toward low-mass star forming regions that are known to harbor water maser emission, using the NASA's 70m antenna at Robledo de Chavela, Spain. One of our main aims is to find the best candidates to make interferometric CCS studies as the one made in B1-IRS with the VLA, but applying the cross-calibration technique for obtaining high-quality maps. Interferometric maps of CCS, ammonia and water masers would allow us to study kinematical and physical properties of star forming regions at high resolution. Another purpose of the present survey is the search for the youngest star forming regions, assuming that the presence of water masers and CCS emission are signs of youth in star forming regions. Moreover, we present a search for the relation between the CCS emission and the physical characteristics of the star forming regions of the survey, focusing on the molecular outflow properties, considering the possible association between CCS emission and the outflow suggested for B1-IRS.  This paper is structured as follows: in section \\S 2 we describe our observations and data reduction. In section \\S 3 we show the survey results, as well as a short description of the sources detected. We discuss the results in \\S 4 and, finally, we summarize our conclusions in \\S 5. ",
        "conclusions": "We present a single-dish survey of CCS and NH$_{3}$ emission at 1 cm, carried out with NASA's 70m antenna at Robledo, toward a sample of low- and intermediate-mass young star forming regions that are known to harbor water maser emission. Our aim was to find the best candidates for a VLA high resolution study of the kinematical and physical properties of young Class 0 objects, to search for the youngest protostars, and to determine the relationship between CCS and NH$_{3}$ in star forming regions. Our general conclusions are the following: \\begin{itemize}  \\item We have detected 6 low-mass sources that show CCS emission, out of a sample of 40 star forming regions associated with water maser emission. Four of these (L1448C, L1448-IRS3, GF9-2, and L1251A) had not been previously detected in any CCS transition, another one (NGC 2071-North) has been detected for the first time with the CCS line at 1 cm, while the other one (B1-IRS) had already been reported to show CCS emission in this and other transitions. All our CCS detections also show ammonia emission.  \\item From the CCS detection rate, and the duration of the water maser emission in low-mass star forming regions, we derive a lifetime of this molecule of $\\simeq$(0.7$-$3) $\\times$ 10$^{4}$ yr in these regions, after star formation started.  \\item Three of the six sources detected in CCS are catalogued as Class 0 protostars (L1448-IRS3, L1448C and B1-IRS), one could be a transitional object between pre-stellar and protostellar Class 0 stage (GF9-2), and the last two sources (NGC2071-North and L1251A) are catalogued as more evolved, a young T Tauri and a Class I source, respectively. Since CCS is considered an early-time molecule, to explain the CCS detections in more evolved objects, we speculate with two possibilities: either the classification of NGC2071-North and L1251A should be revised, or the star formation activity and the physical properties of each cloud could influence in the production and destruction of the CCS molecules.   \\item We did not find any statistically significant trend that may relate the presence of CCS emission with different parameters of the molecular outflows or the central sources of these star forming regions.    \\item We found that the distribution and mean of the peak intensity of NH$_{3}$ in the group of sources detected in CCS are significantly different from those in the group of undetected ones, with the ammonia mean peak intensity higher in regions with CCS. However, no significant difference is found with respect to the integrated intensity of NH$_{3}$. Stronger NH$_{3}$ peak line intensities with indistinguishable integrated intensities suggests that the lines are narrower and the emitting regions less turbulent (i.e. younger).    \\item The linewidths of the CCS and NH$_3$ lines are noticeably different, with CCS being 3 times narrower in some cases. This suggest that emission from these lines arises from gas with different kinematical properties within the telescope beam. The cases with more similar widths (B1-IRS, L1448C) may trace strong interactions between molecular outflows and the CCS-emitting gas.   \\end{itemize}"
    },
    "0512/astro-ph0512566_arXiv.txt": {
        "abstract": "{We present a pilot study of atmospheres of accreting magnetic white dwarfs irradiated by intense fluxes at ultraviolet to infrared wavelengths. The model uses a standard LTE stellar atmosphere code which is expanded by introducing an angle-dependent external radiation source. The present results are obtained for an external source with the spectral shape of a 10000\\,K blackbody and a freely adjustable spectral flux. The model provides an explanation for the observed largely filled-up Lyman lines in the prototype polar AM~Herculis during its high states. It also confirms the hypotheses (i) that irradiation by cyclotron radiation and other radiation sources is the principle cause for the large heated polar caps surrounding the accretion spots on white dwarfs in polars and (ii) that much of the reprocessed light appears in the far ultraviolet and not in the soft X-ray regime as suggested in the original simple theories. We also briefly discuss the role played by hard X-rays in heating the polar cap. ",
        "introduction": "The white dwarfs in magnetic cataclysmic variables (mCVs, more specifically polars) accrete in restricted regions near their magnetic poles. In most of them, the infalling matter is heated in a free-standing shock above the surface of the white dwarf and cools by the emission of X-rays and cyclotron radiation. Early models assumed that the stand-off distance of the shock is small and much of the emission by the post-shock plasma is absorbed by the atmosphere below and immediately around the accretion spot. The implied irradiation of this small section of the atmosphere by intense fluxes of hard X-ray bremsstrahlung and cylotron radiation leads to an expected temperature of the heated atmosphere of $kT \\simeq 25$~eV ($T \\simeq 3\\times 10^5$\\,K), to re-emission in the soft X-ray regime with a quasi-blackbody spectrum, and to an energy balance $L_\\mathrm{sx} = f\\,[L_\\mathrm{cyc}+(1-A_\\mathrm{hx})\\, L_\\mathrm{hx}]$ \\citep{kinglasota79,lambmasters79}, where the three luminosity components refer to soft X-rays (sx), hard X-rays (hx), and cyclotron radiation (cyc), the geometry factor $f\\la1/2$ accounts for the fraction of the luminosity intercepted by the white dwarf, and the hard X-ray albedo $A_\\mathrm{hx}$ reduces the efficiency of X-ray heating \\citep{vanteeselingetal94}. Later on, it was realized that the accretion region is highly structured \\citep{kuijperspringle82, franketal88} with individual sections receiving vastly different mass flow densities and displaying substantially different emission characteristics. The plasma in tenuous tall shocks cools preferentially by cyclotron radiation, the plasma in dense low-lying shocks by hard X-ray bremsstrahlung, and for the highest mass flow densities, the shocks are buried in the atmosphere and the primary bremsstrahlung is reprocessed into soft X-rays, a component physically different from but observationally difficult to distinguish from the originally suggested soft X-ray blackbody component. This ``soft X-ray puzzle'' was the subject of an extended debate over more than two decades \\citep[e.g.,][ and references therein]{beuermann04, ramsaycropper04}.  Quite unexpectedly, observational support for the reprocessing scenario came from studies of the far ultraviolet spectra of the prototype polar AM~Herculis, which showed that photospheric emission of the white dwarf dominates the FUV \\citep{heiseverbunt88, gaensickeetal95-1, gaensickeetal98-2, maucheraymond98, greeleyetal99, gaensickeetal05-1}. The pronounced orbital modulation of the FUV flux indicates a large heated polar cap, which covers $\\sim 10$\\% of the white dwarf surface.  The cap reaches a peak temperature of $\\sim 3.5\\times 10^4$\\,K in low states when accretion nearly ceases and becomes much hotter in high states. \\citet{gaensickeetal95-1} demonstrated that the excess FUV flux of the polar cap quantitatively agrees with the sum of cyclotron and X-ray fluxes from AM~Her, both in the high and the low states, thus confirming the energy balance predicted by the simple models, with the decisive difference, however, that the heated polar cap covers a much larger area than the proper accretion spot and the reprocessed flux emerges in the FUV and not at soft X-ray wavelengths. AM~Her is the best observed polar, but FUV studies of other polars suggest that the situation encountered in AM~Her is quite typical of the class \\citep[e.g.,][]{araujobetancoretal05-1}.  The dominant FUV spectral feature of white dwarfs in polars are the Lyman absorption lines, which are deep and broad in low states, mimicking the pure hydrogen spectra of DA white dwarfs, but become filled up almost entirely in states of high accretion \\citep{heiseverbunt88, gaensickeetal95-1, gaensickeetal98-2, maucheraymond98, greeleyetal99, gaensickeetal05-1}. \\citet{gaensickeetal98-2} first tried to model the filled-up Lyman lines using a simple toy model and concluded that the temperature stratification in the white dwarf atmosphere is severely affected over the entire irradiated polar cap.  A spot due to reprocessing covering 10\\% of the white dwarf surface requires a point-like source at a height of 0.25 white dwarf radii or an extended source at a lower height. The best candidates are optical/infrared cyclotron and to some extent hard X-ray emission from tall shocks, supplemented by the FUV to optical emission from the pre-shock accretion stream. A further potential source is shockless accretion at very low mass flow densities, $\\dot m < 10^{-3}$\\,\\gcs. The infalling particles do not directly heat the lower atmosphere but create a hot corona and the atmosphere below is again heated radiatively \\citep{thomsoncawthorne87, woelkbeuermann92}. The task is, therefore, to consider the properties of atmospheres heated by a large external flux of ultraviolet/optical/infrared radiation. Modeling the UV-emitting spots in polars requires us to consider incoming radiation fluxes that exceed the flux emerging from the unheated atmosphere of the white dwarf by factors up to $\\sim200$ in the high state. Irradiated stellar atmospheres have been studied by several authors for a variety of conditions which differ, however, from those considered here \\citep[e.g.,][]{londonmccray81, vanteeselingetal94, brettsmith93, barmanetal04}. In this paper, we present a pilot study of the structure and the angle-dependent spectra of white dwarfs irradiated by an intense source of infrared to ultraviolet radiation.  The paper is arranged as follows. In Sect.\\,2 we describe the stellar atmosphere code used to calculate the irradiated atmosphere models, Sect.\\,3 introduces the essential parameters of the accretion regions in polars, in Sect.\\,4 we report the results for simple irradiation geometries, which are applied to the prototype polar AM~Her in Sect.\\,5. In Sect.\\,6, we discuss the limitations of the approach and lines for future research. ",
        "conclusions": "We have quantitatively reproduced the $1150-1450$\\,\\AA\\ continuum spectra of AM~Her in its high state of accretion by a model that considers a large polar cap heated by irradiation from the accretion region.  Irradiation strongly affects the temperature stratification of the white dwarf atmosphere in the heated polar cap. Strong heating leads to an inversion of the radial temperature profile in the atmosphere, causes the Lyman absorption lines to fill up completely, and produces substantial emission in the Lyman continuum. The known irradiation sources can account for the observed reprocessed flux longward of the Lyman limit, but an additonal source is needed to account for the predicted Lyman continuum flux. The implied luminosity of the polar cap is larger than considered previously \\citep{gaensickeetal95-1} and stresses the point that the reprocessed flux emerges primarily in the FUV and not in the soft X-ray regime as predicted by the early simple models of polars \\citep{kinglasota79, lambmasters79}.  Since X-ray heating \\citep{vanteeselingetal94} has not been considered here it is appropriate to discuss the potential contribution from that source in some detail. The horizontal temperature profile of the spot vs. distance from the spot center and the vertical atmospheric temperature profile at a given position may both be expected to differ for irradiation with either optical/infrared radiation or hard X-ray bremsstrahlung. Cyclotron radiation will produce a larger spot than bremsstrahlung because it originates higher up in the post-shock region and its beaming properties allow it to reach regions further away from the spot center. This difference is strenghtened by the lower shock heights of the denser, bremsstrahlung-dominated shocks (Eq.~14) and by the fact that bremsstrahlung is emitted isotropically and irradiation is thereby concentrated closer to the spot center.  For a given surface element, one can judge the efficiency of heating by cyclotron radiation and by bremsstrahlung as follows. As seen from Fig.~3 (right panel), cyclotron heating is most intense at Rosseland optical depths $\\tau_{\\rm ross} \\la 0.01$ but reaches down to the photosphere and beyond. Hence, cyclotron radiation is completely reprocessed fairly high in the atmosphere. Estimating the efficiency of bremsstrahlung heating requires us to consider the competition between photoabsorption and Compton scattering. Hydrogen is 99\\% ionized at the relevant levels of our atmospheres, which still allows atomic transitions to determine $\\tau_{\\rm ross}$ and causes the electron scattering optical depth to become $\\tau_{\\rm es} \\ll \\tau_{\\rm ross}$. For the typical 10\\,keV bremsstrahlung photons, on the other hand, photoabsorption is negligible with $\\tau_{\\rm ph} \\ll \\tau_{\\rm es}$. Hence, bremsstrahlung photons will penetrate far below the photosphere and are preferentially (back)scattered rather than photoabsorbed with the consequence of ineffective atmospheric heating \\citep{vanteeselingetal94}. Heating by bremsstrahlung may be important, however, if a finite metalicity of the atmosphere drastically increases the absorption cross section. Since the metalicity of the atmosphere will be enhanced in the inner spot where accretion occurs, X-ray heating may dominate there, while cyclotron heating is expected to dominate further away from the spot center. While a more detailed model should take both irradiation sources into account, we presently lack information on parameters like the mass inflow rate per unit area and the ensuing metalicity as functions of distance from the spot center.  The heating of the extended polar cap from the post-shock source tends to zero for vanishing shock height because then all downward directed emission is intercepted by the atmosphere within the proper accretion spot.  Disregarding irradiation by emission from the pre-shock stream for the moment, our model predicts that high-field polars, in which all cyclotron and X-ray emission occurs in low-lying shocks (see Eq.~14), should lack strongly irradiation-heated polar caps. This is, in fact, the case for the 200\\,MG system AR~UMa in its low state with some remnant cyclotron emission and an amplitude of the orbital modulation at 1400\\AA\\ of only 5\\% \\citep{gaensickeetal01-1}.  An investigation of the temperature distributions of a larger sample of white dwarfs in polars would shed light on some poorly understood aspects of accretion physics. Such a study would have to consider cooling of the deeply heated polar cap as well as the relevent heating processes. Note that accreting white dwarfs are generally hotter than single white dwarfs of the same age \\citep{araujobetancoretal05-1} as a result of compressional heating \\citep{townsleybildsten03}, but this process is rooted in the entire non-degenerate envelope (and to some extent in the core) and heats the entire star and not just the polar cap where accretion takes place.  Finally, a source of uncertainty in the calculated spectra warrants mentioning. Our atmospheric model disregards the compression of the outer atmosphere by accretion and is applicable strictly only outside the region of the extended source, assumed to be fed by accretion. The bombardment of the atmosphere by charged particles adds an external pressure and affects the structure of the upper atmosphere where the emission lines originate. The physics of bombarded atmospheres is complicated and has been treated to a certain extent by \\citet{woelkbeuermann92}. A simple estimate suggests that the effect may not be negligible. For a mass flow rate $\\dot m = 10^{-4}$\\,g\\,cm$^{-2}$s$^{-1}$, which is below present detection techniques, the ram pressure $P_{\\rm ram} = \\dot m \\upsilon_{\\rm ff}$, with $\\upsilon_{\\rm ff}$ the free-fall velocity, amounts to a few percent of the photospheric pressure and severely compresses the atmosphere only outside $\\tau_{\\rm ross} \\simeq 3\\times 10^{-4}$. A mass flow rate ten times higher, however, reaches down to $\\tau_{\\rm ross} \\simeq 0.05$. Under these circumstances, a more detailed study of the line profiles and the strength of the Lyman continuum in irradiated atmospheres of mCVs, including NLTE effects, requires a substantially larger effort."
    },
    "0512/astro-ph0512085_arXiv.txt": {
        "abstract": "{We have shown earlier that, contrary to popular belief, Einstein--de Sitter (E--deS) models can still fit the {\\sl WMAP} data on the cosmic microwave background provided one adopts a low Hubble constant and relaxes the usual assumption that the primordial density perturbation is scale-free. The recent {\\sl SDSS} measurement of the large-scale correlation function of luminous red galaxies at $z \\sim 0.35$ has however provided a new constraint by detecting a `baryon acoustic peak'. Our best-fit E--deS models do possess a baryonic feature at a similar physical scale as the best-fit $\\Lambda$CDM concordance model, but do not fit the new observations as well as the latter. In particular the shape of the correlation function in the range $\\sim 10-100 h^{-1}$ Mpc cannot be reproduced properly without violating the CMB angular power spectrum in the multipole range $l \\sim 100-1000$. Thus, the combination of the CMB fluctuations and the shape of the correlation function up to $\\sim 100 h^{-1}$Mpc, if confirmed, does seem to require dark energy for a homogeneous cosmological model based on (adiabatic) inflationary perturbations.} ",
        "introduction": "The detection by {\\sl COBE} of large angular scale fluctuations in the cosmic microwave background (CMB) opened a new era in modern cosmology and established the inflationary model, in which the initial conditions of the present Universe are set by fundamental physical processes occurring in the very early Universe. The subsequent detection of fluctuations on small angular scales and the considerable improvement in the precision of their measurement in the last decade has led to remarkable progress. The flatness of the universe was the first key result to emerge from identification of the `acoustic peak' in the angular power spectrum of the CMB (Lineweaver et al, 1997) which allowed a measurement of the angular distance to the last scattering surface.  Together with the observed low matter density of the universe and the cosmic acceleration implied by the Hubble diagram of Type Ia supernovae (Riess et al, 1998; Perlmutter et al. 1999), this led to the `concordance model' of cosmology in which the dominant component of the universe is a mysterious dark energy which behaves essentially like a cosmological constant (see e.g. Peebles and Ratra 2003).  This model has been remarkably successful at fitting a large body of cosmological data, in particular of large-scale structure (LSS). However it should be emphasized that {\\em direct} evidence for dark energy, through the detection of the expected correlations between the CMB and LSS induced by the late integrated Sachs--Wolfe effect, still has rather weak ($<3\\sigma$) significance (see e.g. Boughn \\& Crittenden 2004, Padmanabhan et al. 2005). Dimming of distant supernovae by grey dust (Aguirre 1999, Goobar et al. 2002, Vishwakarma 2005) or evolution of their progenitors (Drell et al. 2000, Wright 2002) could well mimic the effects of acceleration in the SN~Ia Hubble diagram (see Figure 7 in Riess et al. 2004). Moreover there are many `degeneracies' in the fitting of cosmological models to CMB and LSS observations which allow rather different parameter combinations to fit the same data. Given the complete lack of theoretical understanding of dark energy, this motivated us to reexamine (Blanchard et al. 2003, BDRS hereafter), whether Einstein--de Sitter (E--deS) models could still be compatible with the extant CMB and LSS data. This required assumptions at odds with current beliefs: we adopted a low value of the Hubble constant (46 km/s/Mpc) and, motivated by theoretical ideas about inflation (Adams, Ross \\& Sarkar 1997, Martin \\& Brandenberger 2003),\\footnote{The `glitches' appearing in WMAP and Archeops data may well be the signature of such new physics (Ringeval \\& Martin 2004, Hunt \\& Sarkar 2004).} we relaxed the assumption that the spectrum of the primordial density perturbation is a simple power-law. The fact that data from WMAP as well as other CMB experiments can be reproduced with E--deS models was a direct demonstration that the CMB data by themselves do {\\em not} require the presence of dark energy.  Moreover we showed that the power spectrum of large-scale structure as measured in the {\\sl 2dF} galaxy redshift survey, and inferred from observations of the Lyman-$\\alpha$ forest, can be adequately reproduced if there is a small component (12\\%) of unclustered matter e.g. in the form of 0.8 eV mass neutrinos or a pressureless scalar field. This lowers $\\sigma_8$, the amplitude of matter fluctuations on the scale $8h^{-1}$ Mpc, to 0.64 and 0.5 respectively, thus allowing agreement with the value{\\bf s} inferred from clusters and from weak lensing for a critical matter dominated universe. In both models, the baryon density is $\\Omega_{\\rm b}h^2 \\simeq 0.02${\\bf ,} in agreement with the value inferred from primordial nucleosynthesis  (see Fields \\& Sarkar 2004 and Charbonnel \\& Primas 2005 for recent discussions). Thus it appeared that an E--deS universe could indeed be made consistent with all data, and possibly favored by the evolution of the number density of distant clusters as inferred from {\\sl XMM} observations (Vauclair et al 2003) as well as the baryon fraction in clusters (Sadat et al. 2005).  It had been emphasized by BDRS that further observations of large-scale structure would provide a key test: {\\it ``The most stable difference between our E-deS models and the $\\Lambda$CDM concordance model is in fact the matter power spectrum shape in the range $k \\sim (0.01-0.03)~h$/Mpc, which galaxy surveys may be able to investigate, provided the possible biasing is reliably understood on these scales.''}  The recent results in this context from {\\sl 2dFGRS} (Cole et al. 2005), as well as {\\sl SDSS} (Eisenstein et al. 2005) have prompted us to return to this question. ",
        "conclusions": "Our main conclusion is that the new information contained in the {\\sl SDSS} LRG correlation function, in conjunction with observations of CMB anisotropy, in principle allows discrimination between cosmological models with and without dark energy. Although we have shown earlier that Einstein--de Sitter models can indeed fit most cosmological data (with the exception of the SNIa Hubble diagram and the Hubble Key Project value of the Hubble constant), it now appears possible to exclude them by more precise measurements of the correlation function of galaxies on large scales, provided that biasing is well understood. The concordance model does not have any basis yet in fundamental physics and should therefore be regarded as a convenient parameterization of the data in the context of the standard FRW cosmology, rather than as a `Standard Model'. In particular there is no physical explanation of why the universe should be embarking an a new inflationary period at this late stage in its history. Nevertheless we acknowledge that the new observations of the galaxy correlation function, jointly with small-angle anisotropies in the CMB which probe the same scales of $\\sim 10-100 h^{-1}$ Mpc, provide a remarkable geometric test of the concordance model which it passes successfully."
    },
    "0512/hep-ph0512323_arXiv.txt": {
        "abstract": "\\begin{center} {\\bf{Abstract }}\\\\ Axions are pseudo-scalar particles, those arise because of breaking of Peccei Queen (PQ) symmetry. Axions have a tree level coupling to two photons. As a consequence there exists a tree level coupling of axion to photon in a magnetic field. However, in an external magnetic field, there exists a new loop induced, axion  photon vertex, that gives rise to axion photon coupling. The strength of the tree level axion photon coupling in magnetic field is known to be model dependent. However in a magnetic field, the new loop induced coupling has some interesting features. This note discusses the new axion photon vertex in a magnetized medium and the corrections arising from there. The magnitude of the correction to axion photon coupling, because of magnetized vacuum and  matter is estimated in this note. While making this estimate we note that the form of the axion photon vertex is related to the axial polarization tensor. This vertex  is shown to satisfy the Ward identity. The coupling is shown to have a momentum dependent piece in it.  Astrophysical importance of this extra modification is also pointed out. \\end{center} \\noindent ",
        "introduction": "\\label{sec:intro} \\setcounter{equation}{0} \\setcounter{footnote}{0}  \\noindent Axions play an important role in the conceptual aspects of particle physics today. They are believed to be associated with  spontaneous breaking of global Chiral symmetry $U(1)_{PQ}$ (Peccei Queen symmetry), postulated to provide an elegant solution to strong CP problem ~\\cite{Peccei77,WW}. They are ultralight pseudo-scalar  field \\cite{Peccei96}. In the Weinberg-Wilczek-Peccei-Queen model ( original ), the symmetry breaking scale was assumed to be around weak scale, $f_w$. Although the original model, associated with the spontaneous breakdown of the global PQ symmetry at the Electro Weak scale (EW) $f_w$, is excluded experimentally, modified versions of the same with their associated axions are still of interest; where the symmetry breaking scale is assumed to lie between the EW scale and $~10^{12}$ GeV. Since the breaking scale of the PQ symmetry, $f_a$, is much larger than the electroweak scale $f_a \\gg f_w$, the resulting axion turns out to be  very weakly interacting (coupling constant $\\sim f_a^{-1}$), very light ($m_a \\sim f_a^{-1}$) and is often called ``the invisible axion model''.   \\noindent Till date, it remains elusive to experimental confirmation, however  there have been some  efforts to constrain  its parameters  through various cosmological or astrophysical considerations. For instance, cosmological observational constraints put  bounds on its mass (such that the universe is not over closed). Through such arguments, the allowed range for the axion mass $m_a$ has turned out to be ~\\cite{Turner,Raffelt90,Raffelt-castle97,Raffelt-school97}, \\begin{equation} 10^{-5} \\, \\mbox{eV} \\lesssim m_a \\lesssim  10^{-2} \\, \\mbox{eV}. \\label{eq:MsAx} \\end{equation} \\noindent Apart from the one mentioned above, there are astrophysical considerations too that constrain axion coupling to photons or fermions. For instance, if they exist,  being very weakly interacting particle, axions can drain away energy from stellar interiors. Since they are produced through processes like, $~~e^+ + e^- \\to \\gamma + a~~$ or the cross channel reaction, $~~e^-  + \\gamma \\to e^- + a ~~$ or $~~\\gamma_{plasmon} \\to \\gamma + a~~$, $\\gamma + \\gamma \\to a$ etc.; for a given coupling constant and mass,  one can estimate the rate at which the axions would draw away energy form the stars. The bounds are placed by fact that, for a given energy budget of a star, the amount of energy drained by axion emission should be less than it's observed luminosity. Apart from the ones mentioned above, there also have been experimental search for solar axions, through the conversion of an axion into a photon in a cavity, in presence of an external magnetic field. These searches also have placed some bound on the axion photon coupling. Incidentally its worth  noting that since most of the astrophysical objects are associated with magnetic field the same process (or the reverse of it) can take place even in astrophysical environments too.\\\\   \\noindent The coupling of axion to photon is realized through a term in the Lagrangian of the following form, \\begin{equation} {\\cal L} = \\frac{1}{M \\,}\\, a \\,\\mathbf{E}\\cdot\\mathbf{\\cal B}. \\label{axlagrangian} \\end{equation} Where $a$ is the axion and $M$ is the axion coupling mass scale. The experimental bound on $M$ coming from the study of  solar axions is set to be, $M > 1.7 \\times 10^{9 - 11}GeV$~\\cite{Moriyama}. It may be worth pointing out here that, though it is usually believed that, $M > 1.7 \\times 10^{9 -11}GeV$, but this is a model dependent number. It should be noted here that the PQ symmetry breaking scale parameter $f_a$ is proportional to $M$. A detailed survey of various astrophysical bounds on  the parameters of axion models and constraints on them, can be found in Ref.~\\cite{Raffelt-book}.\\\\    \\noindent Since most of the bounds on axion parameters arise from astrophysical and cosmological studies where medium  and a magnetic field are present, it becomes important to seek the modification of   the axion coupling to photon, in presence of a medium or magnetic field or both. Particularly in some astrophysical situations where the magnetic component, along with medium (usually referred as magnetized medium) dominates. Examples being, the Active Galactic Nuclei (AGN), Quasars,  Supernova, the Coalescing Neutron Stars or  Nascent Neutron Stars, etc., to name a few. Apart from the ones discussed before, lately there seem to be some observational signature for possible existence of astrophysical objects,( called Magnetars ), with magnetic field strength, ${\\cal B} \\sim 10^{15} - 10^{17}$~G, i.e.,  significantly above the critical, Schwinger value ${\\cal B}_e = m^2_e / e \\simeq 4.41 \\times 10^{13}$~G ~\\cite{toroidal,poloidal}. However we would like to emphasize that even normal astrophysical objects are always associated with magnetic field, though the strength of the field may not be as strong as ${\\cal B}_e$. Thus justifying the role of magnetized medium on axion properties. Also recently there has been an attempt to describe the observed faintness of Type Ia Supernova remnants, based on axion photon oscillation in the magnetic field of the intergalactic medium.\\\\  \\noindent In view of these interesting physical applications of axion physics in astrophysics as well as cosmology it seems timely to find out the effect of medium and magnetic field, on the couplings of axions to photons. \\\\   \\noindent In this note we would investigate the matter induced photon axion coupling in a magnetized medium. Where the particle in the plasma will be considered to be mostly electrons, though it could be any fermion. Since the temperature in these astrophysical objects are not too large (of the order of hundred MeV or so at the most) this seems to be a reasonable approximation. Of course our formulas are general enough to be extended to any temperature and density.\\\\   \\indent The organization of this document is as follows, in section II we would discuss  about the physics of axion photon coupling and the model dependent uncertainties that enter in the axion photon coupling parameter $M$ given in eqn. [\\ref{axlagrangian}]. In the same section we would also try to give a brief over view of the existing axion models and their type of coupling with fermions. In the next section (i.e. section II), following Schwinger's approach \\cite{schw}, we would elaborate on the details of the magnetized propagators. As would be discussed later, matter induced Axion photon coupling in a magnetized medium has two contributions in it, one coming from the magnetized vacuum and the other from the magnetized medium. Sections III and IV  would deal with the details of those contributions. \\ Finally at the end we would conclude by justifying our results through a general analysis and  order of magnitude estimation of the modifications one is getting from the presence of magnetized plasma. Finally we  would like to conclude by pointing out the possible applications of our result. ",
        "conclusions": "As has already been mentioned at the beginning, the purpose of this note is to find out the modifications in the axion photon vertex in the presence of an external magnetic field and matter. Field induced part of the photon axion vertex has been worked out earlier in \\cite{raff-m}, however the effect of magnetized matter wasn't taken into account. In this note we have considered the effect of both external magnetic field and medium on the vertex. More over we have also written down the explicit formula for the vertex involving the PQ charges of the fermions. In the light of these estimates,  it is possible to write down the axion photon mixing Lagrangian, for low frequency photons in an external magnetic field, in the following way: \\begin{eqnarray} {\\cal L}^{Total}_{a \\gamma} = {\\cal L}^{vac}_{a \\gamma}+ {\\cal L}^{{\\cal B}}_{a \\gamma}+ {\\cal L}^{{\\cal B},\\mu,\\beta}_{a \\gamma}. \\end{eqnarray} Where ${\\cal L}^{vac}_{a \\gamma}$ is the usual axion photon mixing term in vacuum, resulting form the electromagnetic anomaly, and is given by \\cite{kim}, \\begin{eqnarray} \\mbox{\\hskip 3.25 cm} {\\cal L}^{vac}_{a \\gamma}&=& - g_{a \\gamma \\gamma} \\frac{e^2}{32\\pi^2} a \\mbox{F}{\\tilde{\\mbox{F}}}, \\mbox{~~where~} g_{a \\gamma \\gamma}= \\frac{A^{c}_{PQ}}{f_a} \\left[\\frac{A^{em}_{PQ}}{A^{c}_{PQ}}- \\frac{2(4+z)}{3(1+z)}\\right] \\!\\!\\!\\!\\!\\!\\! \\nonumber \\\\ {\\cal L}^{\\cal B}_{a\\gamma}& =& \\frac{-1}{32 \\pi^2 } \\left[ 4 + \\frac{4}{3} \\left( \\frac{k_{\\para}}{m} \\right)^2 \\right] \\!\\!\\!\\sum_f g_{ af} (eQ_f)^2 a \\mbox{F}\\tilde{\\mbox{F}}. \\nonumber \\\\ {\\cal L}^{{\\cal B},\\mu,\\beta}_{\\gamma a}&=&\\frac{16}{32\\pi^2} \\cdot \\left(\\frac{k_{\\para}}{\\omega}\\right)^2 ({\\widetilde\\Lambda}) \\sum_f g_{af}(eQ_f)^2 a\\mbox{F}\\tilde{\\mbox{F}}. \\label{axpefl} \\end{eqnarray} To remind ourselves, in the equations above, $z=\\frac{m_u}{m_d}$, where $m_u$ and $m_d$ stands for the masses of the light quarks. The anomaly factors are given by the following relations, $A^{em}_{PQ} = \\mbox{Tr}(Q^2_f)X_{f}$ and $\\delta_{ab}A^{em}_{c}= \\mbox{Tr}\\lambda_a\\lambda_b X_{f}$ (where the trace is over the fermion species.). Hence, in the limit of $|k_{\\perp}| \\to 0 $ and $\\omega << m_f $, using eqn. [\\ref{axpefl}], one can write the total axion photon effective Lagrangian:  \\begin{eqnarray} {\\cal L}^{Total}_{a \\gamma}&=& - \\left[g_{a \\gamma \\gamma} + \\left( 4 + \\frac{4}{3}\\left(\\frac{k_{\\para}}{m} \\right)^2 \\right) \\sum_f g_{af}(Q_f)^2 -16 \\left(\\frac{k_{\\para} }{\\omega}\\right)^2 ({\\widetilde\\Lambda}) \\sum_f g_{af}(Q_f)^2 \\right] \\frac{e^2}{32\\pi^2} a\\mbox{F}\\tilde{\\mbox{F}}. \\label{axpefl.b} \\end{eqnarray} In order to describe the axion photon interaction, for static magnetic field and real photon, the effective Lagrangian employed in the literature is usually given by the following relation, \\begin{eqnarray} \\frac{1}{M}a \\vec{\\mbox{E}}.\\vec{\\cal{B}}^{ext}. \\label{usualeffectivelag} \\end{eqnarray} Where the parameter $M$, defines an energy scale in terms of inverse of the axion photon photon coupling constant in vacuum, i.e $ M \\equiv \\frac{32\\pi^2}{e^2 g_{a \\gamma \\gamma}}$. In principle this factor is a model dependent quantity, and depending on the particular choice of PQ charges, $g_{a\\gamma \\gamma}$ can vary between zero to hundred. This is one of the model dependent uncertainty. On the other hand, in some experiments, attempts have been made to give a bound on $M$ thus $g_{a\\gamma \\gamma}$ through the detection of solar axions in a magnetic cavity haloscope. The astrophysical bounds on $M$ come from the constraint that the stars don't lose energy through axion photon conversion at a rate faster than the generation of energy. As a result of these investigations, the present day bound on the energy scale is believed to be, $M \\le 10^{12} GeV$.  On the other hand in this note, we have been able to point out that apart from the model dependent uncertainties, there is also medium dependent uncertainties that can introduce some variation on the bounds on $M$. The amount of this  variation would depend on the kind of environment one is interested in. The details of its implications for various models and physical situations are extremely interesting by their own right \\cite{massimo,cameron,semertzidis,soa,long,stod,wilc,miani,gas,par}; however these points are beyond the scope of this work. \\begin{figure} \\begin{center} \\input{ag22.tex} \\end{center} \\caption[]{$\\widetilde{\\Lambda}$ vs. $\\frac{\\mu}{T}$ for  relativistic degenerate medium. The parameters are as follows: $m \\le \\mu \\le 6$ when the temperature is held fixed at, $T=10^{-3}$ in units of fermion mass. } \\label{Fig:f2} \\end{figure}"
    },
    "0512/astro-ph0512200_arXiv.txt": {
        "abstract": "We present a generalisation of surface photometry to the higher-order moments of the line-of-sight velocity distribution of galaxies observed with integral-field spectrographs. The generalisation follows the approach of surface photometry by determining the best fitting ellipses along which the profiles of the moments can be extracted and analysed by means of harmonic expansion. The assumption for the odd moments (e.g. mean velocity) is that the profile along an ellipse satisfies a simple cosine law. The assumption for the even moments (e.g velocity dispersion) is that the profile is constant, as it is used in surface photometry. We test the method on a number of model maps and discuss the meaning of the resulting harmonic terms. We apply the method to the kinematic moments of an axisymmetric model elliptical galaxy and probe the influence of noise on the harmonic terms. We also apply the method to {\\tt SAURON} observations of NGC~2549, NGC~2974, NGC~4459 and NGC~4473 where we detect multiple co- and counter-rotating (NGC~2549 and NGC~4473 respectively) components. We find that velocity profiles extracted along ellipses of early-type galaxies are well represented by the simple cosine law (with 2\\% accuracy), while possible deviations are carried in the fifth harmonic term which is sensitive to the existence of multiple kinematic components, and has some analogy to the shape parameter of photometry. We compare the properties of the kinematic and photometric ellipses and find that they are often very similar, but a study on a larger sample is necessary.  Finally, we offer a characterisation of the main velocity structures based only on the kinemetric parameters which can be used to quantify the features in velocity maps. ",
        "introduction": "\\label{s:intro}  Over the last three decades broad band observations of early-type galaxies were successfully analysed by a method commonly called surface photometry or, simply, photometry. This method is based on the analysis of isophotal shapes of the projected surface brightness. The development of the method was stimulated by the empirical discovery that the isophotes of early-type galaxies are reproduced by ellipses to better than 1 per cent \\citep{1984ApJS...56..105K, 1985ApJS...57..473L, 1985AJ.....90..169D, 1987MNRAS.226..747J, 1990AJ....100.1091P}.  Although the isophotes are elliptical in shape to high accuracy, under careful examination, many isophotes of early-type galaxies do show differences from pure ellipses at a level of $\\approx0.5\\%$ \\citep[e.g.][]{1988A&AS...74..385B, 1990AJ....100.1091P}. The true success of photometry was the ability to measure these deviations and classify early-type galaxies accordingly into disky and boxy objects \\citep{1985MNRAS.216..429L,1987MNRAS.226..747J,1987A&A...177...71B}. When combined with information on total luminosity and spatially resolved spectroscopy, it followed that the duality of photometric properties of early-type galaxies is reflected in a duality of kinematic properties, where faint disky objects were found to rotate faster than luminous boxy objects \\citep{1983ApJ...266...41D, 1983ApJ...266..516D, 1988A&A...193L...7B, 1988A&A...195L...1N, 1988A&A...195L...5W, 1989A&A...217...35B, 1989A&A...215..266N, 1992A&A...262...52B, 1994MNRAS.269..785B}. High resolution imaging studies with the Hubble Space Telescope \\citep{1994AJ....108.1567J, 1995AJ....110.2622L, 2001AJ....121.2431R}, again based on photometric analysis, revealed new properties of early-type galaxies that deepened the division of the galaxies in two groups. These remarkable set of discoveries resulted in a revised classification scheme of galaxies \\citep{1996ApJ...464L.119K}.  It is fair to say that our increased knowledge of early-type galaxies (partially) comes from photometry and the ability of the method to harvest and describe in a compact way the information from two-dimensional images. However, the integrated light remains a limited source of information about the internal structure and thus the true nature of galaxies. Kinematic information is also essential to fathom the complexity of galaxies and, especially, two-dimensional kinematic information is required \\citep[e.g.][]{1991ApJ...383..112F, 1991AJ....102..882S, 1994ApJ...425..481S, 1994ApJ...425..458S, 1994ApJ...425..500S, 1994MNRAS.271..924A}.  Two-dimensional velocity maps were, until recently, only possible for objects with clear emission lines. Such maps were the result of studies of, e.g., H{\\small I} with radio interferometers, \\citep[e.g.][]{2002A&A...390..829S}, CO with millimetre interferometers \\citep[e.g.][]{2003ApJS..145..259H} and H$\\alpha$ with Fabry-Perot spectrographs \\citep[e.g.][]{2005MNRAS.360.1201H}.  The advent of integral-field spectrographs (e.g. {\\tt TIGER}, {\\tt OASIS}, {\\tt SAURON}, {\\tt PMAS}, {\\tt GMOS}, {\\tt SINFONI}, {\\tt OSIRIS}), however, has brought two-dimensional kinematic measurements to classical optical and near-infrared wavelengths. Here it is possible to probe both the stellar absorption- and gas emission-lines, which may co-exist in the same potential with very different spatial distributions and dynamical structures. The wealth of features seen in stellar kinematic maps of early-type galaxies \\citep{2004MNRAS.352..721E} confirms the usefulness of two-dimensional data, but also poses a problem to efficiently harvest and interpret the important features from the maps.  An approach using harmonic expansion was developed for the analysis of two-dimensional velocity maps of disk galaxies. This method divides a velocity map into individual rings \\citep[the so-called tilted-ring method,][]{1987PhDT.......199B} and performs a harmonic expansion along these rings \\citep[e.g.][]{1978MNRAS.183..779B, 1991wdir.conf...40T,1994ApJ...436..642F, 1997MNRAS.292..349S,2004ApJ...605..183W}. This analysis, however, is based on the assumption that the emission-line emitting material is confined to a thin disk structure.  Clearly, a spheroidal distribution of stars typical of early-type galaxies does not have the same dynamical properties as a gas disk. To explore those intrinsic properties, one needs a more general method which will not be based on assumptions about the nature of the observed system (e.g. a disk), but will rely solely on the properties of the investigated observables (e.g. surface brightness or velocity).  Here we explore such a general method.  Photometry earned its spurs in intensive applications over the last three decades and offers a natural starting point for the analysis of two-dimensional kinematic maps, although one can, in principle, invent a number of basis functions which can describe a two-dimensional distribution. In this paper, we present a generalisation of photometry to the higher-order moments of the line-of-sight velocity distribution (LOSVD). This generalisation is based on the theoretical fact that surface brightness is itself a moment of the LOSVD, and on our empirical discovery that the velocity maps of many early-type galaxies can be well reproduced by a simple cosine law along sampling ellipses.  We call our method {\\it kinemetry}\\footnote{This name was introduced by \\citet{2001sf2a.conf..289C} who presented a preliminary discussion on this topic.}, which reduces to photometry for the investigation of surface brightness distributions.  In Section~\\ref{s:back} we present the theoretical background and motivation for the new method. Section~\\ref{s:method} presents the technical aspects of the method. The meaning of the kinemetric coefficients and their diagnostic merits for different model maps are presented in Section~\\ref{s:kinpar}. In Section~\\ref{s:appli} we apply the method on kinematic maps of an axisymmetric model galaxy, present the application of kinemetry to actual observations and characterise typical structures on velocity maps.  We summarise our conclusions in Section~\\ref{s:sum}. ",
        "conclusions": "\\label{s:sum}  We have presented a generalisation of surface photometry to the higher-order moments of the LOSVD, observed with integral-field spectrograph. We call our method \\emph{kinemetry}. For even moments of the LOSVD, kinemetry reduces to photometry. For odd moments, kinemetry is based on the assumption that it is possible to define an ellipse such that the kinematic profile extracted along the ellipse can be well described by a simple cosine law. This assumption is satisfied in the case of simple axisymmetric rotators.  Kinematic profiles that deviate from axisymmetry or contain multiple kinematic components will be more complex, and the residuals can be measured through a harmonic analysis of the profiles. We also find that the assumption is well satisfied for a number of observed galaxies.  Since the profiles are extracted along the best fitting ellipse, the number of harmonic parameters necessary to describe the profile is usually small.  A total of six harmonic terms is generally necessary to determine the best fitting ellipse ($A_0$, $A_1$, $A_2$, $B_2$, $A_3$ and $B_3$). The deviations from the simple cosine law are carried in higher harmonic terms. It is not expected that many more higher-order terms will be necessary in the general case, and we stop the analysis at the $A_5$ and $B_5$ harmonics.  The most important parameters of the analysis are the kinematic position angle $\\Gamma$, flattening $q$ of the ellipses, and the harmonic terms: $k_1$ and $k_5$, where we combined the terms of the same order.  We briefly summarise their properties:   \\begin{itemize}  \\item[-] the kinematic position angle traces the maximum velocity and describes the orientation of the map.  \\item[-] the flattening is related to the opening angle of the iso-velocities and describes the different components of the maps, which can appear {\\it flat} or {\\it round}. In the special case of a thin disk it gives the inclination of the galaxy;  \\item[-] $k_1$ is the dominant harmonic coefficient and describes the amplitude of bulk motions. It is related to the circular velocity. Higher order coefficients, which we normalise to $k_1$, describe the deviations from simple circular motion;  \\item[-] $k_5$ is the first higher-order term which is not fitted and quantifies higher-order deviations from rotational motions. It indicates complexity on the maps and is sensitive to the existence of separate kinematic components. A photometric analog of this term is the coefficient that describes the diskiness/boxiness of the isophotes. \\end{itemize}  This method can be straight-forwardly applied to all moments of the LOSVD observed from different objects (e.g. early- and late-type galaxies). However, the analysis and the intepretation of the harmonic terms has to be related to the physical nature of the observed objects. An example are velocity maps of thin disks. In this case, the kinemetric flattening and position angle are directly related to the inclination and orientation of the disk, and the results of kinemetry are equal to results of the tilted-ring method. If $q$ and $\\Gamma$ are kept fixed, kinemetry reduces to the harmonic analysis of \\citet{1997MNRAS.292..349S} and \\citet{2004ApJ...605..183W}. In this case, since the velocity profiles are not extracted along the best fitting ellipses, all harmonic terms will carry information about departures from the simple circular motion.  We applied the method on one model and four observed galaxies. The model galaxy was a two-integral model of an observed axisymetric galaxy. We constructed maps of observable odd kinematic moments of the LOSVD: mean velocity, velocity dispersion and Gauss-Hermite coefficient $h_3$. We analysed both maps to show that the method works well on an axisymmetric galaxy (all higher-order terms were small or consistent with zero). We also introduced an intrinsic scatter and currently realistic measurement uncertainties to show the difficulties expected analysing the higher-order odd moments (e.g. $h_3$) obtained from state-of-the-art integral-field observations. The position angle, flattening and amplitude of the moment can be realistically recovered, but the signal-to-noise ratio is too low for expansion to higher-order harmonics.  The four galaxies were taken from the {\\tt SAURON} sample. We used them to show the descriptive power of the new method: detection of co- or counter-rotating subcomponents well hidden in the kinematics, which could be detected previously only through detailed modelling. We also showed that there are early-type galaxies with rotation velocity described by a simple cosine law. The kinemetric position angle and flattening of these galaxies are in a good agreement with the photometric position angle and flattening.  Based on the kinemetric parameters we characterise several features on the velocity maps: Disk-like Rotation, Multiple Components, Kinematic Twists and Low-level Rotation.  Other, more specific and instrument dependant features on the maps can be constructed from these basic groups.  We wish to stress that the method works well on velocity maps because many of them satisfy the basic assumption of the method to the level of a few per cent. This empirical fact is related to the internal structure of the galaxies. We developed the method of kinemetry to harvest the information from the observed maps of moments of the LOSVD probing the nature of galaxies.  An IDL implementation of the described algorithm is available on the web address: http://www-astro.physics.ox.ac.uk/$\\sim$dxk/.\\\\     \\noindent{\\bf Acknowledgements} We thank Eric Emsellem for providing a model axisymmetric galaxy and Maaike Damen for providing the photometric data used in this work. DK thanks Martin Bureau and Eric Emsellem for a careful reading of the manuscript and Marc Sarzi, Glenn van den Ven and Richard McDermid for fruitful discussions. This research was supported by NOVA, the Netherlands Research School for Astronomy and by PPARC grant PPA/G/O/2003/00020 'Observational Astrophysics in Oxford'. MC acknowledges support from a VENI grant 639.041.203 awarded by the Netherlands Organization for Scientific Research (NWO)."
    },
    "0512/astro-ph0512170_arXiv.txt": {
        "abstract": "The ability to accurately measure the shapes of faint objects in images taken with the {\\it Advanced Camera for Surveys} (ACS) on the {\\it Hubble Space Telescope} (HST) depends upon detailed knowledge of the Point Spread Function (PSF). We show that thermal fluctuations cause the PSF of the ACS Wide Field Camera (WFC) to vary over time. We describe a modified version of the \\tinytim\\ PSF modeling software to create artificial grids of stars across the ACS field of view at a range of telescope focus values. These models closely resemble the stars in real ACS images. Using $\\sim10$ bright stars in a real image, we have been able to measure HST's apparent focus at the time of the exposure. \\tinytim\\ can then be used to model the PSF at any position on the ACS field of view. This obviates the need for images of dense stellar fields at different focus values, or interpolation between the few observed stars. We show that residual differences between our \\tinytim\\ models and real data are likely due to the effects of Charge Transfer Efficiency (CTE) degradation. Furthermore, we discuss stochastic noise that is added to the shape of point sources when distortion is removed, and we present \\multidrizzle\\ parameters that are optimal for weak lensing science. Specifically, we find that reducing the \\multidrizzle\\ output pixel scale and choosing a Gaussian kernel significantly stabilizes the resulting PSF after image combination, while still eliminating cosmic rays/bad pixels, and correcting the large geometric distortion in the ACS. We discuss future plans, which include more detailed study of the effects of CTE degradation on object shapes and releasing our \\tinytim\\ models to the astronomical community. ",
        "introduction": "Accurate shape measurements of faint, small galaxies are crucial for certain applications, most notably the measurement of weak gravitational lensing. Quantifying the slight distortion of background galaxies by foreground matter relies on detecting small but coherent changes in the shapes of many galaxies (see Refregier 2003 for a recent review). To extract the lensing signal, it is crucial to remove instrumental effects from galaxies' measured shapes. On the {\\it Hubble Space Telescope} (HST), these include:  \\newpage  \\begin{itemize} \\item Convolution of an image with the telescope's Point Spread Function (PSF).  \\item Geometric distortion of an image. This is particularly large in the {\\it Advanced Camera for Surveys} (ACS) because of its location off HST's optical axis.  \\item Trailing of faint objects in the CCD readout direction due to degraded Charge Transfer Efficiency (CTE). \\end{itemize}  \\noindent In this proceeding, we describe a method to model and correct for the telescope's temporally and spatially varying PSF.  The geometric distortion has already been shown to be successfully removed during image processing by \\multidrizzle\\ (Koekemoer et al.\\ 2002). Removing the distortion does change the PSF, and we present recommendations to minimize stochastic changes introduced during the repixellization  stage of image processing. The effect  of continuing CTE degradation on galaxy shapes is only becoming apparent as the ACS spends longer in orbit, and is not yet completely understood. That is therefore beyond the scope of this proceeding. A separate method to remove CTE effects will be presented in Rhodes et al.\\ (2006), and the application of all these corrections in a weak lensing analysis will be presented in Massey et al.\\ (2006). Other branches of astronomy, including stellar photometry in crowded fields, the study of AGN, and proper motions also require detailed knowledge of the PSF and will benefit from the models we describe here.  In weak lensing, to deconvolve galaxy shapes from the PSF, we must accurately know the shape of the PSF at the position of the object and at the time of the observation For example, see Rhodes, Refregier \\& Groth (2000) for a description of the method we use on the Cosmic Evolution Survey (COSMOS; Scoville et al.\\ 2006) images we use to test the PSF models we describe in this paper. If the HST PSF were stable over time, it would be straightforward  to build a catalog of stellar images across the entire field of view. However, thermal fluctuations in HST that change its effective focus (the distance between the primary and secondary mirrors) lead to temporal PSF variations. As an example, Figure~\\ref{fig:cosmospsfs} shows the PSF pattern in two sets of COSMOS images. The left hand panel shows stars from images taken when the telescope was near optimal focus, and the right hand panel shows stars observed when the telescope was several microns below optimal focus. Each tick mark in the figure represents the ellipticity of one star, measured using the standard weak lensing definition,  \\begin{equation} \\label{ellipticity} |e|=\\frac{\\left[(I_{xx}-I_{yy})^2+(2I_{xy})^2\\right]^{1/2}}{I_{xx}+I_{yy}} ~, \\end{equation}  \\noindent where star's weighted second order moments  \\begin{equation} \\label{moments} I_{ij}=\\frac{\\sum w I x_i x_j }{\\sum w I} \\end{equation}  \\noindent involve summations over all pixels. $I$ is the intensity of a pixel, $w$ is some weighting function (in our case a Gaussian with a width of about the FWHM of the PSF), and $x_i$ is the distance of a pixel from the centroid of the object. It is apparent from Figure~\\ref{fig:cosmospsfs} that changes in the PSF over time are sufficient for a temporally stable ACS PSF model to be inadequate in demanding applications, when using data collected over a period of more than a few days. Other effects, including CTE, introduce additional variation on longer time scales.   \\begin{figure} \\plottwo{cosmosstars_timedepe1hi.ps}{cosmosstars_timedepe1lo.ps} \\caption{The ellipticity of stars in the COSMOS survey observed with the ACS WFC  while HST happened to be at nominal (left panel) and low (right panel) focus. The orientation and size of each tick mark represents the ellipticity of one star; both panels contain stars from several different fields. The difference in the PSF patterns is apparent, and demonstrates the need for a time dependent PSF model.} \\label{fig:cosmospsfs} \\end{figure}  This proceeding is organized as follows. In \\S\\ref{TinyTim} we introduce the \\tinytim\\ software package that we used for PSF modeling, and discuss modifications that we have made. In \\S\\ref{multidrizzle} we discuss \\multidrizzle\\ and how to minimize the aliasing of point sources that occurs during distortion removal. In \\S\\ref{models} we show  how we have used our \\tinytim\\ models to quantify the temporal variation of the ACS PSF, and describe how the same models can be used to correct for it. In \\S\\ref{conclusions} we draw conclusions and outline a plan to release our \\tinytim\\ PSF models. ",
        "conclusions": "\\label{conclusions}  We have shown that \\tinytim\\ can produce model PSFs that are very close to those observed in real data (for example the COSMOS 2-Square Degree Survey). This required some modifications to the \\tinytim\\ code, most importantly adding the ability to mimic the distortion correction and dithering normally implemented via \\multidrizzle, and to produce grids of PSFs across the entire ACS WFC field of view.  We used \\tinytim\\ model stars to find the best values for \\multidrizzle\\ and found that using a Gaussian kernel, \\texttt{pixfrac }$ =0.8$, and an output pixel scale of 0.03 arcseconds greatly reduced the ``aliasing'' of point sources introduced during repixellization.  Discrepancies between our models and the COSMOS data can be attributed largely to a degradation in the ACS CTE since launch. We are currently studying this problem and will present a complete PSF solution including how to correct for CTE in Rhodes et al.\\ (2006).  We plan to correct science images for CTE on a pixel-by-pixel basis, like the Bristow code developed for STIS (Bristow et al./ 2002), moving charge back to where it belongs, rather than including the effects of CTE in our model PSF (like Rhodes et al. 2004). Thus, the model PSFs we present here are the ones we plan to use in our weak lensing analysis.  At the time of press, we have thoroughly tested PSF models in only the F814W filter. However, our IDL routines preserve \\tinytim's ability to create PSFs in other filters, and for sources of any colors. Our routines are therefore easily adaptable to other data sets.  The whole method is intentionally designed to be as adaptable as possible for many methods. The desire to know the PSF at any arbitrary position on the sky is far from unique to weak lensing. But even in lensing, advanced methods like Shapelets will, in the near future, be able to take advantage of more detailed information about the PSF shape than it is reasonable to expect from interpolation between a few stars (Massey \\& Refregier 2005; Refregier \\& Bacon 2004). This is even more exaggerated when considering higher order shape parameters, with an intrinsically lower signal to noise. The creation of noise-free, oversampled stars at any position on an image allows such analysis in any ACS data.  In the near future, we plan to release our PSF models to the community along with the wrapper we have written for\\tinytim\\ which will allow users to create PSF models in different filters and at user-defined positions in the ACS field of view. Interested readers are advised to contact the authors for these resources."
    },
    "0512/astro-ph0512346_arXiv.txt": {
        "abstract": "We present line strengths in the bulges and inner disks of 38 galaxies in the local universe, including several galaxies whose bulges were previously identified as being disk-like in their colors or kinematics, to see if their spectral properties reveal evidence for secular evolution.  We find that red bulges of all Hubble types are similar to luminous ellipticals in their central stellar populations.  They have large luminosity-weighted ages, metallicities, and $\\alpha$/Fe ratios.  Blue bulges can be separated into a metal-poor class that is restricted to late-types with small velocity dispersion and a young, metal-rich class that includes all Hubble types and velocity dispersions.  Luminosity-weighted metallicities and $\\alpha/Fe$ ratios are sensitive to central velocity dispersion and maximum disk rotational velocity.  Red bulges and ellipticals follow the same scaling relations.  We see differences in some scaling relations between blue and red bulges and between bulges of barred and unbarred galaxies.  Most bulges have decreasing metallicity with increasing radius; galaxies with larger central metallicities have steeper gradients.  Where positive age gradients (with the central regions being younger) are present, they are invariably in barred galaxies.  The metallicities of bulges are correlated with those of their disks.  While this and the differences between barred and unbarred galaxies suggest that secular evolution cannot be ignored, our results are generally consistent with the hypothesis that mergers have been the dominant mechanism responsible for bulge formation. ",
        "introduction": "Bulges are important relics of the galaxy formation process.  An analysis of their structure, kinematics, dynamics, and stellar content can potentially reveal the physical mechanisms responsible for the formation and evolution of galaxies as well as the nature of the Hubble sequence.  Similarities between bulges and ellipticals have long been recognized but recent observations suggest that at least some bulges may be related to disks.  This has led to the suggestion that the large bulges of early-type spirals are more similar to ellipticals while late-type bulges are more disk-like \\citep{1997ARA&A..35..637W}.  As a consequence of these observations, formation scenarios have emerged for bulges that are either identical to those for ellipticals or involve the secular evolution of disks.  However, the degree to which formation mechanisms are homogeneous is still open to question.   Early models for elliptical formation involved the monolithic collapse of a primordial gas cloud \\citep{1974MNRAS.166..585L, 1984ApJ...286..403C, 1987A&A...173...23A}.  This model naturally explains several observed properties of ellipticals including the mass-metallicity relation and the presence of metallicity gradients but large-scale collapse is inconsistent with present day cold dark matter cosmology and with recent observations showing that massive ellipticals were not fully assembled since after z$=$1 \\citep{2004ApJ...608..752B, 2005astro.ph..6044F}. It is now widely believed that ellipticals formed hierarchically through mergers of smaller fragments \\citep{1993MNRAS.264..201K}.  Mergers are frequently caught in the act \\citep{1999ApJ...520L..95V, 2004A&A...428..837F}  and photometric and kinematic evidence for past mergers is abundant in ellipticals \\citep{2004MNRAS.352..721E, 2005astro.ph..6661V}.  The merger model has been extended to bulges due to the many observed similarities between bulges and ellipticals.  For example,\\cite{1997AJ....114.2366C} found that bulges were well-fit by the R$^{1/4}$ law used for ellipticals.  The fundamental plane relation of bulges is nearly the same as that of ellipticals, with late-types perhaps lying below early types \\citep{2002MNRAS.335..741F}.  In the secular evolution scenario, bulges are produced through radial and vertical transport of disk material as the result of instabilities and resonances (see Kormendy \\& Kennicutt 2004 for a review).  These models come in several flavors, most of which involve bars.  Simulations that do not include gas have found that bars can buckle, heating the inner disk and increasing its scale height to resemble a bulge. Hydrodynamical simulations have found that bars can transport gas towards the center, triggering intense star-formation \\citep{1993gabu.symp..387P, 1995A&A...301..649F, 1996ApJ...462..114N, 2000MNRAS.312..194N, 2004A&A...413..547I}.  The presence of neighbors may also drive this \\citep{2004AJ....127.1371K}.  Support for secular evolution comes from observed correlations between the scale lengths of bulges and their disks \\citep{1996ApJ...457L..73C, 2003ApJ...582..689M}.  Recent work has also shown that the light profiles of many bulges are closer to exponential than R$^{1/4}$ \\citep{2003ApJ...582L..79B, 2003ApJ...582..689M,2004MNRAS.355.1155D}.  Furthermore, the ratio of rotational to random motions in bulges is often typical of disks \\citep{1982ApJ...256..460K, 2004ARA&A..42..603K}.  Comparisons between the morphology and kinematics of observed galaxies with simulated ones have shown that boxy and peanut-shaped (b/p) bulges are bars viewed at high inclination \\citep{1999AJ....118..126B, 2003Ap&SS.284..753A, 2004AJ....127.3192C, 2005MNRAS.358.1477A}.  Athanassoula (2005) distinguishes between b/p bulges, which are formed through the buckling of the bar, and what she calls ``disky bulges\", which are smaller cold components that formed out of the gas driven inward by the bar.   Bulges that could have been formed through secular evolution are often referred to as ``pseudobulges\" to distinguish them from the ``classical\" bulges that may have formed through mergers.  Since pseudobulge signatures are generally found in later-typed spirals, Kormendy \\& Kennicutt (2004) suggest that early-type spirals (Sa's, Sab's, and some S0's) contain classical bulges while late-type spirals (Sb's, Sc's and some S0's) contain pseudobulges.  On the theoretical side, \\cite{1993gabu.symp..387P} found that secular evolution can produce small bulges but not those having a characteristic radius much larger than the disk scale length.  However, it is not at all clear that the spectrum of observed bulge properties points towards two distinct formation scenarios.  Since the stability of bars continues to be debated \\citep{2004ApJ...604..614S, 2004ApJ...604L..93D, 2005MNRAS.tmpL..89B}, it is also not clear whether or not pseudobulges should exist only in present-day barred galaxies.   Stellar population (SP) studies can potentially place important constraints on the formation mechanisms.  A succesful formation scenario has to reproduce the observed distribution of ages and metallicities.  In a collapse model, bulges and ellipticals are universally old and have radial metallicity gradients.  In his dissipative collapse simulation \\cite{1984ApJ...286..403C} found that the steepness of the metallicity gradient was correlated with galaxy properties such as mass and luminosity.  If ellipticals and bulges formed through mergers, it is important to keep in mind that their assembly histories might be very different from their star formation histories.  $\\Lambda$CDM simulations suggest that most massive ellipticals (and therefore presumably bulges) were not fully assembled until recently (z$<$1) whereas the bulk of star formation occurred much earlier (z$>$2) in the progenitor galaxies \\citep{2005astro.ph..9725D}.  This is consistent with observational studies of merger activity, number counts, and the luminosity function \\citep{2005astro.ph..6044F, 2005astro.ph..6661V}.  de Lucia et al. find that the star formation histories of massive ellipticals peak at z$\\approx$5 while those of less massive ellipticals peak at progressively smaller redshifts and are more extended.  These simulations predict a mass-metallicity relation, with the most massive ellipticals having solar metallicity and the least massive ones being a factor or ten smaller in metallicity.  Gradients in SPs are difficult to model within the framework of hierarchical formation.  Mergers between disks can presumably preserve existing gradients in the subcomponents and produce new gradients through gas infall, but mixing from successive mergers might erase any correlations between gradients and global properties.  White (1980) found that the metallicity gradient in a disk galaxy was halved after three mergers with similar sized disks.  However, \\cite{1999ApJ...513..108B} found that more massive galaxies had steeper metallicity gradients.  The impact of secular evolution on gradients is not straightforward. Since the resulting pseudobulge has a smaller scale length than the progenitor disk, an existing disk gradient could become amplified (A Klypin, private communication). However, mixing during secular evolution can have the effect of washing out existing gradients.  Adding gas only complicates the picture.  If gas is fueled towards the central regions by bars, this could result in a nucleus that is younger and more metal-rich than the outer regions of the bulge.  Simulations by Friedli et al. (1994) resulted in a flattening of metallicity gradients in all but the innermost regions where a starburst, fueled by infalling gas, produced a metal-rich nucleus.   Abundance ratios can place additional constraints.  Mg and other $\\alpha$-elements are primarily produced in Type II Supernovae (SN II) while a substantial fraction of the Fe-peak elements Fe and Cr are produced in Type Ia Supernovae (SN Ia). Therefore, $\\alpha$-enhancement is generally attributed to a cessation of star formation before the bulk of SN Ia occurred.  Through chemical evolution modeling, Thomas et al. (1999) found that a clumpy collapse model produced uniform $\\alpha$-enhancement or positive gradients (increasing $\\alpha$/Fe with radius) while a merger model produced uniformly solar $\\alpha$/Fe or negative gradients.  For the case of secular evolution, \\cite{2004A&A...413..547I} make different predictions for abundance ratios in bulge stars depending on whether it is the gas disk or the stellar disk which first becomes unstable.  In the former case, gas clumps merge together and spiral inward, causing massive starbursts and producing large $\\alpha$/Fe ratios.  In the latter case, a bar forms and then channels gas towards the center.  This occurs on long timescales, resulting in smaller $\\alpha$/Fe ratios.   The only bulges where individual stars can be resolved are those of the Milky Way (MW) and M31.  The majority of the stars in the MW bulge are old (t$\\geq$7 Gyr), although young (t$\\leq$200 Myr) and intermediate-age (200 Myr $\\leq$t$\\leq$7 Gyr) stars are also detected \\citep{1995MNRAS.275..605I, 1996AJ....112..171S, 2000A&A...355..949F, 2003MNRAS.338..857V, 2003A&A...399..931Z}.  As a barred late-type spiral, the MW might be a good candidate for secular evolution but that hypothesis is challenged by the mean stellar age of its bulge.  If the bulge were produced through a rearrangement of disk stars, this must have occurred several Gyr ago if the inner disk has the same age distribution as the disk at the solar neighborhood.  The young stars in the bulge are mainly found in the innermost regions.  While the gas that formed them could have been driven by a bar, it could just as well have been provided by a recent merger.  The mean metallicity of the Galactic bulge is slightly sub-solar \\citep{1996ApJ...459..175M, 2000A&A...355..949F,2003MNRAS.338..857V, 2003A&A...399..931Z}(Recent Fullbright et al. 2006).  The stellar content of M31's bulge is not as well understood as that of the MW since we cannot reach its main sequence turnoff.  Observations of giant stars are consistent with M31's bulge being similar in age to the MW bulge and slightly more metal-rich\\citep{1999ASPC..192..215R, 2000A&A...359..131J, 2001AJ....122.1386D, 2003AJ....125.2473S, 2005AJ....130.1627S}.   SPs in more distant bulges have to be studied through photometry or spectroscopy of integrated light.  An important limitation of such studies is that integrated light is dominated by the most luminous stars.  Colors have been studied more extensively as they have the advantage of higher signal-to-noise (S/N).  Pioneering work by \\cite{1994AJ....107..135B} found that color variations from galaxy to galaxy are much larger than color differences between disk and bulge in each galaxy.  Similarly, \\cite{1996A&A...313..377D} found that bulge and disk colors are correlated and that the color differences between bulge and disk suggested that the SPs did not vary much from one to the other.  Unfortunately, color studies suffer from degeneracies between ages, metallicities, and extinction.   Line strengths are nearly insensitive to dust \\citep{2005ApJ...623..795M}, provide information on the abundances of several elements and molecules, and allow for breaking the age-metallicity degeneracy.  \\cite{1994ApJS...95..107W} obtained line strengths for a range of single age, single metallicity SPs (SSPs) on the Lick/IDS system \\citep{1984ApJ...287..586B, 1985ApJS...57..711F} and found that while individual indices are sensitive to both age and metallicity, the relative sensitivity varies from index to index.  Spectral indices have also been defined at high resolution \\citep{1999ApJ...513..224V}, allowing better age determinations than otherwise possible.  One of the limitations of the original models is that they were calibrated using galactic stars, few of which had abundance ratios different from solar.  Much progress has since been made in extending Lick indices to non-solar abundance ratios \\citep{1995AJ....110.3035T, 2000AJ....119.1645T, 2003MNRAS.339..897T, 2004MNRAS.351L..19T, 2005ApJS..160..176L}.   Line strengths have been used extensively to characterize the SPs of ellipticals.  The luminosity-weighted ages of cluster ellipticals are large while those of field ellipticals are on average smaller, with a large spread \\citep{1994AJ....108.2054R, 2000AJ....119.1645T, 2001ApJ...551L.127V, 2002MNRAS.333..517P, 2005MNRAS.358..813D, 2005ApJ...621..673T}. This goes against the collapse model and confirms, at least qualitatively, the prediction of the merger model by Kauffmann (1996).  Mg-sensitive indices in ellipticals are more tightly correlated with central velocity dispersion than Fe-sensitive indices, resulting in a correlation between Mg/Fe and $\\sigma_0$ \\citep{1993ApJ...411..153B}.  \\cite{2003ApJ...586...17W} found that that the Mg-$\\sigma$ relation of ellipticals is consistent with these objects having been formed through around 50 mergers with merger probability constant or midly declining with time.   There have been fewer studies of line strengths in bulges.  Integrated light studies on the bulges of the MW and M31 have arrived at similar ages and metallicities as the resolved studies \\citep{2002A&A...395...45P, 2005A&A...434..909P}.  Both bulges have large SSP ages.  M31 is slightly super-solar in SSP metallicity while the Milky Way is solar.  Both are $\\alpha$-enhanced with M31 being more so in line with its larger $\\sigma_0$.  Early studies on extragalactic bulges found them to be similar to ellipticals in their central line strengths \\citep{1996AJ....112.1415J, 1996AJ....112.2541I}.  \\cite{2002MNRAS.333..517P} found that bulges have smaller average luminosity-weighted age than ellipticals.  These authors did not find the correlation predicted by Kauffman (1996) for bulges and suggested that it might have been erased by secular evolution in late-types. The largest sample of bulges to date was that of \\cite{2001A&A...366...68P}, who identified three classes of bulges: a) young bulges which are small, have ionized gas, low velocity dispersions, and low metallicity; b) old bulges that are alpha-enhanced and follow the mass-metallicity relation of ellipticals; and c) bulges that have a mixture of young and old populations, which are less alpha-enhanced than those of class (b), and deviate from the Mg$_2$ relation of ellipticals.  Prugniel et al. and Proctor et al. found that both Fe and Mg were correlated with $\\sigma_0$ in bulges, resulting in the lack of a tight correlation between Mg/Fe and $\\sigma_0$ in bulges.  \\cite{2001A&A...366...68P} found that Mg$_2$ in bulges is more tightly correlated with the V$_{max}$ of the disk than with $\\sigma$, indicating that the SPs are more sensitive to the total galaxy potential (i.e. the dark matter halo) than the bulge potential.    Studies with spatial resolution offer several advantages to studies that only sample the central region.  First, differential studies of ages and abundances are more reliable than absolute estimates.  Second, formation models invariably make predictions for the global properties of galaxies which are better traced by mean observed quantities than central ones; observations with spatial resolution allow estimation of mean values.  Finally, as mentioned already, population gradients can place additional constraints on formation mechanisms.   Line strength gradients have been studied extensively in ellipticals. \\cite{1993MNRAS.265..553C} and \\cite{2005MNRAS.tmpL..47F} find strong correlations between gradients and physical properties while others find weak (Mehlert et al. 2003) or no \\citep{1999ApJ...527..573K} correlations.  There have been relatively few studies on gradients in bulges.  Fisher et al. (1996) found steeper metallicity gradients along the minor axes of nine edge-on S0s than along the major axes, suggesting different formation mechanisms for the bulge and the disk. \\cite{1999Ap&SS.269..109G} found that gradients were correlated with luminosity in 16 bulges.  \\cite{2000MNRAS.311...37P} found that gradients correlated with velocity dispersion, albeit with a sample of only four galaxies, while \\cite{2002Ap&SS.281..367J} found no such correlation.   Integral field spectroscopy has enabled the acquisition of 2D line strengths in bulges with results just starting to emerge (e.g. Sil'chenko et al. 2003; Falc{\\'o}n-Barroso et al. 2004).  Recent work by \\cite{2005MNRAS.358.1337R} shows that tunable filters might be another way to obtain 2D line strengths.   In this paper, we present line strengths and line strength gradients in the bulges and inner disks of 38 galaxies.  Our sample, described in Section 2, was chosen to span a range of bulge properties and specifically targeted several galaxies with blue bulges and similar bulge/disk colors and/or disk-like kinematics in an attempt to look for SP signatures of secular evolution.  Section 3 describes the observations and data analysis.  Section 4 describes the SP results and Section 5 discusses their implications for bulge formation scenarios.  Section 6 contains a summary.  The structure, kinematics, and dynamics and how they relate to the SPs will be discussed in a future paper (hereafter Paper II). ",
        "conclusions": "We have studied line strengths in the bulges and inner disks of 38 galaxies in the local universe.  Our galaxies span a wide range of Hubble types, central velocity dispersions, maximum disk rotational velocities, and inclinations.  The low-inclination galaxies include barred and unbarred objects; the edge-on galaxies include those with and without boxy/peanut-shaped bulges.  We included several galaxies whose bulges were previously identified as being disk-like in their colors or kinematics to see if their spectral properties reveal evidence for secular evolution.  We use the [MgFe]' index and five Balmer indices to characterize the luminosity-weighted metallicities and ages of the SPs and the Mgb/$\\langle$Fe$\\rangle$ index to characterize the $\\alpha$/Fe ratios.  Our main results are the following:  \\begin{itemize}   \\item The central regions of bulges range in SSP metallicity from [Z/H]=-0.8 to +0.7 dex and in SSP age from less than 2 to greater than 15 Gyr.  \\item The central ages and metallicities are sensitive to bulge color which is in turn sensitive to central velocity dispersion and maximum disk rotational velocity.  \\item Red bulges of all Hubble types are similar to luminous ellipticals in their central SPs.  They have large SSP ages and are super-solar in SSP metallicity and $\\alpha$/Fe.   \\item Blue bulges can be separated into two classes: a metal-poor class that is restricted to late-types with small velocity dispersion and a young, metal-rich class that includes all Hubble types and velocity dispersions.  The metal-poor blue bulges are actively forming stars while the metal-rich ones are not.  Low-luminosity ellipticals exhibit a similar range of SSP ages and metallicities as blue bulges.   \\item Luminous ellipticals and the different types of bulges form a continuous and overlapping sequence on diagrams of metallicity- and age-sensitive indices versus $\\sigma_0$.  At fixed $\\sigma_0$, there is no systematic difference between bulges and ellipticals on these diagrams but bulges exhibit larger scatter.  At fixed $\\sigma_0$, age and metallicity are more tightly anticorrelated in ellipticals than in bulges.   \\item $\\alpha$/Fe in red bulges is correlated with $\\sigma_0$ and V$_{max}$.  Red bulges and ellipticals follow the same $\\alpha$/Fe-$\\sigma_0$ relation.   \\item Most blue bulges (11 out of 14) are consistent with having solar $\\alpha$/Fe. At fixed $\\sigma_0$, blue bulges have smaller $\\alpha$/Fe than red bulges and ellipticals.   \\item Barred galaxies appear to have larger central metallicities than unbarred galaxies of the same $\\sigma_0$ and V$_{max}$.   \\item Most galaxies show a steady decrease in metallicity-sensitive indices with radius.  The slope of the gradient is correlated with the central value and therefore with the global kinematics.  The bulge- and disk-dominated regions are distinct in their line strength profiles, with the disks generally having shallower slopes.  The smallest bulges do not have negative line strength gradients; some of these have flat profiles in the central region while others have positive gradients.   \\item There is a correlation between [MgFe]' strength in the bulge and the disk. This correlation holds for all galaxies, not just those with bars, blue bulges, or bulges identified as having disk-like structural or kinematical properties.   \\item Where positive age gradients (with the central regions being younger) are present, they are invariably in barred galaxies.  This suggests that bar-driven star formation has occurred.  However, several red bulges in barred galaxies have large central SSP ages (although it could be younger than the outer regions) which means there has been no significant bar-driven star formation for several Gyr.   \\item Four galaxies have super-solar $\\alpha$/Fe in the disk-dominated region.  The rest are consistent with having solar $\\alpha$/Fe in the disk.   \\item Objects identified as having disk-like structural or kinematic properties do not have noticeably different SPs than other bulges.  They follow the same scaling relations as the red bulges and ellipticals and have metallicity gradients.  The three barred S0s identified as having bulges with disk-like structural and kinematic properties are also $\\alpha$-enhanced and therefore do not resemble the majority of the disks, including the MW disk at the solar neighborhood.   \\item Color profiles agree frequently but not always with line strength profiles. Where there is a discrepancy, due likely to the colors being affected by dust, it is usually in the central regions.  Consequently, color gradients (computed as the difference in color between the center and a characteristic scale length) do not necessarily correlate with [MgFe]' gradients, illustrating the value of spectroscopy.   \\end{itemize}   Overall, our results are consistent with the hypothesis that mergers have been the dominant mechanism responsible for the formation of bulges.  However, some of the observations, such as the correlation between bulge and disk metallicity, pose significant challenges to the merger scenario.  Furthermore, the possibility that barred galaxies follow different scaling relations than unbarred galaxies and are overrepresented among galaxies with age gradients supports the secular evolution picture.  Central line strengths on a statistically significant sample of ellipticals and bulges of barred and unbarred spirals would be invaluable in determining whether more than one formation mechanism is required for bulges.  The necessary data are already available in the databases of large surveys such as the SDSS.  Spatially resolved studies on a smaller, representative sample, would allow for better comparisons between gradients in different types of galaxies."
    },
    "0512/astro-ph0512352_arXiv.txt": {
        "abstract": "We describe a method of estimating the abundance of short-period extrasolar planets based on the results of a photometric survey for planetary transits. We apply the method to a 21-night survey with the 2.5m Isaac Newton Telescope of $\\sim$32000 stars in a $\\sim 0.5\\degr \\times 0.5\\degr$ square field including the open cluster NGC~7789. From the colour-magnitude diagram we estimate the mass and radius of each star by comparison with the cluster main sequence. We search for injected synthetic transits throughout the lightcurve of each star in order to determine their recovery rate, and thus calculate the expected number of transit detections and false alarms in the survey. We take proper account of the photometric accuracy, time sampling of the observations and criteria (signal-to-noise and number of transits) adopted for transit detection. Assuming that none of the transit candidates found in the survey will be confirmed as real planets, we place conservative upper limits on the abundance of planets as a function of planet radius, orbital period and spectral type. ",
        "introduction": "Photometric surveys for transiting extra-solar planets have become very popular since the detection of the transits exhibited by the planet-host star HD 209458 (\\citealt{cha00}; \\citealt{bro01}). For the first time the radius of an extra-solar planet was determined, and the measurement of the orbital inclination lead to an estimate of the planetary mass, not just a lower limit. The planet HD 209458b was found to have an average density of $\\sim$0.38 g/cm$^3$, significantly less than the average density of Saturn (0.7 g/cm$^3$), leading to the term ``hot Jupiter'' for the class of Jupiter mass planets with short periods (1-10 d). Transiting planets are also very important in that their atmospheric composition may be determined from transmission spectroscopy for the brighter host stars (\\citealt{cha02}; \\citealt{bro02}; \\citealt{vid03}). Careful modelling of the transit morphology and/or timings may be used to constrain the presence of moons or rings and to probe the limb darkening of the star (\\citealt{bro01}).  Since the discovery of the transiting nature of HD 209458b, many transit candidates have been put forwards by various groups (e.g: \\citealt{str03}; \\citealt{dra04}; \\citealt{bra05}). OGLE have been by far the most prolific transit survey with 177 transit candidates from three observational seasons (\\citealt{uda02a}; \\citealt{uda02b}; \\citealt{uda03}; \\citealt{uda04}). However, even with the discovery of numerous candidates, follow-up observations have confirmed the planetary status of only six, bringing the total number of transiting planets to nine (see \\citealt{bra05} and references therein; \\citealt{sat05}; \\citealt{bou05}). This is due to the ubiquity of eclipsing binaries and the many observational scenarios involving these systems that mimic a transit event (\\citealt{bro03}). Spectroscopic and multi-band photometric observations are required to rule out the eclipsing binary scenarios and determine the mass of the companion (e.g: \\citealt{alo04}).  When hunting for new planets, the main advantage of the transit method over the radial-velocity technique is that many stars may be monitored in parallel and to fainter magnitudes, thus probing out beyond the Solar neighbourhood. Even though only a small fraction ($\\sim$0.1\\%) of stars are expected to exhibit a hot Jupiter transit signal, by using a large field of view instrument on a crowded star field one can monitor enough stars to the precision required to detect a number of transiting planets. Consequently large charge-coupled device (CCD) mosaic cameras are essential to the planet catch potential of a transit survey.  A transit survey produces transit candidates that need follow-up observations to determine the nature of the transit signals. Candidates confirmed as transiting planets add to our database of extra-solar planets and constrain their poorly known mass-radius relationship (\\citealt{bur04}). To estimate the fraction of stars that harbour a planet (the planet fraction) as a function of spectral type and planet type we compare the number of transiting planets detected with a calculation of the expected number of transiting planet detections. Even when zero planets are detected (a null result) such a calculation places upper limits on the planet fraction.  \\begin{table*} \\centering \\caption{The different subsets of stars used when calculating the expected number of transiting planet detections and false alarms. \\label{tab:spectypes}} \\begin{tabular}{lcrr} \\hline Set Of Stars & Mass Range & No. Of Stars & No. Of Stars With 1999-07 Lightcurve Data \\\\ \\hline All Stars    & $0.08 M_{\\sun} \\leq M_{*} \\leq 1.40 M_{\\sun}$ & 32027 & 20949 \\\\ Late F Stars & $1.05 M_{\\sun} \\leq M_{*} \\leq 1.40 M_{\\sun}$ & 3129  & 2780  \\\\ G Stars      & $0.80 M_{\\sun} \\leq M_{*} \\leq 1.05 M_{\\sun}$ & 7423  & 6711  \\\\ K Stars      & $0.50 M_{\\sun} \\leq M_{*} \\leq 0.80 M_{\\sun}$ & 15381 & 9690  \\\\ M Stars      & $0.08 M_{\\sun} \\leq M_{*} \\leq 0.50 M_{\\sun}$ & 6094  & 1768  \\\\ \\hline \\end{tabular} \\end{table*}  In this paper we describe a Monte Carlo method for calculating detection probabilities (and false alarm rates) of transiting planets based on photometric data, as functions of various parameters, taking into account the following factors: \\begin{enumerate} \\item Limb darkening effects which tend to make central eclipses deeper and grazing eclipses shallower. \\item The effect of orbital inclination on the shape and width of the transit lightcurve. \\item The distribution of the photometric data in time and the individual error bars on each measurement. \\item The signal-to-noise threshold, number of transits, and number of data points in-transit and out-of-transit required for a detection. \\end{enumerate} We then apply the method to the transit survey described in \\citealt{bra05} (here on referred to as BRA05) to determine the expected number of transiting planet detections and place limits on the planet fraction as a function of star and planet type.  In Section 2 we describe the lightcurve data used in the analysis and in Section 3 we define the detection probabilities and false alarm probabilities for an extra-solar planet based on photometric data. In Section 4 we present the Monte Carlo method that we used to calculate these probabilities and derive limits on the hot Jupiter fraction in the field of NGC~7789 as a function of star and planet type. In Section 5 we discuss the results and in Section 6 we present our conclusions. ",
        "conclusions": "The calculation of the expected number of transiting planet detections for a transit survey has been discussed in varying levels of detail by several authors (e.g: \\citealt{gil00}; \\citealt{wel05}; \\citealt{hid05}; \\citealt{hoo05}; \\citealt{moc05}). A requisite for such a calculation is an estimate of the masses and radii of the stars in the sample from observational constraints or a model for the star population that predicts these properties.  In this paper, we present a method for calculating in detail the detection probabilities (and false alarm rates) of transiting planets for photometric time-series data as functions of detection threshold, planetary radius and orbital period. The calculation is based on Monte Carlo simulations and requires the properties of the host stars to be known, either from observational constraints as in the case of stellar clusters or a model for the star population. We have shown how to convert these probabilities into an expected number of transiting planet detections as a function of star and planet type. For a null result in a transit survey, this information can be used to determine a significant upper limit on the planet fraction.  In the case of the transit survey of NGC~7789 presented in BRA05 we have derived upper limits on the planet fraction for the F, G, K and M stars in the sample and for three relevant period ranges of hot Jupiters. We have also derived how these limits scale with detection threshold and planetary radius. In BRA05, it is estimated that the survey expects to detect $\\sim$2 HD 209458b-like transiting planets or $\\sim$4 OGLE-TR-56b-like transiting planets using simple arguments. Our results indicate that for HD 209458b (3-5~d hot Jupiter with $R_{\\mbox{\\small p}}~\\sim~1.4~R_{\\mbox{\\small J}}$), and under the assumption that $f_{\\mbox{\\small p}} \\approx$~1\\%, we also expect to detect $\\sim$2 such transiting planets. Similarly for OGLE-TR-56b (1-3~d very hot Jupiter with $R_{\\mbox{\\small p}}~\\sim~1.2~R_{\\mbox{\\small J}}$), and again under the assumption that $f_{\\mbox{\\small p}} \\approx$~1\\%, we expect to detect $\\sim$3 such transiting planets. It is encouraging to note the agreement between the two methods although the simpler method from BRA05 may tend to slightly over estimate the planetary detection rate. We conclude that the transit survey presented in BRA05 reached the sensitivity required in order to detect a few hot Jupiters if the abundance of such planets in the field of NGC~7789 is similar to that of the Solar neighbourhood.  Improved survey design, mainly by employing a longer survey duration, will greatly improve the sensitivity to hot Jupiters. It is well known that metal rich stars have a much higher probability of hosting an extra-solar planet (\\citealt{san04}) and hence careful choice of the target star population will increase the probability of a detection. Even in the presence of a null result, a more sensitive survey will allow the derivation of tighter limits on the planet abundance."
    },
    "0512/astro-ph0512487_arXiv.txt": {
        "abstract": "We report on a large scale numerical study of networks of semilocal cosmic strings in flat space in the parameter regime in which they are perturbatively stable. We find a population of segments with an exponential length distribution and indications of a scaling network without significant loop formation.  Very deep in the stability regime strings of superhorizon size grow rapidly and ``percolate'' through the box. We believe these should lead at late times to a population of infinite strings similar to topologically stable strings.  However, the strings are very light; scalar gradients dominate the energy density and the network has thus a global texture-like signature. As a result, the observational constraints, at least from the temperature power spectrum of the CMB, on models predicting semilocal strings, should be closer to those on global textures or monopoles, rather than on topologically stable gauged cosmic strings. ",
        "introduction": "The improvement in observational data during the last years, especially on Cosmic Microwave Background (CMB) radiation, has shown that topologically stable cosmic strings cannot be the dominant component creating the primordial fluctuations that seed large scale structure as we see it in the Universe today. Strings could be a subdominant component~\\cite{Wyman:2005tu} but there is as yet no evidence for their existence (see~\\cite{Sazhin:2003cp} for a candidate string lensing event that turned out to be two galaxies). On the other hand, from the theoretical point of view, the formation of topological defects via the Kibble mechanism~\\cite{Kibble:1976sj} is inevitable in theories that allow stable defects. The formation of cosmic strings after inflation appears to be generic both in supersymmetric grand unified theories breaking down to the Standard Model~\\cite{Jeannerot:1997is,Jeannerot:2003qv} and in brane inflation models within the framework of fundamental superstrings~\\cite{Sarangi:2002yt}. Therefore, models that do not predict topologically stable defects at the end of inflation are somewhat appealing. Recent work~\\cite{Urrestilla:2004eh} suggests that a particular type of non-topological defects, called {\\it semilocal} strings, could form in the early universe both after $D$-term inflation \\cite{Dvali:1994ms} and after brane inflation with $D3/D7$-branes \\cite{Dasgupta:2004dw,Chen:2005ae}.   The semilocal model~\\cite{Vachaspati:1991dz} is a minimal extension of the Abelian Higgs model by a global $SU(2)$ symmetry. With two equally charged Higgs fields and only one $U(1)$ gauge field, there are not enough gauge degrees of freedom to cancel all scalar gradients, even outside the string cores, unless the scalar windings are correlated.  Semilocal (SL) strings are stable for sufficiently large gauge coupling~\\cite{Vachaspati:1991dz,Hindmarsh:1991jq}, but due to their non-topological character they appear after the phase transition as open segments. These are closely related to electroweak dumbbell configurations \\cite{Nambu:1977ag,Vachaspati:1992fi,Urrestilla:2001dd} but whose ends have long-range interactions that resemble those of global monopoles~\\cite{Hindmarsh:1992yy}, with a force roughly independent of distance.  Therefore the network of SL strings has several features that are not present in networks of topologically stable Abrikosov-Nielsen-Olesen (ANO) strings~\\cite{VS} and their cosmological evolution could be quite different~\\cite{Benson:1993at,Achucarro:1999it}.  Semilocal string segments can either contract and eventually disappear or grow to join a nearby segment and form a longer string. Closed loops can be formed by intercommutations, and also if the two ends of a segment join.  The obvious question is if at some point this joining of segments can form infinitely long strings.  If so, the networks present characteristics intermediate between topological ANO strings and global defects -- both of which show scaling behaviour in an expanding universe \\cite{VS} - and we would like to understand their cosmological impact.  The purpose here is to report on a large scale numerical study of SL string networks. The non-topological nature of the SL strings does not permit a Nambu-Goto approach and our results are based on classical field theory simulations.  Our basic results are twofold.  We find that SL strings will, in the right parameter range, form a network that shares features with a network of ANO strings due to the existence of fast growing, extremely long percolating strings, which we believe would lead to an infinite string component at late times (we use the word ``percolating'' in a non-technical sense, to mean that the strings span the simulation box).  On the other hand, the energy density is dominated by scalar gradients, like in models for global defects, even in the case of dense string networks deep in the stable regime.  As the energy density plays a crucial role for the observable cosmological consequences, like the effect on the temperature power spectrum of the CMB, the signature of a SL string network is thus likely to be very similar to that of global defects. ",
        "conclusions": "In this paper we studied the evolution of semilocal string networks deep in the stability regime using high performance computing to carry out numerical field theory simulations in flat space. Our main results are:  $\\bullet$ A population of segments with an initially exponential length distribution. Short segments vanish rapidly and the mean string length grows.  $\\bullet$ Appearance of superhorizon strings (much larger than the box size) percolating through the box deep in the stability regime, $\\sqrt{\\beta} \\lesssim 0.3$. For $\\sqrt{\\beta}\\lesssim 0.25$ these fast growing strings fall outside the exponential distribution.   $\\bullet$ No significant population of small loops. In dense networks $\\sqrt{\\beta} \\lesssim 0.4$ about one third of the simulations show a loop of a larger size.  $\\bullet$ The string network can persist for a long time and for low $\\beta$ the results are consistent with linear scaling.  $\\bullet$ The energy density is dominated by scalar gradients.   These findings suggest that SL networks may have a cosmological signature, in particular in the CMB and large scale structure, closer to that of global defects (monopoles or textures) than to topological strings. Current bounds~\\cite{Durrer:2001cg} on global textures in CMB allow about 13\\% of the signal to come from defects at $\\ell = 10$.   The relevance of these results to the cosmological scenarios in~\\cite{Urrestilla:2004eh,Dasgupta:2004dw} would require further study in expanding backgrounds (see also \\cite{Benson:1993at}) and the development of analytic characterizations of the network to verify the apperance of an infinite string population at late times, as was done in ref. \\cite{Copeland:1998na} for the case of dense networks of cosmic string loops. Apart from the expansion, a complete quantum field theoretical study of the conditions at formation is still lacking, as pointed out in~\\cite{Dasgupta:2004dw}.  If it is confirmed that the observational signal of the semilocal network is similar to global textures or monopoles in the temperature power spectrum \\cite{BKHLU}, there exists a potential problem for distinguishing between models via future data. A dedicated study is needed to find out if there is any difference e.g. in polarisation of vector modes, or in the possibility of gravitational waves radiating from oscillating loops \\cite{VS}."
    },
    "0512/astro-ph0512164_arXiv.txt": {
        "abstract": "s{ Gamma-ray bursts are known to be sources of high-energy $\\gamma$ rays, and are likely to be sources of high-energy cosmic rays and neutrinos. Following a short review of observations of GRBs at multi-MeV energies and above, the physics of leptonic and hadronic models of GRBs is summarized. Evidence for two components in BATSE and EGRET/TASC data suggest that GRBs are sources of high-energy cosmic rays. GLAST observations will reveal the high-energy $\\gamma$-ray power and energy releases from GRBs, and will provide detailed knowledge of anomalous high-energy emission components, but confirmation of cosmic ray acceleration must await 100 TeV -- PeV neutrino detection from GRBs. } ",
        "introduction": "Gamma-ray burst (GRB) studies represent one of the most dynamic fields in contemporary astronomy, and the field continues to evolve rapidly as new instrumentation comes online. Knowledge of GRBs in the high-energy, multi-MeV regime is poised to undergo great advances in the near future.  Right now, Swift is exploring the late prompt and early afterglow phases of GRBs at X-ray and hard X-ray energies with unprecedented detail. The ground-based air Cherenkov telescopes, including HESS, VERITAS, and MAGIC, are reaching better sensitivities and lower thresholds with the goal of detecting GRBs at $\\sim 100$ GeV -- TeV energies. The MAGIC telescope has already demonstrated the ability to slew within $\\approx 30$ s to a GRB. AGILE, a small scientific mission developed by the Italian Space Agency, is due for launch next year. With sensitivity comparable to EGRET, though with a much larger field-of-view (factor-of-two decline in sensitivity at an off-axis angle of $\\sim 60^\\circ$, versus $\\sim 25^\\circ$ for EGRET), it should detect some 5 -- 10 GRBs per year above 100 MeV. The GLAST mission, which includes both the Gamma-ray Burst Monitor and Large Area Telescope that will cover energies from $\\approx 5$ keV and above, is scheduled for launch in September 2007.  Its peak effective area at $\\approx 1$ GeV will be $\\approx 8\\times$ greater than EGRET's at $\\approx 100$ MeV, and its field-of-view represents $\\approx 1/6^{\\rm th}$ of the full sky, compared to $\\approx 1/20^{\\rm th}$ for EGRET. This should lead to the detection of some $30$ -- 100 GRBs per year at $\\gtrsim 100$ MeV energies.  Beyond the electromagnetic realm, there is the real possibility that GRBs will be the first detected sources of high-energy ($\\gg 1$ TeV) neutrinos. The south pole AMANDA neutrino detector is, as new strings are added, evolving into the km$^3$ IceCube neutrino telescope over the course of this decade. AMANDA has already detected thousands of cosmic-ray induced background neutrinos events, and improved sensitivity and background rejection should soon reveal meaningful upper limits on GRB source models if not direct detections.  If GRBs are sources of ultra-high energy cosmic rays (UHECRs), then detection of $\\gtrsim 10^{17}$ eV neutrino coincident with GRBs may be possible with Auger and ANITA. In view of the rapid technological advances expected over the next 5 years, therefore, we can expect that our knowledge of GRBs as high-energy sources will vastly increase.  In this contribution, I summarize what is now known from observations of high-energy radiation from GRBs. This includes primarily the {\\it Compton Gamma Ray Observatory} results, which have provided the most significant detections of multi-MeV emission from GRBs. An important result from EGRET is that there is good evidence for two components in GRBs, and incontrovertible evidence for an extended phase, as observed from GRB 940217. In addition, an anomalous $\\gamma$-ray emission component was detected from GRB 941017. There is, moreover, a tantalizing suggestion of TeV emission from GRB 970417a, made with the Milagrito telescope.  Interpretation of these results is made within the blast-wave scenario, which was originally developed to understand GRB afterglows \\cite{mr97}. Although normally used to model leptonic afterglow emissions from GRBs, this scenario can also be used to predict emission signatures from hadrons accelerated by GRBs. If GRBs accelerate UHECRs, then anomalous $\\gamma$-ray emission signatures are expected from GRBs. Studies of high-energy radiation from GRBs have the potential to answer one of the outstanding questions in contemporary astronomy, namely the origin of the cosmic rays. Detection of $\\gtrsim 10^{14}$ eV  neutrinos from GRBs would provide compelling evidence in support of this solution. ",
        "conclusions": "GRBs are established to be sources of high-energy $\\gamma$ rays.  Here we have argued GRBs also accelerate hadrons, and that GRBs are sources of high-energy cosmic rays. Evidence for this is suggested by the slow decay of $\\gamma$ rays from GRB 941017, and by the anomalous $\\gamma$-ray emission component in GRB 941017.  Several types of observations can test this hypothesis. Most unambiguous is the detection of high-energy neutrinos from a GRB, which would require an ultra-relativistic hadronic component that is much more powerful than the nonthermal electron component that produces the hard X-ray and soft $\\gamma$-ray emissions from GRBs \\cite{da03}. Another prediction is the detection of hadronic emission components in the spectra of GRBs, as observed in GRB 941017 \\cite{gon03}. A third observation that would implicate GRBs as the sources of HECRs is the detection of high-energy neutron $\\beta$-decay halos around star-forming galaxies \\cite{der02}, including $\\beta$-decay emission from GRBs in the Milky Way \\cite{ikm04}.  \\vskip0.1in I would like to acknowledge joint research with Armen Atoyan, James Chiang, Jeremy Holmes, and Stuart Wick.  This work was supported by the Office of Naval Research and NASA GLAST grant DPR S-13756G."
    },
    "0512/astro-ph0512214_arXiv.txt": {
        "abstract": "{% }{ To study turbulence in the Orion Molecular Cloud (OMC1) by comparing observed and simulated characteristics of the gas motions. }{ Using a dataset of vibrationally excited H$_2$ emission in OMC1 containing radial velocity and brightness which covers scales from 70\\,AU to 30000\\,AU, we present the transversal structure functions and the scaling of the structure functions with their order. These are compared with the predictions of two-dimensional projections of simulations of supersonic hydrodynamic turbulence. }{ The structure functions of OMC1 are not well represented by power laws, but show clear deviations below 2000\\,AU. However, using the technique of extended self-similarity, power laws are recovered at scales down to 160\\,AU. The scaling of the higher order structure functions with order deviates from the standard scaling for supersonic turbulence. This is explained as a selection effect of preferentially observing the shocked part of the gas and the scaling can be reproduced using line-of-sight integrated velocity data from subsets of supersonic turbulence simulations. These subsets select regions of strong flow convergence and high density associated with shock structure. Deviations of the structure functions in OMC1 from power laws cannot however be reproduced in simulations and remains an outstanding issue. }{} ",
        "introduction": "Turbulence plays a central role in star-forming molecular clouds, acting both to support the clouds globally and to create local clumps and density enhancements that can undergo gravitational collapse. Simulations have shown that this latter process, known as turbulent fragmentation, may directly determine the initial mass function (IMF). Insight into the effects of turbulence on molecular clouds is thus essential for understanding the mechanisms of star formation. Such insight can only be gained by a close interplay between observations and simulations.  The characterization of turbulence may be achieved by statistical methods. Several techniques, such as the size-linewidth relation \\citep{larson1981,goodman1998,ossenkopf02,gustafsson2005}, probability distribution functions \\citep{falgarone1990,falgarone1994,miesch1999,ossenkopf02,pety2003,gustafsson2005}, structure functions \\citep{falgarone1990,miesch1994,ossenkopf02,gustafsson2005} and $\\Delta$-variance \\citep{bensch2001,ossenkopf02}, have previously been used to characterize the structure of brightness or velocity in molecular clouds. Comparisons between observations and models have earlier been made by \\cite{falgarone1994,falgarone1995,lis1998,joulain1998,padoan1998,padoan1999,pety2000,klessen2000,ossenkopf02,padoan2003,gustafsson2005}; see also the review by \\cite{elmegreen2004}. These earlier comparisons \\citep[save those in][]{gustafsson2005} are based on CO observations, tracing relatively cool and low density gas. Data are limited in spatial resolution and can only be used to address the physics at scales larger than roughly 0.03\\,pc (6000\\,AU).  In the present work we use IR K-band observations of vibrationally excited H$_2$ in the Orion Molecular Cloud (OMC1) to make a first comparison between observations and hydrodynamical simulations at the scales where individual stars are forming. In the region observed the H$_2$ emission is optically thin in the sense that it is not self-absorbed. There will be some obscuration by dust \\citep{rosenthal}, whose spatially differential nature is not known and is ignored here. The observations cover scales from 70\\,AU to 3$\\cdot10^4$AU. OMC1 is the archetypal massive star-forming region and the best studied region in the sky. OMC1 is highly active with widespread on-going star formation, exemplified by the presence of protostars, outflows and larger scale flows \\citep[see][ and references therein]{nissen2005}. In a previous paper \\citep{gustafsson2005}, using the same observational data as in the present work, we quantified the nature of turbulence in OMC1 by calculating size-linewidth relations, probability distribution functions and structure functions. It was shown that the turbulence at the small scales covered in these data generally follows the trends observed in CO data at larger scales. However, analysis also showed clear deviations from the fractal scaling observed at larger scales. These deviations could be ascribed to the presence of star formation and associated structures such as circumstellar disks. Here we use the structure functions of the radial velocities and the scaling of the structure function exponents from our observational data to compare with a numerical simulation of supersonic, compressible, hydrodynamic turbulence. The scaling of structure functions has earlier been analysed in \\cite{padoan2003} where the column densities of $^{13}$CO were used.    The structure function of order $p$ of the velocity vector $\\vec{u}$ is defined here as \\begin{equation} S_p(r)=\\langle \\left|[\\vec{u}(\\vec{x})-\\vec{u}(\\vec{x}-\\vec{r})] \\cdot\\vec{e}\\right|^p \\rangle \\propto r^{\\zeta_p} \\label{eq:struc} \\end{equation} where $\\vec{e}$ is a unit vector parallel (longitudinal structure function) or perpendicular (transversal structure function) to the vector $\\vec{r}$, and $r=|\\vec{r}|$. The average is taken over all spatial positions $x$. The modulus sign in our definition (\\ref{eq:struc}) is adopted to improve the statistics for odd moments. In our case the data consist of projected radial velocities. We measure differences in radial velocity across the plane of the sky and we are therefore dealing with transversal structure functions. The structure functions of fully developed turbulent fields are known to follow power laws in the inertial range, $\\eta \\ll r \\ll L$, where $\\eta$ is the dissipation scale and $L$ is the integral scale. The set of scaling exponents, $\\zeta_p$ in Eq.~(\\ref{eq:struc}), can therefore be determined \\citep{frisch1995}. The scaling exponents are expected to be characteristic of the turbulence involved and universal for all scales and Reynolds numbers. The transversal and longitudinal structure functions are anticipated to have the same scaling in the infinite Reynolds number limit. This may not however be the case at moderate Reynolds numbers, where it has been found that $\\zeta_{p,{\\rm long}} > \\zeta_{p,{\\rm trans}}$ for $p >$~3 in incompressible hydrodynamical experiments and simulations \\citep[][ and references therein]{kerr2001}.  \\cite{kolmogorov1941} found from the energy conservation law in incompressible, isotropic and homogeneous turbulence that $\\zeta_3$=1. However, \\cite{dubrulle1994} suggested that ratios of scaling exponents, say $\\zeta_p/\\zeta_3$, are inherently universal, while the individual scaling exponents may not be universal themselves. In this connection \\cite{frick1995} showed in the context of cascade models that one may have $\\zeta_3 \\neq 1$ and yet recover scaling laws for the structure functions in good agreement with the She-Leveque model (see below) for the ratio $\\zeta_p/\\zeta_3$.  \\cite{she1994} described the scaling of velocity structure functions in incompressible turbulence by: \\begin{equation} \\zeta_p/\\zeta_3={p\\over9} + 2\\left[1-\\left(\\frac{2}{3}\\right)^{p/3}\\right], \\label{eq:she-leveque} \\end{equation} which is confirmed by simulations of nearly incompressible turbulence \\citep{padoan2004,haugen2004a} and by experiments \\citep{anselmet1984,benzi1993}. For supersonic turbulence \\cite{boldyrev2002a} obtained, as an extension to the She-Leveque model, the scaling: \\begin{equation} \\zeta_p/\\zeta_3={p\\over9} + 1-\\left(\\frac{1}{3}\\right)^{p/3}, \\label{eq:boldyrev} \\end{equation} which is confirmed by observations in the Perseus and Taurus molecular clouds \\citep{padoan2003} and simulations \\citep{boldyrev2002b,boldyrev2002c,padoan2004}. This type of scaling was originally proposed by \\cite{politano1995} for magnetohydrodynamic turbulence, where the dissipative structures are thought to be two-dimensional current sheets.    In Sect.~\\ref{sec:obs} we describe the observations and calculate the structure functions from the observed velocities. We then use the method of extended self-similarity (ESS) of \\cite{benzi1993} to find the structure function exponents and show that the scaling of these does not represent any known theoretical scaling as represented by Eqs.~(\\ref{eq:she-leveque}) and (\\ref{eq:boldyrev}). In Sect.~\\ref{sec:sim} we describe the simulation. In Sect.~\\ref{sec:sim_3d} we calculate the longitudinal and transversal structure functions of the 3D simulation and show that the scaling of the structure function exponents is similar to that of \\cite{boldyrev2002a} in contrast to the observations. In Sect.~\\ref{sec:subset} we choose subsets of the simulations which select the shocked gas seen in the observations, project the data onto 2D maps and calculate the structure functions. We show that if only strong shocks are included in the subset the scaling of the exponents is now similar to the scaling found in the observations. In Sect.~\\ref{sec:conclusion} we discuss our results. ",
        "conclusions": "\\label{sec:conclusion}   We have used observational data of shocked H$_2$ emission in OMC1 to show that structure functions at scales 70 -- 3$\\cdot 10^4\\,$AU (3.4$\\cdot 10^{-4}$ -- 0.15\\,pc) exhibit unusual scaling exponents for $p>3$. The scaling exponents are nearly constant for $p>3$ and smaller than predicted by both \\cite{she1994} and \\cite{boldyrev2002a}.  In three simulations we have selected shocked regions by imposing requirements on the value of the velocity divergence, $\\vec\\nabla\\cdot\\vec{u}$. In certain important respects the simulations presented here are then remarkably successful in reproducing the statistical behaviour observed in OMC1. In other equally important areas, they fail. Let us first reiterate the success.  We have found that by only including shocks that are relatively strong ($D_0 < -1.5$), the unusual scaling exponents of the observations are reproduced in the simulations. By contrast, a scaling following that of \\cite{she1994} or \\cite{boldyrev2002a} is found when all points in the simulations are included in the data. An explanation for this behaviour is as follows. Enhanced energy content at small scales, relative to larger scales, implies smaller slopes, that is, smaller values of $\\zeta_p$. Both the observational data of OMC1, and some of those subsets of the simulations selected only to include shocks, show that the values of $\\zeta_p$ are reduced for $p\\ge4$. Since structure functions of high order $p$ are dominated by regions of strong velocity differences, it follows that the observed excess of small scale energy is associated with regions of large velocity differences. These are likely to be the regions of strong shocks, as is evidenced by the fact that reduced values of $\\zeta_p$ are most clearly seen in subsets of the simulations that select the most strongly convergent high density regions.  The present work does not however furnish any explanation of why departure from the She-Leveque or Boldyrev scaling occurs at the specific value of $p\\ge4$. It is possible that the critical value of $p$ is in some way connected with the physical nature of the shocks, for example the fact that they are smoothed in the simulations, mimicking the structure of continuous (C-) type shocks, rather than jump (J-) type shocks \\citep[][ and references therein]{flower2003}.  We now turn to the failure of the simulations. The structure functions of the observations in OMC1 all deviate from power laws and exhibit clear bumps around $10^3$AU, exemplified by the third order structure function in Fig.~\\ref{fig:scal_obs}. This cannot in general be reproduced by the simulations. There appears to be two possible explanations for this observed behaviour. The first is that the deviation from power laws is due to protostar formation and associated outflows at a preferred scale. The second is that the behaviour is in some way inherent in the nature of the turbulence as opposed to the presence of protostars.  Turning to the first suggestion, the process of star formation pulls structure of a certain size out of the cascade and creates outflows, injecting energy into a turbulent background. Gravitational energy and angular momentum is spewed out of the star via such outflows and turned into local turbulence, hence increasing the overall turbulent content of the gas -- and restarting the whole cascade process. Such outflows are of course not present in the simulations.  The second suggestion, that the deviation from power law is somehow inherent in the nature of the turbulence, requires that there is some non-statistical element in this medium which is otherwise governed by statistical considerations. This may arise through our selection of strong shocks as a subset of the whole. We have seen in Fig.\\ref{fig:sim_512g} that traces of bumps in the structure functions are found in one snapshot of Run~2 when highly shocked material is selected. The deviations of the structure functions from power laws are not as pronounced in the simulation as in the observations of OMC1. However this finding provides some evidence that part of the explanation for the deviations from power laws of the structure functions in OMC1 arises from the fact that we observe preferentially shocked gas. As departures from power law behaviour are only evident in a single snapshot and not throughout the simulation at other times, this suggestion remains tentative.  In order to explore the reasons for the departure from power law behaviour, more advanced simulations are necessary. These should include self-gravity and energy feedback from protostellar zones through outflows and should ultimately incorporate ionization and magnetic fields."
    },
    "0512/astro-ph0512508_arXiv.txt": {
        "abstract": "A Monte Carlo simulation exploring uncertainties in standard stellar evolution theory on the red giant branch of metal-poor globular clusters has been conducted. Confidence limits are derived on the absolute $V$-band magnitude of the bump in the red giant branch luminosity function (\\vbump)  and the excess number of stars in the bump, \\rbump.  The analysis takes into account uncertainties in the primordial helium abundance, abundance of alpha-capture elements, radiative and conductive opacities, nuclear reaction rates, neutrino energy losses, the treatments of diffusion and convection, the surface boundary conditions, and color transformations.  The uncertainty in theoretical values for the red giant bump magnitude varies with metallicity between $+0.13/-0.12$ mag at $\\feh = -2.4$ and $+0.23/-0.21$ mag at $\\feh = -1.0$. The dominant sources of uncertainty are the abundance of the alpha-capture elements , the mixing length, and the low-temperature opacities.  The theoretical values of \\vbump\\ are in good agreement with observations.  The uncertainty in the theoretical value of \\rbump\\ is $\\pm 0.01$ at all metallicities studied. The dominant sources of uncertainty are the abundance of the alpha-capture elements, the mixing length, and the high-temperature opacities. The median value of \\rbump\\ varies from 0.44 at $\\feh = -2.4$ to $0.50$ at $\\feh = -1.0$.  These theoretical values for \\rbump\\ are in agreement with observations. ",
        "introduction": "}  Globular clusters are made up of a very large number of stars with varying mass but identical age, composition, and distance. This makes them a rich and productive application of the theory of stellar structure and evolution. Detailed stellar evolution calculations are done numerically using computer programs which incorporate previously calculated nuclear reaction rates and opacities, approximations to complex phenomena such as convection, and assumptions about the chemical composition. To compare theoretical models to observations, moreover, requires converting physical quantities such as luminosity and surface temperature to the observational system of magnitudes and colors using empirical relations or the results of separate stellar atmosphere models. The results of theoretical stellar evolution calculations, therefore, depend upon a set of prior assumptions. To assess the reliability of theoretical models of globular clusters stars one must study how uncertainties in these assumptions of stellar evolution theory propagate to the predictions of the theory.  In this work we are interested specifically in studying uncertainties in stellar evolution theory along the red giant branch (RGB) of metal-poor globular clusters. The RGB is the region in a color-magnitude diagram comprising low-mass stars that have already expended the hydrogen fuel at their center, and are burning hydrogen in a spherical shell around an inert helium core. Stars in this stage of evolution grow larger, brighter, and redder over time. An excellent review of RGB evolution is found in \\citet{sal02}, including a discussion of uncertainties in the theoretical models and observational tests of the theory. Much of our work was guided by this review.  After a low-mass star has consumed all the hydrogen at its center, the envelope of the star expands and it moves across the color-magnitude diagram from the main sequence region towards lower effective temperature and redder colors. Stars in this stage are said to be on the subgiant branch. Initially, hydrogen burning continues through a thick shell covering a region of mass $\\sim0.1M_\\odot$, while outside the hydrogen-burning shell the stellar envelope has the original chemical composition with a convective region in the star's outer layers. As the star progresses along the subgiant branch, the hydrogen-burning shell narrows to a mass of $\\sim0.001M_\\odot$ while the convective region grows steadily deeper. The convective region eventually reaches material previously processed by the hydrogen-burning core during the main sequence phase, and containing a higher abundance of helium which immediately is mixed throughout the convective region.  Eventually, the stellar luminosity begins to grow as the effective temperature continues to drop, and the star is said to be at the base of the RGB. The star progresses up the RGB as the hydrogen-burning shell moves outward through the star leaving behind an increasingly massive core of helium. The star's radius continues to grow larger, its effective temperature lower, and its luminosity brighter. The lower boundary of the convective region, which reached a maximum depth near the base of the RGB, recedes steadily.  A discontinuity in the chemical composition is left at the maximum depth reached by the convective envelope, since convection mixes hydrogen from the envelope into the partially depleted region left behind by the hydrogen-burning core of the main sequence phase. When the hydrogen-burning shell reaches this discontinuity, the sudden increase in available fuel causes the stellar luminosity to drop temporarily, and the star's evolution pauses before continuing its progression up the RGB. This point is called the RGB bump.  There is a very strong correlation between stellar luminosity and the mass of the helium core along the RGB. The luminosity grows as the core mass grows, and since the rate at which mass is added to the helium core is itself proportional to the luminosity, this means that throughout red giant evolution the helium core mass and the luminosity increase at an ever faster rate. The one exception to this is at the RGB bump, where the sudden change in chemical composition causes the growth in luminosity to pause temporarily, while the helium core continues to gain mass.  The probability of observing a star in a given luminosity range on the RGB is inversely proportional to the rate at which stars evolve at that luminosity. In observed globular cluster color-magnitude diagrams, therefore, we find that the number of observed stars steadily decreases along the RGB, except at the red giant bump. In fact, the differential luminosity function of a globular cluster, showing the number $N$ of observed stars as a function of magnitude, descends along the RGB as a nearly straight line in the magnitude-$\\log N$ plane. The slope of this line indicates how the rate of evolution increases along the RGB.  Globular cluster luminosity functions provide an important test of the accuracy of stellar evolution models. The number of stars observed as a function of luminosity indicates the relative timescale of stellar evolution, which in turn conveys information about the internal chemical structure of stars. A great deal of attention has been paid, for example, to the total number of stars on the RGB compared to the main sequence turn-off region. Several authors \\citep{bolte,van98,langer} have found a discrepancy in this quantity between observed luminosity functions for the globular clusters M5 and M30 and the predictions of stellar evolution theory. \\citet{van98} take this discrepancy to suggest the presence of rapidly rotating cores in red giants, while \\citet{langer} take it to indicate the possibility of a chemical mixing process deep in the stellar interior.  The magnitude of the RGB bump in observed globular clusters serves to indicate the depth of the chemical discontinuity left by the convective envelope at its point of deepest extent near the base of the RGB. In one of the first extensive studies of the red giant bump, \\citet{fusi} compared observational determinations of the bump magnitude in 11 globular clusters to theoretical predictions. They found that, taken relative to the horizontal branch magnitude, theoretical values of the bump magnitude are higher than observed values by about 0.4 mag. However, more recent studies using updated stellar evolution models and improved observations \\citep{cas3,zoccali,riello} have found no discrepancy between theory and observations.  \\citet{cas3} have studied how uncertainties in the most important individual stellar evolution parameters impact the magnitude of the red giant bump. They look at the equation of state, mixing length, mass loss along the RGB, opacities, and the $V$-band bolometric correction. In addition, \\citet{cas1} studied the impact of element diffusion in detail. The effects of overshooting from the convective envelope have been considered by \\citet{alongi}.  This paper presents a comprehensive study of \\vbump\\, incorporating all the relevant uncertainties to firmly establish the agreement between standard stellar models and observations of the red giant bump.  The magnitude of the RGB bump depends upon critically upon the maximum depth of the convection zone.  In contrast, the enhancement in the observed number of stars in the bump depends upon the size of the chemical discontinuity. \\citet{bono} introduced the \\rbump\\ parameter, which measures the relative number of stars in the RGB bump. \\citet{bono} found that \\rbump\\  was a robust prediction of stellar evolution theory, as changes in the opacity, equation of state, and nuclear cross sections changed the predicted value of \\rbump\\ by a few percent.  Their work found fair agreement between the observed \\rbump\\ values and those predicted by standard stellar evolution theory.  More recently \\citet{riello} used HST data to determine \\rbump\\ in 54 Galactic globular clusters.  They found this data was in good agreement with standard stellar evolution models.  This paper presents a comprehensive analysis of how uncertainties in the inputs into stellar evolution theory effect the predictions for \\vbump\\ and \\rbump.  Section \\ref{method} provides an overall of the Monte Carlo approach used in the paper, while \\S \\ref{uncertain} presents a detailed analysis of the uncertainties in standard stellar evolution theory.  The theoretical luminosity functions are presented in \\S \\ref{theolf}.  Section \\ref{rgbb} discusses the \\vbump\\ results, while the \\rbump\\ results are presented in \\S \\ref{rbmp}.  A summary of the main results is given in \\S \\ref{summ}. ",
        "conclusions": "}  We use a large Monte Carlo simulation to investigate theoretical uncertainties in the RGB luminosity function, specifically in the $V$-band magnitude of the red giant bump and excess number of stars in the bump.  We find excellent agreement between theoretical and observational values for the bump magnitude, with the uncertainty in theoretical values comparable to the scatter in observational values. Metallicity has a very significant effect on the bump magnitude, while the cluster age impacts the bump magnitude at a level lower than the theoretical uncertainties. The most important sources of uncertainty in the predication of the bump magnitude are the alpha-capture overabundance [$\\alpha$/Fe], the mixing length, low-temperature opacities, and the treatment of surface boundary conditions.  Theoretical uncertainties in stellar evolution models have little impact on the excess number of stars in the bump region. This is a robust prediction of standard stellar evolution is found to be in reasonable agreement with observations."
    },
    "0512/astro-ph0512022_arXiv.txt": {
        "abstract": "The X-ray holes at the centre of the Perseus Cluster of galaxies are not all at the same position angle with respect to the centre of the cluster. This configuration would result if the jet inflating the bubbles is precessing, or moving around, and the bubbles detach at different times.  The orientations which best fit the observed travel directions are an inclination of the precession axis to the line of sight of $120^{\\circ}$ and an opening angle of $50^{\\circ}$. From the timescales for the bubbles seen in the cluster, the precession timescale, $\\tau_{\\rm prec}$, is around $3.3 \\times 10^7 \\yr$.  The bubbles rising up through different parts of the cluster may have interacted with the central cool gas, forming the whorl of cool gas observed in the temperature structure of the cluster.  The dynamics of bubbles rising in fluids is discussed.  The conditions present in the cluster are such that oscillatory motion, observed for bubbles rising in fluids on Earth, should take place. However the timescale for this motion is longer than that taken for the bubbles to evolve into spherical cap bubbles, which do not undergo a path instability, so such motion is not expected to occur. ",
        "introduction": "High resolution radio observations of a number of extragalactic radio sources show that the opposed jets contain inversion symmetries, which provides evidence that the central engine is precessing, e.g. NGC 326 \\citep{Ekers1978} and NGC 315 \\citep{Bridle1979}.  In a review of quasars and radio galaxies which exhibit bent jets, of which some are inversion-symmetric about the core, \\citet{Gower_1982} present ten sources whose morphology can be explained using a precessing jet with periods ranging between $10^4\\to10^7\\yr$.  The extraordinary galactic object SS433 \\citep{HjellmingJohnston81} is successfully described with a precessing jet model by \\citet{Margon84}, and has been the subject of a recent study by \\citet{Blundell}.   The radio source at the centre of the Perseus Cluster, 3C84, is an extended source and shows an ``{\\sf S}'' or ``{\\sf Z}''-shaped morphology at $\\ghz$ frequencies.  These lobes have been observed to anti-correlate spectacularly with the observed X-ray emission from the cluster \\citep{Bohringer,ACF_complex_PER00}.  These depressions in the X-ray emission are thought to be the result of the jet from the central engine blowing bubbles of relativistic plasma into the thermal Intra-Cluster Medium, pushing the ICM aside \\citep{GullNorthover,Bohringer}.  When their buoyancy velocity is greater than their expansion velocity, the bubbles are expected to detach and rise up buoyantly through the ICM \\citep{Churazov00}.  This is observed in the Perseus Cluster, to the North-West and to the South of the core where there are depressions in the X-ray emission which do not have any associated $\\ghz$ radio emission; however, at lower frequencies (e.g. $330\\mhz$) ``spurs'' of emission are observed pointing towards these so-called ``Ghost bubbles'' \\citep{Celotti02}.  \\begin{figure*} \\includegraphics[width=0.95 \\columnwidth]{X_ray.ps} \\includegraphics[width=0.95 \\columnwidth]{Bubbles.ps} \\caption{\\label{Xray_radio_image} {\\scshape left:} The X-ray emission from the centre of the Perseus Cluster from the 200 ks observation of \\citet{ACF_deep_PER03}, corrected and accumulatively smoothed with a $\\sigma=15$. {\\scshape right:} A $330 \\mhz$ radio map of 3C84.  The shapes used for the different bubbles and the lines joining their centres to the radio nucleus are shown along with their names.} \\end{figure*}  The two pairs of bubbles which are clearly visible in the X-ray emission from the cluster belong to two different ``generations,'' the young, radio-active pair and the older ``ghost'' pair.  The bubbles are observed at different position angles with respect to the radio core in the centre of the cluster.  In Section \\ref{Radioemission} we discuss the morphology of the radio emission at the centre of the cluster, linking in other features in Section \\ref{otherfeatures}.  Section \\ref{model} outlines the model for the precessing jet, the resultant parameters of which are presented in Section \\ref{parameters}.  We calculate the timescales for the precession in Section \\ref{timescales} and discuss the model in Section \\ref{discussion}.  We put forward two candidates for causing the precession in Section \\ref{causeprec} and discuss the dynamics of rising bubbles in Section \\ref{daVinci}. ",
        "conclusions": "The observations of the bubbles in the Perseus Cluster show that they are at different position angles with respect to the cluster centre.  We have outlined several interpretations for this behaviour, including a jet which is precessing or moving around.  For the precessing jet we find that the best fit precession parameters are that the projected precession angle offset is $10^{\\circ}$ westwards of North; the precession axis is tilted $120^{\\circ}$ away from the line of sight and the precession opening angle is $50^{\\circ}$.  Using the timescales for the bubbles the precession timescale is around $3.3 \\times 10^7 \\yr$.  We review some work on the motion on bubbles rising in fluids and the onset of path instability.  The conditions of the motion in the cluster are such that the oscillations about the direction of motion are possible for the young bubbles, but the timescales are of order $1\\gyr$, by which time the bubbles have evolved into spherical caps, which do not exhibit this type of motion.  A steady precession or motion of the jet may be responsible for the whorl seen in the temperature structure. Alternatively, we note that the whorl may be an artifact of a partial inner circle and a partial outer ellipse, connected together in the SE. The structures at the centre of the Perseus cluster provide an enigmatic view of the past history of this region."
    },
    "0512/astro-ph0512091_arXiv.txt": {
        "abstract": "{This paper presents a consistent description of the formation and the subsequent evolution of gaseous planets, with special attention to short-period, low-mass hot-Neptune planets characteristic of  $\\mu$ Ara-like systems. We show that core accretion including migration and disk evolution and subsequent evolution taking into account irradiation and evaporation provide a viable formation mechanism for this type of strongly irradiated light planets. At an orbital distance $a \\, \\simeq$ 0.1 AU, this revised core accretion model leads to the formation of planets with total masses ranging from $\\sim$ 14 $\\mearth$ (0.044 $\\mjup$) to $\\sim$ 400 $\\mearth$ (1.25 $\\mjup$). The newly born planets have a dense core of $\\sim$ 6 $\\mearth$, independent of the total mass, and heavy element enrichments in the envelope,  $M_{\\rm Z,env}/M_{\\rm env} $, varying from 10\\% to 80\\% from the largest to the smallest planets. We examine the dependence of the evolution of the born planet on the evaporation rate due to the incident XUV stellar flux. In order to reach a $\\mu$ Ara-like mass ($\\sim$ 14 $\\mearth$) after $\\sim $ 1 Gyr, the initial planet mass must range from 166 $\\mearth$ ($\\sim$ 0.52 $\\mjup$) to about 20 $\\mearth$, for evaporation rates varying by 2 orders of magnitude, corresponding to 90\\% to 20\\% mass loss during evolution. The presence of a core and heavy elements in the envelope affects appreciably the structure and the evolution of the planet and yields $\\sim 8\\%-9\\%$ difference in radius compared to coreless objects of solar composition for Saturn-mass planets. These combinations of evaporation rates and internal compositions translate into different detection probabilities, and thus different statistical distributions for hot-Neptunes and hot-Jupiters. These calculations provide an observable diagnostic, namely a mass-radius-age relationship to distinguish between the present core-accretion-evaporation model and the alternative colliding core scenario for the formation of hot-Neptunes. ",
        "introduction": "Since the discovery of the first exoplanet by Mayor \\& Queloz (1995), planet hunters have been discovering planets with smaller and smaller masses. With the discovery of objects in the range of 10-20 $\\mearth$, i.e. 0.03-0.06 $\\mjup$ \\footnote{1 $\\mjup = 318 \\mearth$.} (Butler et al. 2004; McArthur et al. 2004; Santos et al. 2004), a new step in the quest for Earth-like planets has been taken. While representing an extraordinary achievement from the observational side, theorists still struggle to understand the structure and the origin of these light giant planets. Are they essentially composed of ices and rocks, with possibly a thin atmosphere, like our ice giants Uranus and Neptune? Or do they originate from larger gaseous planets, with a large gaseous envelope and a relatively small central rocky core? The answer to these questions requires an understanding of their formation process.  Current planet formation scenarios, based either on the core accretion model or on gravitational instability, can more or less explain the presence of relatively massive planets with masses $\\simgr \\, 100  \\, \\mearth$ (or even larger in the case of the gravitational instability scenario) at various orbital separations but they do not necessarily predict the formation of a large number of lighter planets (Boss 2001; Ida \\& Lin 2005; Papaloizou \\& Nelson 2005).  In the framework of the core accretion model (Pollack et al. 1996), the general expectation was to find preferentially planets either less massive or more massive than the newly discovered Neptune-like planets, for these latter lie within the domain of critical mass ($\\sim 10-20 \\mearth$) above which runaway accretion of gas begins. Similarly, Ida \\& Lin (2004) suggest a possible deficit of intermediate mass planets ($\\sim 10-100 \\mearth$, i.e. $\\sim 0.03-0.3 \\mjup$) with orbital separation $a<3$ AU. On the opposite, using N-body simulations of colliding cores in a protoplanetary disk, Brunini \\& Cionco (2005) find that Neptune-like planets close to their host-star can form easily as a by-product of planetary formation. If their scenario is correct, these planets should be of ice-rock composition with only a thin atmosphere. They also predict that a large population of these \"hot cores\" should be discovered in a near future.  As an alternative to this colliding core scenario, and given the close orbital distance of the Neptune-like planets  discovered up to now, one cannot  exclude the possibility that they formed initially as larger giant planets wich have undergone atmospheric evaporation during their lifetime (Baraffe et al. 2005). Within this picture, formation and evolution are strongly correlated: a correct understanding of light planet properties thus requires a consistent description of the planet formation and evolution in order to interpret present-day observations.  This paper is a first attempt to derive such a consistent picture from the planet formation to its subsequent evolution. We apply the core accretion models (cf. \\S 2) developed recently by Alibert et al. (2005a) to the 14 $\\mearth$ (0.044 $\\mjup$) planet (modulo sin $i$) orbiting around the G-star $\\mu$ Ara at an orbital distance $a=0.09$ AU (Santos et al. 2004). We focus on $\\mu$ Ara-like planets  because it is still reasonable to apply the Alibert et al. (2005a) formation model at this orbital distance, whereas it is no longer the case at the location of the two other Neptune-like planets which are located at $a \\, < \\, 0.05$ AU from their parent star (Butler et al. 2004; McArthur et al. 2004). Indeed, for $a \\simle 0.1$ AU, the description of the inner part of the disk, including tidal and magnetic interactions with the star, is too crude to provide a reliable formation scenario.   Adopting the main characteristics of the newly born planets as predicted by the Alibert et al. (2005a) formation model (core mass, heavy element content), we follow the later evolution  of these planets, according to Baraffe et al. (2004, 2005), taking into account irradiation and evaporation effects due to the vicinity of the parent star. We examine the sensitivity of the results on the evaporation rate by exploring a range of different rates and present our results in \\S 4. Predictions and uncertainties of our scenario are discussed in \\S 5. ",
        "conclusions": "The present study presents first consistent calculations between planet formation, within an improved version of the core accretion model, and subsequent evolution, including irradiation and evaporation. These calculations can explain, among others, the existence of hot-Neptune planets, like the one recently discovered around $\\mu$ Ara. The formation model constructed by Alibert et al. (2005a) allows the existence of planets with initial mass ranging from $\\sim$ 14 $\\mearth$ (0.044 $\\mjup$) to $\\sim$ 400 $\\mearth$ (1.25 $\\mjup$). The planets end up having a dense core of $\\sim$ 6 $\\mearth$, independently of the total mass, and heavy element enrichment in the envelope varying from 10\\% to 80\\% from the largest to the smallest mass. With maximal evaporation rates,  as predicted by the models of L03, the progenitor of gigayear old $\\mu$ Ara-like  planets has an initial mass $\\sim$ 0.5 $\\mjup$. For evaporation rates a factor 10 to 20 smaller, the progenitor has a mass in the range $\\sim 0.1$-0.15 $\\mjup$ ($\\sim$30-50 $\\mearth$), meaning the planet has lost 1/2 to 2/3 of its initial envelope. For planets already forming in the Neptune-mass range, $\\lesssim 20\\,\\mearth$, the L03 rate must be reduced by a factor 1/100 for the planets to survive to evaporation at 0.09 AU from their Sun. In this latter case, the evaporation rates vary from $\\sim$ 10$^{11}$ g/s (6 $10^{-10}$ $\\mearth$/yr), at an age of 1 Gyr, to $\\sim$ 7 $10^{9}$ g/s (4 $10^{-11}$ $\\mearth$/yr) at 10 Gyr. Our current knowledge of evaporation processes  in exoplanet atmospheres is still too embryonic to favor or exclude any of these values, although recent hydrodynamical simulations favor the intermediate value (Tian et al. 2005). The range of escape rates explored in the present calculations, however, should bracket the \"real\" solution.  Our calculations provide an {\\it observable diagnostic} for such a core-accretion-irradiation scenario for the formation and evolution of hot-Neptunes, namely radii $\\simgr$  0.6 $\\rjup$ for Neptune-mass planets. Also, the scenario with maximal evaporation rate may provide a statistical signature in terms of the number of hot-Neptunes, which is expected to be much smaller than the number of hot-Jupiter-like planets. In the absence of detailed statistical analysis, a crude comparison of evolutionary timescales suggests 10 times  less hot-Neptunes than hot-Jupiters at  $a \\simeq 0.09$ AU, {\\it for the L03 maximal escape rate}. The detection of a larger fraction of hot-Neptunes at such orbital separation would thus favor scenarios with lower evaporation rates.  A complete understanding of the formation and evolution of light, Neptune-mass planets is still far from reach. The present attempt to derive a {\\it consistent picture} from the planet genesis to its present day conditions provides one possible alternative path for the birth and fate of these objects. It also points to several issues to be addressed for a better modelling of planet formation and interior and atmospheric structures. Further improvement requires in particular (i) a better treatment of heavy element enrichment, both in the envelope and in the atmosphere and (ii) a better understanding of evaporation processes.  In parallel with such theoretical developments, observations of planets down to a few Earth-masses and located at orbital separations large enough ($a \\simgr 1 $ AU) to remain unaffected by evaporation would provide strong constraints on the formation scenarios of light planets. This is another appealing challenge for planet hunters."
    },
    "0512/astro-ph0512572_arXiv.txt": {
        "abstract": "{} {To investigate the spectroscopic properties of a selected optical photospheric spectra of core collapse supernovae (CCSNe). Special attention is devoted to traces of hydrogen at early phases. The impact on the physics and nature of their progenitors is emphasized.} {The CCSNe-sample spectra are analyzed with the parameterized supernova synthetic spectrum code ``SYNOW'' adopting some simplifying approximations.} {The generated spectra are found to match the observed ones reasonably well, including a list of only 23 candidate ions. Guided by SN Ib 1990I, the observed trough near 6300\\AA~is attributed to H$\\alpha$ in almost all Type Ib events, although in some objects it becomes too weak to be discernible, especially at later phases. Alternative line identifications are discussed. Differences in the way hydrogen manifests its presence within CCSNe are highlighted. In Type Ib SNe, the H$\\alpha$ contrast velocity (i.e. line velocity minus the photospheric velocity) seems to increase with time at early epochs, reaching values as high as 8000 km s$^{-1}$ around $15-20$ days after maximum and then remains almost constant. The derived photospheric velocities, indicate a lower velocity for Type II SNe 1987A and 1999em as compared to SN Ic 1994I and SN IIb 1993J, while Type Ib events display a somewhat larger variation. The scatter, around day 20, is measured to be $\\sim$5000 km s$^{-1}$. Following two simple approaches, rough estimates of ejecta and hydrogen masses are given. A mass of hydrogen of approximately  0.02 $M_\\odot$ is obtained for SN 1990I, while SNe 1983N and 2000H ejected $\\sim$0.008 $M_\\odot$ and $\\sim$0.08 $M_\\odot$ of hydrogen, respectively. SN 1993J has a higher hydrogen mass, $\\sim 0.7$ $M_\\odot$ with a large uncertainty. A low mass and thin hydrogen layer with very high ejection velocities above the helium shell, is thus the most likely scenario for Type Ib SNe. Some interesting and curious issues relating to oxygen lines suggest future investigations.} {} ",
        "introduction": "Stripped-envelope SNe, namely Type Ib (helium-rich) and Type Ic (helium-poor), being hydrogen deficient objects, are undoubtedly amongst the most mysterious  SN classes. Recently efforts have started to understand the nature of these objects through studies of samples (Matheson et al. 2001\\nocite{Math01}; Branch et al. 2002\\nocite{Br02}).  However the rarity of cases with well sampled observations, photometry and spectra, hampers a more direct inference of the physical situation behind the explosions and hence a clear view of the progenitor nature. Nevertheless, the discovery of metamorphosing events as SN 1987K and SN 1993J (recognized as SNe IIb), that evolve from Type II to Type Ib-c as they age, together with the similarity of the environments in which they occur have linked Type Ib-c SNe to a core-collapse scenario in massive stars. At present, indeed, the most widely accepted models relate Type Ib-c SNe to both relatively low mass progenitors within the context of close binary system evolution (i.e. mass-loss as consequence of mass transfer) and massive stars that have undergone significant mass-loss due to a wind (i.e. Wolf-Rayet stars). So far observations have not discriminated between the two scenarios, although recently, using HST data, a high spatial resolution search for the progenitor of the Type Ic SN 2004gt in a wide wavelength range from the far UV to the near IR has suggested that the event might result from an evolved Wolf-Rayet star, although the observations could not constrain models invoking less massive progenitors in binary systems (Gal-Yam et al. 2005).  Photometrically, the  lack of significant hydrogen in the outer layers of SNe~Ib-c probably inhibits the most important characteristic of Type II SNe light curves, namely the plateau phase resulting from the hydrogen recombination wave.  At late phases, the steeper decline rate, compared to the ``$^{56}$Co to $^{56}$Fe'' decay slope, is indicative of significant $\\gamma$-ray escape as a result of the low mass ejecta in this class of objects (Clocchiatti $\\&$ Wheeler 1997\\nocite{Cloc97}); there are rare exceptions where the late slope approaches the full trapping rate (e.g. SN Ib 1984L; Schlegel $\\&$ Kirshner 1989\\nocite{Schl89}).  Spectroscopically, a clear separation scheme within the stripped-envelope SNe subclasses is still lacking. Part of the problem is the absence of meaningful statistics. A direct classification was earlier proposed by Harkness et al. (1987)\\nocite{Har87} on the basis of He I strengths in the photospheric optical spectra of Type Ib SNe. Wheeler et al. (1994\\nocite{Whee94}), however, have claimed the presence  of He I 10830\\AA~ in SN Ic 1990W, although He I lines were not noted in the optical region. The authors presented the idea of adopting, instead, the OI 7773\\AA~ line as a distinguishing feature. The absorption seems stronger in Type Ic than in Ib SNe. Matheson et al. (2001)\\nocite{Math01} came to the same conclusion when analyzing a sample of Ib and Ic events. In addition, it has been argued that in Type Ib SNe He I lines 5876\\AA~ and 7065\\AA~ gradually grow in strength with respect to the one at 6678\\AA~(Matheson et al. 2001)\\nocite{Math01}. SN 1998bw (GRB980425) is another example of classification as Type Ic, but where IR lines of He I have been clearly identified (Patat et al. 2001)\\nocite{Pat01}.  Branch et al.(2002)\\nocite{Br02} have presented a relation between the velocity at the photosphere, measured using synthetic-spectrum analysis, and the time since maximum light for a sample of Type Ib SNe. The relation is well fitted by a power-law, indicating a fairly homogeneous behaviour. Interestingly, SN Ib 1990I does not follow this trend (Elmhamdi et al. 2004)\\nocite{Elm04}.  An important issue concerns features in the 6000$-$6500\\AA~region of the early spectra in Ib-c SNe. Deng et al. (2000)\\nocite{Den00}, when analyzing spectra of SN Ib 1999dn, have attributed the absorption feature seen at 6300\\AA~ around maximum brightness to H$\\alpha$  which later on disappears or is overwhelmed by C II 6580\\AA~when H$\\alpha$ optical depth decreases as a consequence of the envelope expansion. The possibility of the presence of C II 6580\\AA~in Ib spectra has been earlier suggested by Harkness et al.(1987)\\nocite{Har87}. Alternatively, the 6300\\AA~absorption was identified to be due to Ne I 6402\\AA~in SN Ib 1991D (Benetti et al. 2002). Ne I lines and H$\\alpha$ had already been proposed to account for the deep absorption in SN Ib 1954A by Branch (1972)\\nocite{Br72}.  The presence, or not, of hydrogen and/or helium in Type Ib-c, with the possibility of quantifying the amount, is of great importance in identifying the progenitor stars that may give rise to these classes of objects.  The papers by Matheson et al. (2001)\\nocite{Math01} and Branch et al. (2002)\\nocite{Br02} have provided an impetus towards an advanced understanding of SNe Ib and Ic. The present work may be regarded as a continuation of those efforts. However in our comparative study of early spectra of Type Ib SNe we include representatives of all the various types of CCSNe, namely Type IIb, Type Ic and Type II. This is established by means of synthetic spectra generated with the parametrized SN synthetic-spectrum code ``SYNOW''.  Our main goal was to understand the similarities and differences among SNe Ib objects in the available sample, and to also compare to properties of the wider CCSNe family. The paper is organized as follows. First the analysis method and parameters are illustrated in Sect. 2. Data description is briefly given and the best fit synthetic spectra are presented and compared with the observed ones in Sect. 3. This is done separately for each individual object among  representatives of CCSNe classes. Sect. 4 presents two methods to obtain spectroscopic mass estimates. The complete results will be discussed in detail and conclusions will be drawn in Section 5. ",
        "conclusions": "One of our main goals was to identify traces of hydrogen in Type Ib-c, and to identify any systematic similarities and differences correlating with other physical properties. Consequently, guided by modeling early spectra of SN Ib 1990I, we have explored different combinations of ions and shaping the $6000-6500$\\AA~wavelength range. Special attention has been devoted to the feature seen near 6300\\AA. Even though the number of events with spectra well sampled in wavelength, is limited, it appears that hydrogen in varying amounts is identifiable in most Type Ib objects especially at very early times. Only in two cases, namely SNe 1991D and 1991L, does the Ne I 6402\\AA~line remain as an alternative possibility with large departure from LTE. Even in SN 1996aq, that we classified as a transition Type Ib/c event, and SN 1999ex the presence of the H$\\alpha$ trough seems highly preferred. \\begin{figure} \\includegraphics[height=9cm,width=9cm]{4366fg34.ps} \\caption{Comparison of SN 1990I with SNe 1984L and 2000H, at two different phases of evolution. The vertical dashed lines indicate the locations of H$\\alpha$ and H$\\beta$ troughs in SN 2000H. He I optical series are marked by vertical ticks.} \\end{figure}  At later phases, more than $\\sim$ 3 weeks, SNe 1984L, 1988L, 1991ar, 1998dt, 1998T and 1999dn behave quite similarly, namely still showing evident He I lines and a ``$flat$'' $6000-6500$\\AA~region of the spectrum, with a weak H$\\alpha$ absorption feature. Even SN 1991D with its particularly low photospheric velocity and narrow lines belongs to this class. SN 1991L appears to be similar to SN 1991D in having shallow He I optical features. SNe 1999di and 2000H are the unique events among the sample that display obvious He I lines, together with a pronounced and deep H$\\alpha$ absorption line. It is interesting to note here the peculiar behaviour of SN 1997dc. The event has He I troughs as deep as in SNe 1999di and 2000H, however it has a flat $6000-6500$\\AA~region. SN 1990I at this phase shows distinct He I lines and a moderate H$\\alpha$ trough. Figure 34 compares SN 1990I with SNe 1984L and 2000H at two different phases of evolution. The vertical dashed lines indicate the locations of H$\\alpha$ and  H$\\beta$ troughs in SN 2000H. On the one hand, while at early phases (upper panel) the trough assigned to H$\\alpha$ is clear in all the three SNe, it disappears later on in SN 1984L (lower panel). On the other hand, H$\\beta$ appears clearly in SN 2000H and absent in SNe 1984L and 1990I. Moreover, the comparison emphasizes the high velocity behaviour of SN 1990I as is evident from the blueshifted features in SN 1990I compared to the others. He I optical lines are also shown. For the three events the He I troughs at 5876\\AA, and 7065\\AA~are clearly identified and get narrower with time. At early phase the He I 6678\\AA~is recognized in SNe 1990I and 1984L while it is just hinted in SN 2000H. Later on, Fig.34-bottom panel, the  He I 6678\\AA~line grows in strength in SNe 1984L and 2000H while it fades in SN 1990I. The comparison of these three events, in terms of changes in line visibility and velocity, demonstrates the complexity involved in demonstrating any thing like a continuous sequence.  \\begin{figure} \\includegraphics[height=10.5cm,width=9cm]{4366fg35.ps} \\caption{The photospheric velocity evolution of Ib SNe sample, compared to SNe 1987A, IIP 1999em and Ic 1994I (upper panel). Middle panel: the H$\\alpha$ contrast velocity ratio evolution (i.e. v$_{min}^{ratio}$(H$\\alpha$)$/$v$_{phot}$). Lower panel: the H$\\alpha$ contrast velocity evolution (i.e. v$_{min}$(H$\\alpha$)-v$_{phot}$). Results for SN IIP 1999em and SN IIb 1993J are also shown.} \\end{figure}  Hydrogen manifests its presence in a different way in the other CCSNe types. A part from SN 1996aq that we re-classify as a transient Ib/c event rather than a pure Type Ic SN, there is no evidence for the presence of hydrogen in Type Ic objects analyzed here (i.e. SNe 1987M and 1994I). The ``$hybrid$'' SN IIb 1993J displays typical Type II features at early phases, such as strong and broad H$\\alpha$ P-Cygni emission component, which is absent in Type Ib objects. This can be explained within the context of the ``detachment'' concept. In fact, H$\\alpha$ P-Cygni profile would lose its obvious emission component when it is highly detached. At early epochs, both SN IIb 1993J and SN II 1999em show clear evidence of the other H I Balmer lines since the corresponding optical depths are too large to allow them to be distinctly visible, contrary to normal Type Ib objects, with the exception of some cases with deep and conspicuous H$\\alpha$ troughs that present signature of H$\\beta$ too (e.g. SNe 1999di and 2000H). Two factors make H$\\beta$ barely discernible in Type Ib: the optical depth found to fit H$\\alpha$ is small and the contrast velocity of H$\\alpha$ is high. The opposite is seen in Type II and IIb  events. As time goes on, the H$\\alpha$ emission peak in SN 1993J changes to a double-peaked structure which we recognized as the emergence of the He I 6678\\AA, with more clearer He I optical lines. These laters disappear at the nebular phase and the spectrum is dominated by typical Ib features such as [O I], [Ca II] and Ca II in emission. The situation in SN~1999em, and in Type II  in general, is different. The He I lines are found to disappear very early, about one weak after explosion, and even at later nebular epochs the H$\\alpha$ is still the most prominent feature in the spectrum.  In Figure 35, upper panel, we report the resulting photospheric velocities from our best fits for the CCSNe sample. Data for SN 1987A, corresponding to Fe II 5018\\AA~absorption, are also shown for comparison (dashed line;  Phillips et al. 1988\\nocite{Phil88}). Additional points for SN 1994I are taken from Millard et al. 1999. The plot indicates the low velocity behaviour of Type II SNe, both 1987A and 1999em, at early as well as intermediate epochs. SN IIb 1993J follows somewhat similar behaviour as SN Ic 1994I in having higher velocities. As far as Type Ib SNe are concerned, they appear to display a different velocity evolution. The scatter seems to increase at intermediate  phases (around 20$-$30 days). This fact can be simply due to the paucity of available observations outside that range. Around day 20, for example, a scatter as high as 5000 km s$^{-1}$ is measured. SNe 1990I and 1998dt belong to a class with the higher ``$v_{phot}$'', while objects such as SNe 1991D and 1996aq have the lowest estimated velocities, approaching even Type II objects. The remainder of the Ib events follow a similar trend, namely the one described by Branch et al. (2002). SN 1999ex is found to belong to this class. Surely more data are needed to obtain meaningful statitical conclusions. \\begin{figure} \\includegraphics[height=9cm,width=9cm]{4366fg36.ps} \\caption{ The H$\\alpha$ velocity contrast evolution, $v_{cont}^{ratio}$(H$\\alpha$), versus H$\\alpha$ optical depth. Results for SNe IIP 1999em and IIb 1993J are displayed for comparison. Note that results are displayed independently of the phase.} \\end{figure}  The middle and bottom panels in Fig. 35 display the evolution the of H$\\alpha$ contrast velocities, ``$v_{cont}^{ratio}$'' and ``$v_{cont}$'' respectively,  for the sample. As discussed before, when discussing differences in the H$\\alpha$ profile in CCSNe, we would expect an increasing value of ``$v_{cont}$(H$\\alpha$)''  going from Type II to IIb to Ib SNe. This trend is in fact illustrated in the plot (bottom panel). A similar trend is seen in the ``$v_{cont}^{ratio}$(H$\\alpha$)'' evolution as indicated by the middle panel. While in SN IIb 1993J and SN~IIP~1999em the line is found to be either undetached or slightly detached, it is highly detached in Type Ib events. Moreover, the ``$v_{cont}$(H$\\alpha$)'' is found to increase within the first 15 days, reaching values as high as 8000  km s$^{-1}$, and then follows an almost constant evolution. According to Table 2 and up to $\\sim$60 days, Type Ib SNe have hydrogen down to 11000-12000 km s$^{-1}$, while in SN 1993J hydrogen is down to 8000 km s$^{-1}$. SNe II appear to have hydrogen down to even lower velocities ($\\sim$5000 km s$^{-1}$ in SN 1999em). In addition, hydrogen in Type Ib SNe is found to have very small optical depths independently of the contrast velocity (also independently of the phase). This is shown by Figure 36. In fact, Type Ib objects populate the region of extremely low optical depths, although the ``$v_{cont}^{ratio}$(H$\\alpha$)'' spans a large range, contrary to SNe 1993J and 1999em. \\begin{figure} \\includegraphics[height=9cm,width=9cm]{4366fg37.ps} \\caption{ The optical depth of the O I 7773\\AA~line according to the best fit spectra. Data for different SN types are reported for comparison.} \\end{figure}  Table 2 also displays properties of the helium line. In Type Ib SNe, He I lines are found to be not always detached as is the case for hydrogen. The He I 5876\\AA~line is found to be clearly distinguishable in Type Ib objects. A contribution from Na ID is argued for in some cases, such as the broad feature in SN 1996aq and cases with an observed bump shortward of the minimum absorption as seen in SNe 1990B and 1997dc. It is important to recall here that in some cases we obtained a better fit to the He I optical lines adopting an $e-$folding SSp, which allows a two-component treatment of the line and a gradually decreasing optical depth below the detachment velocity, rather than a discontinuity in it. In the case of SN 1999ex, we extend our He I line fits beyond 1 micron, confirming the presence of various IR lines, namely at 1.083, 1.284, 1.700(?), and 2.085 microns (Fig. 20), in accord with optical He I parameters. In our analysis, we also pointed to some interesting line identifications.  For example the contribution of Sc II lines in nearly all the CCSNe spectra, especially for reproducing the double absorption features blueward of the He I 5876\\AA~line. We note here that except for He I and H I, almost all the lines in the CCSNe synthetic spectra are undetached.  Another important study is the behaviour of the O I 7773\\AA~line. Figure 38 reports the resulting optical depths corresponding to our best synthetic spectra fits. It is not easy to draw clear-cut conclusions from this figure since one needs, for instance, to populate the figure with more Type Ic objects. However, at intermediate phases, it seems that Type Ib objects tend to concentrate in the low optical depth region, while SN Ic 1987M is found to display the deepest profile. SNe 1993J and 1999em, at similar phases, are the objects with the lowest O I 7773\\AA~optical depth. At somewhat later epochs, transient Type Ib/c objects display deep O I 7773\\AA~troughs. The stronger and deeper permitted oxygen lines at early phases of SNe Ic and Ib/c spectra might imply that they are less diluted by the presence of a helium envelope. Indeed one might expect oxygen lines to be more prominent for a ``naked'' C/O progenitor core. Despite the paucity of well sampled CCSNe observations, two observational aspects tend to reinforce this belief: $\\bf{First}$, the forbidden lines, especially [OI]6300,~6364\\AA, seem to appear earlier following a SNe sequence ``Ic$-$Ib$-$IIb$-$II''. In fact the oxygen line emerges at an age of 1-2 months in Type Ic SN 1987M (Filippenko 1997)\\nocite{Fil97}. SN Ic 1994I displayed evidence for the line at an age of 50 days, although some hints may even be seen in the $\\sim$36day spectrum (Clocchiatti et al. 1996b). While in SN Ib 1990I it was hinted at the 70day spectrum (Elmhamdi et al. 2004). In other Type Ib SNe it appears earlier than in SN 1990I. In SN IIb 1993J, a transition object, the line was visible in the 62day spectrum (Barbon et al. 1995)\\nocite{Barb95}. SN 1996cb, another well observed IIb event, showed evidence of the [OI]6300,~6364\\AA~line around day 80 (Qiu et al. 1999). In SNe II, however, the line appears later: around day 150 in SN 1987A (Catchpole et al. 1988)\\nocite{Cat88} and after day 138 in SN 1992H (Clocchiatti et al. 1996a). In SN II 1999em it is suggested at a somewhat earlier phase compared to SNe 1987A and 1992H, namely at day 114. Whether all this can be understood in terms of the lower progenitor mass of SN 1999em and its presumed lower oxygen mass remains to be investigated (Elmhamdi et al. 2003). $\\bf{Second}$, it seems that the nebular emission line decreases in breadth following the SNe sequence above. Of course much work is still needed in this respect based on larger CCSNe samples at both photospheric and nebular phases, relating them to the photometry of the CCSNe variety (Elmhamdi $\\&$ Danziger, in preparation). Furthermore, one important point is to check for any possible correlations between the oxygen minimum velocity (at early epochs), their line widths, and to relate this to a progenitor mass indicator [Ca II]/[O I] ratio (at nebular phases; Fransson $\\&$ Chevalier 1989), the progenitor properties (i.e. masses, energies). Other factors such as mixing, variation in envelope densities and metallicity may play an additional complicating role.  Finally we presented two methods to determine the ejecta and hydrogen mass in CCSNe and especially in Type Ib events. Although the methods are very approximate\\footnote{``Method 2'' is a direct method, while ``Method 1'' is undirect and gives upper limits on the hydrogen mass.}, our results do not conflict with more detailed estimates, those based on hydrodynamical and NLTE models. A thin layer of hydrogen, ejected at high velocities down to 11000-12000 km s$^{-1}$, appears to be present in almost all the Type Ib events studied here. These results suggest possible directions for further more sophisticated work. The necessity to introduce lines of Ne I, Sc II, Ba II, Ca I, Fe II in some cases but not in others must inevitably raise concerns about the identifications and other reasons for these variations."
    },
    "0512/astro-ph0512602_arXiv.txt": {
        "abstract": "We perform a statistical analysis of the temporal structure of long Gamma-Ray-Bursts (GRBs). First we consider a sample of bursts in which a long quiescent time is present. Comparing the pre-quiescence with the post-quiescence emission we show that they display similar temporal structures, hardness ratios and emitted powers, but, on the average, the post-quiescence emission is roughly twice as long as the pre-quiescence emission. We then consider a sample of long and bright GRBs. We show that the duration of each emission period is compatible with the duration of an active period computed in various inner engine models. At the contrary, if the inner engine is assumed to be always active, i.e. also during the quiescent times, in several cases the total duration of the burst largely exceeds the theoretical durations. Our analysis therefore does not support the interpretation of long quiescent times in terms of stochastic modulation of a continuous wind. Instead the quiescent times can be interpreted as dormancy periods of the inner engine. Before and after a dormancy period the inner engine produces similar emissions. ",
        "introduction": "The time structure of GRBs is usually complex and it often displays several short pulses separated by time intervals lasting from fractions of second to several ten of seconds. The analysis of the light curves can provide hints on the activity of the inner engine although the relation between the observed signal and the Physics of the inner engine is not yet completely understood.  A previous statistical analysis \\citep{NP} has shown that there are three time-scales in the GRB light curves: the shortest one is the variability scale determining the pulses' durations and the intervals between pulses; the largest one describes the total duration of the bursts and, finally, an intermediate time scale is associated with long periods within the bursts having no activity, the so called {\\it quiescent times} \\footnote{QTs are not the same as the gaps between the precursors and the main pulses. In the case of the precursors, the gaps are between a softer and weaker component (the precursor) and a harder and much stronger one (the main burst). Here we observe gaps between pulses with the same characteristics.}. The origin of these periods of quiescence is still unclear.   Here we show, through a statistical analysis, that if a quiescent time longer than a few ten of seconds is present in the light curve then the pre-quiescence and the post-quiescence emissions (PreQE, PostQE) have similar variability scales but, on the average, the PostQE is longer and only marginally softer than the PreQE. The similarities between the first and the second emission periods strongly suggest that both emissions are produced by the same mechanism. Moreover we will show that the average durations of PreQE and of PostQE, separately, are compatible with the theoretical durations predicted by various inner enngine models. ",
        "conclusions": "In this paper we have studied the problem of quiescent times in long GRBs borrowing the technique developed by Nakar and Piran. The main results obtained in our analysis are the following:  \\noindent --- the cumulative distribution of time intervals between peaks within a same active period is well fitted by a lognormal distribution, while long quiescent times are well fitted by a powerlaw distribution. This strengthens \\citet{NP} conclusion about the different origin of long quiescent times respect to shorter time intervals;  \\noindent --- the two components of the emission, preceding and following the quiescent time are statistically similar from the viewpoint of their temporal microstructure, their hardness ratio and emitted power. These results suggest a unique mechanism at the origin of both the pre-quiescence and the post-quiescence emission. Interestingly, the second emission lasts on the average twice the first, what provides an important constraint to the inner engine models;  \\noindent --- the durations of the activity periods of the inner engine, computed within the existing theoretical models do not exceed few ten seconds. These theoretical estimates compare favorably with the durations of the pre-quiescent and of the post-quiescent emissions, separately. On the other hand, in several bursts the total duration including the quiescent time, is roughly three times longer. This is an indication in favor of a dormant inner engine respect to wind modulation models."
    },
    "0512/astro-ph0512434_arXiv.txt": {
        "abstract": "Experimental conclusions from air shower observations on cosmic-ray photons above $10^{19}$~eV are based on the comparison to detailed shower simulations. For the calculations, the photonuclear cross-section needs to be extrapolated over several orders of magnitude in energy. The uncertainty from the cross-section extrapolation translates into an uncertainty of the predicted shower features for primary photons and, thus, into uncertainties for a possible data interpretation. After briefly reviewing the current status of ultra-high energy photon studies, the impact of the uncertainty of the photonuclear cross-section for shower calculations is investigated. Estimates for the uncertainties in the main shower observables are provided. Photon discrimination is shown to be possible even for rapidly rising cross-sections. When photon-initiated showers are identified, it is argued that the sensitivity of photon shower observables to the photonuclear cross-section can in turn be exploited to constrain the cross-section at energies not accessible at colliders. ",
        "introduction": "\\label{sec-intro}  Photons around and above $10^{19}$~eV might provide a key to understanding the origin of cosmic rays. Substantial fluxes of these ultra-high energy (UHE) photons are predicted in top-down models of cosmic-ray origin. In addition, UHE photons are produced by the GZK process of resonant photoproduction of pions~\\cite{gzk}, in analogy to GZK neutrinos.  Experimentally, UHE photons can be discerned from nuclear primaries due to differences in the expected shower signatures. So far, no claim of a photon detection exists and upper limits to UHE photons were set. Any conclusions about UHE photons rely on the comparison of data to detailed simulations of photon-induced showers. Although these photon showers are dominated by electromagnetic interactions, a source of uncertainty is the photonuclear cross-section, which has to be extrapolated over several orders of magnitude in energy from laboratory data for calculating showers induced by UHE photons.  In this work, after giving an overview of the current status of UHE photon studies, both the impact of the uncertainty from the photonuclear cross-section to photon shower features and possible prospects for constraining the cross-section are investigated. The plan of the paper is as follows. In Section~\\ref{sec-cr}, scenarios for producing cosmic-ray photons above $10^{19}$~eV = 10~EeV are briefly described. Features of showers induced by photons are discussed in Section~\\ref{sec-eas}. In Section~\\ref{sec-data}, the experimental situation concerning upper limits to photons is summarized. The role of the photonuclear cross-section is investigated in Section~\\ref{sec-sigma}. ",
        "conclusions": ""
    },
    "0512/hep-th0512120_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\hspace*{15pt} According to analyses of CMB[1]+SN-Ia[2]+HST[3]+LSS[4] data, there are strong evidences to indicate that our universe has recently entered a phase of accelerating expansion and that the universe is flat. It implies that if this is not a signal of modifying the standard theory of gravity, that may be a evidence of existing a \"exotic homogeneous matter\" with negative pressure termed \"dark energy\"(DE)[5]. An equation of state parameter (w$=p/\\rho$) is usually used to describe this energy component. The value of w is required to less that $-\\frac{1}{3}$ for a accelerating expansion universe. Due to the existence of ordinary matter and radiation, we actually need a more negative value of w to drive accelerating expansion of the universe. Analysis shows w can lie in the rang of $-1.32<$w$<-0.82$ with the 2-$\\sigma$ confidence levels[6]. Up to now a cosmological constant with w$=-1$ is in good agreement with all the data. However, the result doesn't rule out the scalar field and phantom field models as the dark energy candidate. For the $\\Lambda$CDM model[7], there are two boring problems: the fine-tuning problem and cosmic coincidence problem[8]. Though same problems remain in other models, it can be alleviated in some models[9]. \\par The scalar field appeared in cosmology is not the first time. As early as more than twenty years ago, theorist normally consider a homogenous scalar field $\\phi$ in an inflationary universe[10]. A scalar field $\\phi$ slowly evolving down its potential $V(\\phi)$ can drive both inflation and late-time accelerating expansion. The standard quintessence scenario[11](a canonical scalar field described by a lagrangian $L=\\frac{1}{2}{\\dot\\phi}^2-V(\\phi)$) is a simplest scalar field. The universes they predict will have significant differences with respect to different potentials. Another more complicated scalar field is K-essence. The idea of K-essence was firstly introduced as a possible model for inflation[12]. Later it was noted that K-essence was introduced as a possible models for dark energy[13]. K-essence can be defined as any scalar field with non canonical kinetic terms. Its lagrangian usually takes the form $L_A=V(\\phi)F(X)$. A more general form of lagrangian for K-essence[14] is $L_B=f(\\phi)g(X)-V(\\phi)$, where $\\phi$ is the scalar field, and $X=\\frac{1}{2}\\nabla_{\\mu}\\phi\\nabla^{\\mu}\\phi$. \\par The Born-Infeld[15] theory has been considered widely in string theory and cosmology[16]. There are two Born-Infeld type scalar field models. The first is rolling tachyon field[17] with the lagrangian form $L_{tach}=-V(\\phi)\\sqrt{1-{\\dot\\phi}^2}$, which can be classified as $L_A$. Its interesting features have been widely studied[18]. Anther one is the Nonlinear Born-Infeld(NLBI) scalar field theory with the lagrangian form $L_{NLBI}=\\frac{1}{\\eta}[1-\\sqrt{1-\\eta{\\dot\\phi}^2}]-V(\\phi)$(noted, when field velocity $\\dot\\phi \\rightarrow0$,$L_{NLBI}=L_{quin}=L=\\frac{1}{2}{\\dot\\phi}^2-V(\\phi)$ by Taylor expansion). Obviously it can be regarded as one form of $L_B$. W.Heisenberg proposed this Lagrangian density in order to describe the process of meson multiple production connected with strong field regime[19]. H.P.de Oliveira qualitatively studied the solutions of a three-dimensional dynamical system describing the static and spherically symmetric solutions for this NLBI scalar field [20]. The dark energy model with this NLBI scalar field was recently suggested by H.Q.Lu[21]. However, comparing with tachyon field, the role of NLBI scalar field in cosmology is far beyond study. \\par One direct thought motivated us to consider the NLBI field is, besides the attractive properties[19,22], being able to study the role of nonlinearities in the matter fields. Furthermore, we will be able to see whether it can provide more interesting physical results than those generated by ordinary quintessence model. \\par It is important to choose the potential $V(\\phi)$ for the scalar field. In many cases, we use a potential predicted by particle physics. However, we do not really know which theory of particle physics best describes the universe, we should keep an open mind as to the form of $V(\\phi)$. One ideal approach is to consider some simple and possible forms of potential, explore the cosmological behavior in detail using the qualitative theory of dynamical systems, and then study the evolution of universe in more detail to quantitatively fit with current observation data. We choose two simple potentials $-s\\phi$ and $\\frac{1}{2}m^2\\phi^2$ for further study in this paper. Constraints on the linear potential $-s\\phi$ in quintessence  and phantom models from recent supernova data have been argued in[23]. It has been argued[24] that such a potential is favored by anthropic principle considerations and can provide a potential solution to the cosmic coincidence problem. The square potential $\\frac{1}{2}m^2\\phi^2$ has been considered in a chaotic inflationary universe[25]. \\par In this paper, we will focus on the differences between NLBI scalar field and linear scalar field(quintessence), explore the cosmological scenario in detail and compare them with SN-Ia Gold data. This paper is organized as follows: In section 2, we simply review the theoretic model with the two scalar fields and derive their cosmological evolution with two different potentials. In section 3, we fit the Hubble parameter to the SN-Ia Gold data, obtain constraint for the potential parameters and compare the NLBI scalar field with the linear scalar field. Section 4 is conclusion and discussion. ",
        "conclusions": "Through the analyses of the two potentials in quintessence model and NLBI model, we can conclude that: \\par Firstly, the principle we choose the potential is that the theoretical prediction should be consistent with the observational universe(such as the observation of universe age, the CMBR measurement, the SN-Ia observation, the structure formation and so on). It is to say the predicted universe need to have nearly same \"history\"( to account for the observational data), but the fate of the predicted universe could have significant differences. \\par Secondly, for non-negative potentials, the common feature of the further universe is that they will continue expanding for ever, though the fate of the universes with different potentials may be dramatically altered. \\par Thirdly, if the potential can evolve into negative value, the universe will evolve continually from expansion($H>0$) to contraction($H<0$). This is not the case for positive potential(see fig1 and fig5). The cosmology with negative potential has been discussed in Ref[32]. \\par Finally, comparing our theoretic models with the observational data of SN-Ia, the age of universe $H_0t$, the equation of state w and the transition redshift $z$, we can conclude that the NLBI modes is consistent with all the observations. Furthermore the result shows that the NLBI model slightly excels quintessence model. However, the result also shows that a smaller value of $s\\&m$ (a slower rolling of field $\\phi$ correspondingly ) can provide better fits with observational data, but in this case the difference between NLBI scalar field and quintessence is not distinct. In order to get a more convincible result, the effect of NLBI scalar field in CMB (for instance the late-time ISW effect) should be studied carefully. \\hspace*{15 pt}"
    },
    "0512/astro-ph0512384_arXiv.txt": {
        "abstract": "% I summarize current knowledge of galaxy formation with emphasis on the initial conditions provided by the $\\Lambda$CDM cosmology, integral constraints from cosmological quantities, and the demographics of high-redshift protogalaxies. Tables are provided summarizing the number density, star formation rate and stellar mass per object, cosmic star formation rate and stellar mass densities, clustering length and typical dark matter halo masses for Lyman break galaxies, Lyman alpha emitting galaxies, Distant red galaxies, Sub-millimeter galaxies, and Damped Lyman $\\alpha$ absorption systems. I also discuss five key unsolved problems in galaxy formation and prognosticate advances that the near future will bring. ",
        "introduction": "\\label{sec:boundary}  \\subsection{Initial Conditions:  \\boldmath $\\Lambda$CDM Cosmology} \\label{sec:initial}  The initial conditions for the formation of galaxies are provided by the now-standard $\\Lambda$CDM cosmological model. The combined results of the WMAP satellite study of Cosmic Microwave Background anisotropies, large-scale structure, and Type Ia supernovae observations yield best-fit values for the cosmological parameters of roughly $\\Omega_\\Lambda = 0.7$, $\\Omega_m=0.3$, $\\Omega_b = 0.04$, and $H_0 = 70h_{70}$km s$^{-1}$ Mpc$^{-1}$ \\citep{bennettetal03}.\\footnote{We include $h_{70}$, analogous to the traditional parameter $h\\equiv h_{100}$, even though its value appears quite close to 1.} The original model of galaxy formation was Monolithic Collapse \\citep*{eggenls62}, where gravitational collapse of a cloud of primordial gas very early in the lifetime of the Universe formed all parts of each galaxy at the same time.  Modern evidence rules out this model on two fronts; the widely varying ages of different components of the Galaxy provide a counter-example, and the $\\Lambda$CDM cosmology predicts ``bottom-up'' i.e. hierarchical rather than ``top-down'' structure formation.   Hierarchical structure formation is a generic feature of Cold Dark Matter (CDM) models.  Small overdensities are able to overcome the cosmological expansion and collapse first, and the resulting dark matter ``halos'' merge together to form larger halos which serve as sites of galaxy formation. This process continues until the present day, making galaxy formation an ongoing process. The nearly-scale-invariant primordial power spectrum inferred from combining WMAP with large-scale structure observations provides power on all scales in the distribution of CDM. The baryons fall into the CDM potential wells after decoupling, leaving only trace evidence of their previous acoustic oscillations as a series of low-amplitude peaks in the matter power spectrum. The non-linear collapse of dark matter overdensities occurs on larger and larger scales, so the typical collapsed halo mass grows with time, but no preferred scale is introduced. $\\Lambda$CDM therefore provides a distribution of halos where galaxies can form, with the details of the process up to baryonic physics.  Despite the lack of preferred galaxy scales in the distribution of dark matter halos, baryonic physics causes galaxies to have minimum and maximum masses. The maximum mass is that of CD galaxies in cluster centers with baryonic masses $\\sim 10^{12}\\mathrm{M}_\\odot$ and virial masses  $\\sim 10^{13}\\mathrm{M}_\\odot$; there are $\\sim10^{14}$M$_\\odot$ of baryons available in a rich cluster but virialization of galaxies and heating of gas to the high virial temperature prevent most of this mass from finding its way to the central galaxy.  The minimum mass observed today is that of dwarf galaxies, $\\sim10^8\\mathrm{M}_\\odot$, but galaxies may initially have formed as small as  $10^6\\mathrm{M}_\\odot$ (the baryonic Jean's mass after recombination i.e. the minimum mass for which gravity overwhelmed pressure support).  Explaining the lack of observed galaxies with circular velocities below 30 km/s is a major goal; it is suspected that feedback from supernovae explosions may have quenched star formation in such low-mass objects immediately after a single burst of star formation \\citep{dekels86}.  The growth of cosmological structure and collapse of dark matter halos is a feature of the matter-dominated epoch.  During radiation-domination, perturbations on scales smaller than the sound horizon were unable to grow due to acoustic oscillations in the photon-baryon fluid that gave rise to the famous peaks in the CMB angular power spectrum and the lower-amplitude peaks in the matter power spectrum.  Now that we have entered a phase of dark energy domination, structure growth is slowing and will cease entirely as the universe enters a new phase of inflation.  This cosmological ``freeze-out'' in structure formation is recent, since equality between the dark energy and matter densities occurred at $z_{eq}=0.4$. The slowing of structure formation occurs gradually, so the growth of cosmological structure continued nearly unabated until $z_{eq}$, even though we see strong observational evidence for ``downsizing'' at $z<1$ where high-mass galaxies grow far more slowly than lower-mass galaxies \\citep[e.g.][]{treuetal05,smith05}. Another term being used by some is ``anti-hierarchical'', which is basically a synonym for ``downsizing'' but seems to imply inconsistency with hierarchical cosmology.  However, the observed freeze-out in galaxy (and possibly supermassive black hole) formation in massive galaxies is not inconsistent with CDM models; rather, it appears to be caused by baryonic feedback which is not well understood at present (see \\S \\ref{sec:problems}). The slowing of cosmological structure growth since $z\\simeq0.4$ may, however, play a role in the recent decline of the cosmic star formation rate density discussed by \\citet{belletal05}.  \\subsection{Final Conditions:  Low-redshift Galaxies} \\label{sec:lowz} The study of galaxy formation is made easier by having full boundary conditions.  The final conditions are the Hubble sequence of mature galaxies we see in the nearby universe at redshift zero.  Indeed, much has been learned about galaxy formation from ``archaeological'' evidence in the ages and chemical abundances of various Galactic stellar populations, and expanding these studies to the rest of the Local Group and beyond is quite useful.  Nonetheless, there are great advantages to observing galaxies in the act of formation, which motivates the study of high-redshift galaxies. At $z>2$, galaxy-mass halos are rare so the majority of galaxies we observe reside in dark matter halos that have only recently collapsed i.e. at high-redshift most galaxies are young. In this sense, $z>2$ can be considered the epoch of galaxy formation. ",
        "conclusions": "\\label{sec:attractions}  The speakers have been asked to discuss major advances expected in the coming decade.  For galaxy formation, I will go on record with three promising predictions and one slightly fascetious warning.   {\\it The coming years will see the unification of galaxy formation and evolution}. Until very recently, galaxy formation was studied at $z>2.5$ and galaxy evolution was studied at $z<1$ and the period $1<z<2.5$ was referred to as the ``redshift desert''.  But technological advances in NIR imaging and spectroscopy have made the rest-frame Balmer/4000\\AA\\ break and nearby emission lines available for study in distant galaxies. Development of these ``needle-in-a-haystack'' techniques now allows us to successfully find evolved galaxies at $z>2$ even though these objects may be rare at those epochs.  Hence we are beginning to study objects at $z\\sim3$ that formed at $z>6$ which may turn out to be much easier than observing $z>6$ galaxies directly. Imaging with the Spitzer satellite is enabling the first studies of rest-frame near-infrared emission from $z>2$ galaxies, breaking degeneracies between age and dust. Deep imaging and slitless spectroscopy with the GALEX satellite are revealing the analogs of Lyman break galaxies at low redshift \\citep{burgarellaetal05}. These combined studies may make it possible to piece together a rough evolutionary sequence, e.g. DLA$\\rightarrow$LAE$\\rightarrow$LBG$\\rightarrow$SMG$\\rightarrow$DRG, that would form part of a {\\it grand unified} model of high-redshift galaxies and AGN.  {\\it We will be able to study the interstellar medium in emission at high-redshift}. ALMA will enable studies of molecular gas in young galaxies through high-order CO lines. The [CII] 158 micron line, which dominates the cooling of the Cold Neutral Medium phase at both low and high redshift, should be detectable for galaxies with large gas mass or rapid cooling equilibrating the heating due to starbursts. The current set of Early Universe Molecular Emission Line Galaxies consist mostly of quasars and are reviewed (and assigned the questionable TLA ``EMG'') by \\citet{solomonv05}. Both CO and [CII] have now been detected in $z>6$ SDSS quasars, where they provide the best direct probes of the quasar host galaxies \\citep{bertoldietal03,walteretal04,maiolinoetal05}. Detecting these lines and the sub-millimeter dust continuum from protogalaxies with ALMA will allow us to probe a multivariate mass function of gas mass, molecular mass, dust mass, and stellar mass.  Even ALMA sensitivity may only allow detections of the tip of the gas-mass function, but this will provide a complementary set of objects to the tip of the rest-frame-UV and rest-frame-optical luminosity functions currently studied at high redshift, and much can be learned from the intersection and union of these samples.  {\\it High-redshift galaxies will be used to constrain dark energy properties}.   It has recently been shown \\citep{seoe03,linder03,blakeg03} that the scale of baryon acoustic oscillations provides a ``standard rod'' that can be measured in the clustering of high-redshift galaxies. The measurement will constrain the dark energy equation-of-state as a function of redshift, $w(z)$, via its influence on the expansion history of the universe.  The measurement can be performed at any redshift where the line-of-sight starting at $z=0$ is sufficiently influenced by the dark energy, making $z=1$ and $z=3$ equally acceptable. Of order a million redshifts are needed, and the most likely surveys to accomplish this ambitious goal appear to be HETDEX using the VIRUS instrument under construction for HET and the wide-field multi-fiber spectrograph KAOS proposed for Gemini.  {\\it The rapidly increasing sophistication of studies of the high-redshift universe will generate even more jargon}.  We are already debating proper nomenclature for special categories of DLAs at lower column density (sub-DLAs) and those found in gamma-ray burst afterglows (burst-DLAs or bDLAs).  Four-letter object acronyms (FLOAs?) are going to be part of the future."
    },
    "0512/astro-ph0512498_arXiv.txt": {
        "abstract": "The two-point autocorrelation function  of ultra-high energy cosmic ray (UHECR) arrival directions has a broad maximum around 25 degrees, combining the data with energies above $4\\times 10^{19}$~eV (in the HiRes energy scale) of the  HiRes stereo, AGASA, Yakutsk and SUGAR experiments. This signal is not or only marginally present analyzing events of a single experiment, but becomes  significant when data from several experiments are added. Both the energy dependence of the signal and its angular scale might be interpreted as first signatures of the large-scale structure of UHECR sources and of intervening magnetic fields.   \\begin{small} PACS: 98.70.Sa % \\end{small} ",
        "introduction": "The sources of ultra-high energy cosmic rays (UHECR) are despite of more than 40~years of research still unknown. Main obstacle for doing charged particle astronomy are deflections of the primaries in the Galactic and extragalactic magnetic fields. While the magnitude and the structure of extragalactic magnetic fields are to a large extent unknown, already deflections in the Galactic magnetic field alone are large enough to prevent UHECR astronomy if the primaries are heavy nuclei~\\cite{nuc,KST05}. Assuming optimistically that the primaries are protons, typical deflections in the Galactic magnetic field are around five degrees in most part of the sky at $E=4\\times 10^{19}$~eV~\\cite{KST05}. Therefore, it might be possible to perform charged particle astronomy, if moreover deflections in extragalactic magnetic fields are sufficiently small.  This scenario can be divided in two quite different sub-cases: In the first one, a small number of point sources results in small-scale clusters of arrival directions around or near the true source positions. Accumulating enough events, the identification of sources will become possible using e.g.\\ correlation studies. Various studies have been pursued in this direction~\\cite{ssc,corr}. In the second sub-case, the number of point sources is too large to receive two or more UHECRs from the same source with the present statistics. However, the sky density of UHECRs reflects the large-scale structure of the sources and, possibly, of the intervening magnetic fields~\\cite{lss,napoli}. Therefore, the measured UHECR distribution is anisotropic and over-/underdense regions exist that reflect the angular size of up-to 15--20 degrees of  typical structures in the galaxy distribution. Obviously, Nature might have chosen a mixture of these two extreme possibilities: The vast majority of UHECR sources might produce only singlet events, while a subclass of sources with extreme luminosity might be detectable as point sources via small-scale clustering studies. Furthermore, point sources might be easier to identify at the highest energies, if the number density of sources decreases with  the maximal energy $E_{\\rm max}$ to which they can accelerate as argued in Ref.~\\cite{spec}.   In this work, we study the arrival direction distributions of the UHECRs, putting emphasis in contrast to most earlier studies on intermediate angular scales. Since these two-dimensional  distributions average three-dimensional structures (with typical scale $L$) over the mean free path $l$ of UHECRs, no anisotropies reflecting the large-scale structure of sources are expected for $l\\gg L$. To obtain an optimal compromise between the number of events used, the mean free path $l$ of UHECRs and deflections in magnetic fields, it is important to use a consistent energy scale when combining different experiments. Therefore, we discuss first in Sec.~\\ref{data} the publicly available UHECR data and how we rescale the data of different experiments to a common energy scale. In Sec.~\\ref{analysis}, we analyze then the arrival direction distributions of the used UHECRs and discuss in Sec.~4 our results, before we summarize in Sec.~\\ref{sum}. ",
        "conclusions": "Our results, if confirmed by future independent data sets, have several important consequences.  Firstly, anisotropies on intermediate angular scales constrain the chemical composition of UHECRs. Iron nuclei propagate in the Galactic magnetic field in a quasi-diffusive regime at $E=4\\times 10^{19}$~eV and all correlations would be smeared out on scales as small as observed by us. Therefore, models with a  dominating extragalactic iron component at the highest energies are disfavored by anisotropies on intermediate angular scales.  Secondly, the probability that small-scale clusters are indeed from point sources will be reduced if the clusters are in regions with an higher UHECR flux. For example, the AGASA triplet is located in an over-dense spot (cf. map in Fig.~\\ref{skymap1}) and the probability to see a cluster in this region by chance is increased. In contrast, the observation of clusters in the \"voids\" of Fig.~\\ref{skymap1} would be less likely by chance than in the case of an UHECR flux without medium scale anisotropies.  However, the most important consequence of our findings is the prediction that astronomy with UHECRs is possible at the highest energies. The minimal energy required seems to be around $E'=4\\times 10^{19}$~eV, because, as one can see from Figs.~\\ref{skymap2} and \\ref{p_ch_E}, at lower energies UHECR arrive more and more isotropically. This trend is expected, because at lower energies both deflections in magnetic fields and the average distance $l$ from which UHECRs can arrive increase. Since the two-dimensional skymap corresponds to averaging all three-dimensional structures (with typical scale $L$) over the distance $l$, no anisotropies are expected for $l\\gg L$. Thus, if the signal found in this analysis will be confirmed it has to be related to the local large scale structure.  Finally, we note that to check this signal an independent data set of order $O(100)$ events with $E'>4\\times 10^{19}$~eV is required. This agrees roughly with the finding of Ref.~\\cite{napoli} that around 100 events are needed to reject the hypothesis that the UHECR sources trace the galaxy distribution."
    },
    "0512/astro-ph0512451_arXiv.txt": {
        "abstract": "\\emph{INTEGRAL} and \\emph{RXTE} performed three simultaneous observations of the nearby radio galaxy Centaurus A in 2003 March, 2004 January, and 2004 February with the goals of investigating the geometry and emission processes via the spectral/temporal variability of the X-ray/low energy gamma ray flux, and intercalibration of the \\emph{INTEGRAL} instruments with respect to those on \\emph{RXTE}. Cen~A was detected by both sets of instruments from 3--240 keV. When combined with earlier archival \\emph{RXTE} results, we find the power law continuum flux and the line-of-sight column depth varied independently by 60\\% between 2000 January and 2003 March. Including the three archival \\emph{RXTE} observations, the iron line flux was essentially unchanging, and from this we conclude that the iron line emitting material is distant from the site of the continuum emission, and that the origin of the iron line flux is still an open question. Taking X-ray spectral measurements from satellite missions since 1970 into account, we discover a variability in the column depth between 1.0$\\times10^{23} \\rm cm^{-2}$ and 1.5$\\times10^{23} \\rm cm^{-2}$ separated by approximately 20 years, and suggest that variations in the edge of a warped accretion disk viewed nearly edge-on might be the cause. The \\emph{INTEGRAL} OSA~4.2 calibration of JEM-X, ISGRI, and SPI yields power law indices consistent with the \\emph{RXTE} PCA and HEXTE values, but the indices derived from ISGRI alone are about 0.2 greater. Significant systematics are the limiting factor for \\emph{INTEGRAL} spectral parameter determination. ",
        "introduction": "At a distance of $\\sim$3.5 Mpc \\citep{Hui93}, the radio galaxy Centaurus~A is one of the nearest and brightest active galactic nuclei (AGN).  Since its discovery over three decades ago \\citep{Bowyer70} (Note, however, \\citealt{Byrum70}), many X-ray to gamma-ray instruments have shown it to have non-thermal, power law emission extending to the MeV range \\citep{Baity81,Steinle98}.  \\emph{Hubble Space Telescope} (HST) observations have revealed evidence for a small ($\\sim$40 pc), inclined disk of ionized gas (Schreier et al. 1998) around a $\\sim$$10^9$ M$_\\odot$ black hole, and perhaps a region near the black hole evacuated by the jet \\citep{Marconi00}. \\citet{Karovska03} suggested that mid-IR observations resolved the nuclear region of size $\\sim$3 pc, and together this might explain the lack of a hidden broad line region in Cen~A \\citep{Alexander99}.  The power law index of Cen~A below 100 keV has remained at $\\sim1.8$ for the last 40 years with the exception of 1972--1973 when \\emph{OSO-7} found the index to be $\\sim1.2-1.4$ \\citep{Winkler75,Mushotzky76}. The X-ray spectrum does not show any reflection component or a significant broad iron K$\\alpha$-line \\citep{Wozniak98,Rothschild99,Benlloch01}, and yet it is rapidly variable \\citep{Morini89}. This would imply, in contrast to radio-quiet Seyfert galaxies, that a cold accretion disk extending down close to the black hole may not be the source of the high energy radiation and the reprocessor is relatively far away. Furthermore, the hard X-ray spectra of individual Seyfert~1s show an underlying continuum which is a power law, with a nearly exponential rolloff with a folding energy of $\\sim$100--300 keV \\citep{Johnson97}. Cen A, on the other hand, does show $\\sim$MeV range emission, thereby indicating it is quite different from the radio-quiet Seyferts.  Current observations of Cen~A are presented in \\S~2, and the methods of analyzing both the \\emph{INTEGRAL} and \\emph{RXTE} data will be found in \\S~3. In \\S~4 we present the results from fitting the combined instrument data from each mission from the 3 simultaneous observations and from reanalyzing the 3 previous \\emph{RXTE} observations with HEASOFT release (5.3.1). In \\S~5 we use these results to discuss the emission processes in the nuclear region of Cen~A, and present our conclusions in \\S~6. We present an extensive analysis of each instrument's response to Cen~A as well as an inter-calibration of the 2 missions in the Appendix. Included in the Appendix is also a study of the stability of the PCU2 calibration over the \\emph{RXTE} mission as determined from observations of Cas~A. ",
        "conclusions": "The six \\emph{RXTE} observations over the last nine years have shown that the 3--240 keV Cen~A spectrum can be described by a single absorbed power law plus an iron emission line. We have measured the column density to Cen~A to about 1\\%, the power law index and iron line centroid energy to better than 1\\%, and the iron line flux to approximately 10\\%. While still systematics dominated, \\emph{INTEGRAL} determines the power law index to a few percent and the column depth to 10--20\\% from simultaneous observations with the last three \\emph{RXTE} observations in 2003 and 2004. We have provided an in depth comparison of \\emph{RXTE} and \\emph{INTEGRAL} instruments using the latest knowledge of the instrument responses and techniques for addressing systematic errors. Appendix~\\ref{sec:50} gives the improved \\emph{INTEGRAL} spectral results using OSA~5.0 analyzed after submission of this paper.  The mean values of $N_{\\rm H}$ and $\\Gamma$ resulting from the spectral analyses of individual instruments on \\emph{INTEGRAL} and \\emph{RXTE} and of simultaneous fitting of all insturments on a given satellite are given in Table~\\ref{tab:mean_values}. All five instruments' spectral parameters are in agreement at the 90\\% confidence level, except for the ISGRI determination of the power law index. The ISGRI value of $\\Gamma$ is significantly larger ($\\Delta \\Gamma \\sim$ 0.2$\\pm$0.1). This discrepancy is $\\sim0.1$ using OSA~5.0.  From the \\emph{RXTE} observations, we have detected a 60\\% increase in the mean column depth to Cen~A between 2000 and 2003, and this increase was not correlated to either the spectral index or the 2--10 keV flux. The increase in column depth was accompanied with a small drop in iron line flux that is not significant at the 99\\% confidence level. By considering past satellite measurements of the absorbing column, we note two episodes of $N_{\\rm H}=1.0\\times10^{23} \\rm cm^{-2}$ separated by $\\sim$20 years, and speculate that variability in the structure of the outer edge of the warped accretion disk could explain the observed variability. Since the continuum shape and iron line flux did not vary significantly, we suggest that they are separate from each other and the intervening material.  Where, then, does the Fe K$\\alpha$ line originate? Since the line strength does not correlate with $N_{\\rm H}$, we do consider it unlikely that the Fe line is produced in the absorbing material.  The similarity of the Cen A line parameters with those now seen with \\emph{Chandra} or \\emph{XMM-Newton} in many AGN, low-luminosity Seyferts, and high-luminosity QSOs, which all have narrow lines (to the resolution of the observation) at an energy consistent with emission by neutral Fe and equivalent widths of less than about 150\\,eV \\cite[see, e.g.,][]{Pounds02,Yaqoob04,Reynolds04,Jimenez05} points at a similar origin for these features. The line parameters are characteristic for a line origin as a fluorescence line in material that is irradiated by X-rays. The possible location of the line emitting region is either the outer regions of the accretion disk or a medium separate from the disk, such as the torus posited in unifying models for AGN. As shown by \\citet{Ghisellini94} and \\citet{Leahy93}, for parameters typically assumed for the torus and the central source, equivalent widths of 100\\,eV are expected. Alternatively, as discussed by \\citet{Yaqoob01} and \\citet{Jimenez05}, such lines could also originate in the broad line region, with the size of the region again being sufficiently large that one would not expect a correlation between the flux from the central source and the line strength.  We note, however, that the presence of a jet is one crucial difference between radio quiet Seyfert galaxies and objects such as Cen A or radio-loud QSOs. As has been recently shown for both Galactic black hole candidates and for low luminosity AGN, a significant fraction of the X-ray emission from these systems could also be explained by synchrotron-self Comptonization radiation from the base of the radio jet and thus not be due to thermal Comptonization \\cite[][and therein]{Markoff05,Falcke04}. The lack of a roll-over and reflection component in Cen A, as opposed to Seyfert galaxies, could therefore also be due to jet emission, with beaming or the small footprint of the jet reducing the amount of reflection expected. In these models, the advected flow at the base of the jet produces the hot electrons and SSC radiation could extend to high energies. In Seyfert galaxies, the jet could be less beamed/thermal, or thermal Comptonization could be more dominant, resulting in the observed roll-over. If this is the case, then the observed line emission could also originate from a region that, due to beaming, is significantly more illuminated by the jet than the disk. As we would not see the illuminating radiation but only the Fe K$\\alpha$ emission, which is emitted isotropically, no correlation between the continuum and the Fe line flux would be expected. In addition, such a model could also explain the absence of hidden broad lines from radio galaxies like Cen A, as the jet may sweep out the material that otherwise would form the broad line region. It is beyond the scope of this paper to quantify these effects, as they all strongly depend on the unknown jet-kinematics at the base of the jet.  Naturally, the iron line emission could be a combination of these effects, where the fraction transmitted through a torus (where the line flux would be proportional to the primary flux) is small compared to that reradiated at a distance. The lack of an iron edge at 7.1 keV in the \\textsl{XMM-Newton} and \\textsl{Chandra} spectra \\citep{Evans04} as well as the \\textsl{RXTE} data, further complicates the question of the origin of the iron line in Cen~A."
    },
    "0512/astro-ph0512337_arXiv.txt": {
        "abstract": "We study the frequency and angular dependences of Cherenkov radio pulses originated by the excess of electrons in electromagnetic showers in different dense media. We develop a simple model to relate the main characteristics of the electric field spectrum to the properties of the shower such as longitudinal and lateral development. This model allows us to establish the scaling of the electric field spectrum with the properties of the medium such as density, radiation length, Moli\\`ere radius, critical energy and refraction index. We normalize the predictions of the scaling relations to the numerical results obtained in our own developed GEANT4-based Monte Carlo simulation, and we give a unified parameterization of the frequency spectrum and angular distribution of the electric field in ice, salt, and the lunar regolith, in terms of the relevant properties of the media. Our parameterizations are valid for electromagnetic showers below the energy at which the Landau-Pomeranchuk-Migdal effect starts to be relevant in these media. They also provide an approximate estimate of radio emission in hadronic showers induced by high energy cosmic rays or neutrinos. ",
        "introduction": "Neutrinos are the most suitable candidates to extend astronomical observations to the ultra high energy (UHE) regime, above $\\sim 1$ TeV. Unlike charged particles, they point directly to the source where they were produced. Unlike neutrons, they are stable. Unlike photons, they can penetrate large amounts of matter, and they are not attenuated by cosmic backgrounds. They are not optimal candidates though, due to their small interaction cross section, and to their small expected fluxes at ultra high energy \\cite{gaisser95,learned00,halzen02}. However these two latter difficulties may be overcome with a large enough detector volume of at least $1~{\\rm km^3}$ \\cite{halzen94}.  The ``conventional\" technique to achieve a large detector volume exploits the observation of Cherenkov light emitted by high energy neutrino-induced muons and showers in water and ice \\cite{andres01,AMANDA,Antares,Baikal,spiering03}. An alternative to it is the search for coherent Cherenkov radio pulses in the MHz-GHz radio frequency range, produced by neutrino induced showers in dielectric, transparent, dense media. When the wavelength of the emitted radiation is larger than the typical dimensions of the shower the emission is coherent. The contribution to the electric field from positive and negative particles would approximately cancel out were it not for the existence of an excess of electrons over positrons in the shower \\cite{askaryan62}. The charge asymmetry arises from photon and electron scattering on the surrounding medium that sweeps electrons into the shower, and from positron annihilation in flight that contributes to the excess charge by terminating positron trajectories. Due to the coherent nature of the radio emission the power in radio waves scales as the square of shower energy.  Dense media have the advantage that shower dimensions are of the order of meters, and coherence extends to high frequencies, typically $\\sim$ GHz, where more power is available because the coherently radiated electric field is proportional to frequency. Besides, large formations of dense, transparent media such as ice and salt exist in nature having attenuation lengths in radio frequencies of a few hundred meters \\cite{Icebook,SalSA02,vandenBerg_ARENA}. They have the advantage of being located in low noise environments such as the South Pole or deep salt mines. A single antenna in such a medium makes a detector of effective volume of $\\sim 0.1~{\\rm km^3}$ water equivalent \\cite{price96,Frichter95}. The cost-effectiveness of antennas also adds to the attractiveness of the radio technique.  Several experiments are searching for neutrinos exploiting the radio technique in dense media. The RICE experiment \\cite{RICEdetector} is an array of antennas buried in the Antarctic ice cap that has imposed stringent limits on astrophysical neutrino fluxes in the energy range above $\\sim 10^{17}$ eV \\cite{RICElimits}. The Goldstone Lunar Experiment (GLUE) has used two large radiotelescopes to search for pulses produced by neutrino-induced showers on the Moon's regolith \\cite{GLUE00}. With no neutrino candidates in 120 hours of observation, GLUE has established competitive upper bounds on astrophysical neutrinos of energy above $\\sim 10^{19}$ eV \\cite{GLUElimits}. The technique could be of interest to search for ultrahigh energy cosmic rays \\cite{markov86,alvarez_RADHEP} and is being explored using the Westerbork array of radiotelescopes in the Netherlands \\cite{Bacelar_ARENA}. It will also be exploited by the proposed LUNASKA experiment in Australia \\cite{LUNASKA}. A recently approved experiment is ANITA \\cite{ANITA_ARENA}, a balloon-borne antenna that circles the Antarctic continent scanning for large radio pulses. The idea of using large salt domes as targets for neutrino interactions has also been explored  and there is a proposed experiment named SalSA \\cite{gorham-aspen}, as well as other initiatives such as the ZESANA proposal in the Netherlands \\cite{vandenBerg_ARENA}.  It is also very important to remark that the dominant emission mechanism in dense media has been experimentally confirmed in accelerator measurements at SLAC, using bunches of bremsstrahlung photons as projectiles, and sand and salt as target medium \\cite{saltzberg_SLAC,gorham_salt}. This experiment has tested the theoretical predictions originally made by Askar'yan in the 1960's \\cite{askaryan62}.  Given the current and future experimental efforts, the motivation and aim of our work are clear: A good understanding and comprehensive characterization of the dependence of the radio signal on frequency and observation angle in different dense media are needed to interpret the data that is being collected \\cite{RICElimits,GLUElimits}, as well as to evaluate the capabilities and potential of the planned experiments looking for neutrinos and cosmic rays.  In this paper we show that many of the observables of the electric field spectrum emitted by an electromagnetic shower developing in dense media, scale to a few percent level with several properties of the media, such as density ($\\rho$), radiation length ($X_0$), Moli\\`ere radius ($R_M$), critical energy ($E_c$), and index of refraction ($n$). This scaling allows us to predict the properties of radio pulses in different media without performing time-consuming Monte Carlo simulations. Our work provides further insight into the close relation between shower development and radio emission, which may help designing future experiments exploiting the radio technique, as well as assessing the potential of those currently under construction.  This paper is structured as follows. In Section \\ref{pulse} we briefly review radio emission from electromagnetic showers. In Section \\ref{toy} by means of a simple toy model we relate the main features of the radiopulse to general characteristics of shower development. This allows us to obtain the scaling of radiopulse with the relevant parameters of the medium. In Section \\ref{media} we normalize these scaling relations using our own developed GEANT4-based Monte Carlo simulation \\cite{GEANT4,alvarez03} to numerically calculate the electric field spectrum in ice, salt and the lunar regolith. We also give in this section unified parameterizations of our numerical results to be used in practical applications and establish their range of applicability. ",
        "conclusions": ""
    },
    "0512/astro-ph0512101_arXiv.txt": {
        "abstract": "The Advanced Camera for Surveys is equipped with three prisms in the Solar Blind (SBC) and High Resolution (HRC) Channels, which together cover the 1150 -- 3500 \\AA\\ range, albeit at highly non-uniform spectral resolution. We present new wavelength- and flux calibrations of the SBC (PR110L and PR130L) and HRC (PR200L) prisms, based on calibration observations obtained in Cycle 13.  The calibration products are available to users via the ST-ECF/aXe web pages, and can be used directly with the aXe package.  We discuss our calibration strategy and some caveats specific to slitless prism spectroscopy. ",
        "introduction": "The Advanced Camera for Surveys is currently the only instrument on HST which provides spectroscopy in the UV-optical range.  With STIS being unavailable (at least for the moment), the interest in the ACS spectroscopic modes has increased substantially. An extensive calibration effort was undertaken in Cycle 13 in order to provide much improved wavelength- and flux calibrations for the prism modes which had seen little use in earlier cycles. The ACS has two prisms (PR110L and PR130L) installed in the Solar Blind Channel (SBC), both covering the wavelength range from roughly 1200\\AA\\ -- 2000\\AA . The main difference between the two SBC prisms is that the sensitivity of PR110L extends below the geocoronal Ly $\\alpha$, while that of the PR130L does not. The sensitivity below 1216 \\AA\\ results in a significantly higher sky background for the PR110L.  The one prism (PR200L) in the High Resolution Channel covers the range $\\sim1800$\\AA -- 3500\\AA . In this paper we present the calibration observations and discuss the derived trace-, wavelength and flux calibrations. The calibrations presented here are implemented in configuration files for the aXe package (K{\\\"u}mmel et al., these proceedings) and are made available to users via the aXe web pages ({\\tt http://www.stecf.org/software/axe}).   Spectroscopic observations with the ACS prisms share many similarities with the G800L grism in the WFC and HRC channels, but there are also some differences. In both cases the object spectrum (or spectra) may fall anywhere on the detector, and the reference point for the wavelength scale is typically established using a direct image taken immediately before (or after) the prism exposure. The direct image is even more crucial for the prism modes, which do not show a zeroth order spectrum (nor any higher- or negative orders).  Contrary to the case of grism spectroscopy, the wavelength scale of the prism spectra is highly non-linear, with the spectral resolution decreasing towards longer wavelengths (see below). For the PR200L, this causes a ``red pile-up'' with wavelengths between 4000 \\AA\\ and 10000 \\AA\\ being compressed into only 7 pixels.  This pile-up can be particularly troublesome for red objects, where light from the diffraction spikes and outer halo of the PSF may contaminate the bluer parts of the spectrum. These effects are still poorly quantified.  No such pile-up is seen for the SBC prisms, due to the lack of sensitivity at wavelengths $>2000$\\AA\\ for the MAMA detector. ",
        "conclusions": "The SBC and HRC prisms are now fully calibrated, both in wavelength and flux. The calibration products are available for use with the aXe package, and can be downloaded from the aXe web pages.  Future calibration observations will aim to monitor the stability of the prism modes and possibly provide various improvements, such as quantifying the effect of scattered light from the red pile-up in the PR200L. A better understanding of the general properties of the SBC and HRC detectors, such as CTE effects, may also lead to improvements in prism spectroscopy."
    },
    "0512/astro-ph0512271_arXiv.txt": {
        "abstract": "We explore the consequences of new observational and theoretical evidence that long gamma-ray bursts prefer low metallicity environments. Using recently derived mass-metallicity correlations and the mass function from SDSS studies, and adopting an average cosmic metallicity evolution from \\citet{kewley2005} and \\citet{savaglio2005} we derive expressions for the the relative number of massive stars formed below a given fraction of solar metallicity, $\\varepsilon$, as function of redshift. We demonstrate that about 1/10th of all stars form with $\\varepsilon < 0.1$. Therefore, a picture where the majority of GRBs form with $\\varepsilon < 0.1$ is not inconsistent with an empirical global SN/GRB ratio of 1/1000. It implies that (1) GRB's peak at a significantly higher redshift than supernovae; (2) massive star evolution at low metallicity may be qualitatively different and; (3) the larger the low-metallicity bias of GRBs the less likely binary evolution channels can be significant GRB producers. ",
        "introduction": "It become clear in the last few years that long gamma-ray bursts are associated with the endpoints of massive star evolution. They occur in star forming regions at cosmological distances \\citep{jakobsson2005}, and are assiciated with supernova-type energies. The collapsar model explains gamma-ray burst formation via the collapse of a rapidly rotation massive iron core into a black hole \\citep{woosley1993}. The short time scale of gamma-ray emission requires a compact stellar size, of the order of lightseconds. This constraint leaves only massive Wolf-Rayet stars as possible progenitors.  However, this poses a difficulty: Wolf-Rayet stars in the local universe are known to have strong stellar winds \\citep{nugis1998}, which lead to a rapid spin-down \\citep{langer1998} --- in agreement with the absence of signatures of rapid rotation in the Galactic Wolf-Rayet sample \\citep{eenens2004}.  It is the current understanding that the ratio of gamma-ray bursts to supernovae is about 1/1000, based on about 1 - 2\\, 10$^{-6}$ observed bursts per supernova in the BATSE sample \\citep{porciani2001}, and a beaming factor of $\\sim 500$ \\citep{frail2001,yonetoku2005}. This implies that about 1 out of 100 Wolf-Rayet stars produces a gamma-ray burst \\citep{putten2004}. These low values are thought to support the idea that rather exotic binary evolution channels might constitute the main evolutionary paths towards gamma-ray bursts \\citep{podsiadlowski2004,fryer2005}, corroborated by the two basic problems of single star models to produce collapsars \\citep{petrovic2005}: (1) the spin-down of the stellar core due to magnetic core-envelope coupling \\citep{spruit2005,petrovic2005} --- which is required to understand the slow rotation of young pulsars \\citep{ott2005} and white dwarfs \\citep{berger2005}; and (2) the spin-down due to Wolf-Rayet winds mentioned above.  However, recent single star models have overcome both problems for suitable initial conditions. The models by \\citet{woosley2005} and \\citet{yoon2005} avoid problem (1) by rapid rotation --- which keeps the stars nearly chemically homogeneous and thus avoids the formation of a massive envelope --- and problem (2) by choosing a low enough metallicity --- which, according to recent evidence \\citep{crowther2002, vink2005}, reduces the Wolf-Rayet mass loss rates. These are the first single star evolution models fulfilling the requirements of the collapsar model which are at the same time fully consistent with the slowly rotating stellar remnants in our Galaxy. However, they predict long GRBs only for  metallicities of about $Z/Z_{\\odot} \\lesssim 0.1$.  In this context, the growing empirical evidence that the long bursts indeed prefer a low-$Z$ environment is remarkable. While indirect evidence from gamma-ray burst host galaxies is pointing towards low metallicity \\citep{fynbo2003, conselice2005, fruchter2005}, direct metallicity determinations yield sub-solar values down to 1/100th of the solar metallicity \\citep{gorosabel2005, chen2005, starling2005}. While the observational and the theoretical evidence for long gamma-ray bursts occurring at low metallicity needs further confirmation, we are motivated by the findings reported above to explore the consequences of such a possibility. ",
        "conclusions": "To produce a GRB, Wolf-Rayet stars at core collapse are required to have sufficient angular momentum. Stellar evolution models which include magnetic fields predict too slowly-rotating cores for models which develop an extended, massive envelope after the main sequence. Current evolutionary models that include rotation predict extended envelopes for the vast majority of massive stars in the Galaxy or Magellanic Clouds --- in agreement with the number of blue and red supergiants \\citep{maeder2001, maeder2005}. Only the fastest rotators are thought to be able to avoid extended envelopes, for metallicities below about $Z_{\\odot}/10$, or $\\varepsilon = 0.1$ \\citep{yoon2005, woosley2005}.  Most importantly, the numbers derived for $\\varepsilon = 0.1$ in Table~1 appear not to be in conflict with observations. Per 1000 supernovae in the universe, 160 are predicted to occur from stars with $Z > Z_{\\odot}/10$, out of which 22 would produce a black hole.  Thus, producing one GRB per 1000 supernovae globally in the universe (cf. Sect.~1) seems possible.  On the other hand, a large fraction of the 22 black holes may be born without producing a GRB: not all of them may occur in WR stars but rather in more extended stars \\citep{maeder2005}, and the most massive ones would lose too much angular momentum in a wind, even for metallicities as low as $Z > Z_{\\odot}/10$ \\citep{yoon2005}. While within the chemically-homogeneous-evolution scenario for GRB formation \\citep{yoon2005, woosley2005} a GRB/BH fraction of 1/20 can certainly be obtained \\citep{yoon2006}, this is therefore unlikely for exotic binary channels for GRB production --- i.e. for channels through which only a small fraction of stars of any initial mass evolves. Clearly, the more the long GRBs are confined to low metallicities, the more unlikely it is that binary evolution is needed to explain the majority of events.   The empiric cosmic GRB to SN ratio of about 1/1000 (cf. Sect.~1) can not directly rule out more extreme values of $\\varepsilon$, i.e. $\\varepsilon= 0.1 - 0.01$ (cf. Table~1); in fact, for $\\varepsilon= 0.01$, about 10\\% of all massive stars with $Z < Z_{\\odot}/100$ would need to produce a GRB. However, it would imply that the formation of {\\em every} black hole would be accompanied by a GRB. Furthermore, the GRB rate would peak only at a redshift of about $z=10$. A value of $\\varepsilon= 0.01$ appears thus unlikely. Furthermore, our models with $\\beta=5$ produce such a small local GRB/SN ratio that they seem to be ruled out.  We find that a restriction of GRBs to low metallicities ($Z < Z_{\\odot}/10$, i.e. $\\varepsilon= 0.1$) has the following consequences:  \\begin{itemize} \\item GRBs do not follow star formation in an unbiased manner. For example, for an overall star formation rate which predicts a preceived SN peak at a redshift of $z_{\\rm SM} \\simeq 1.8$ we find, for $\\varepsilon =0.1$, that the GRB rate peaks at a redshift of $z_{\\rm GRB}\\simeq 3.2$ (Fig.~2 and Tab.~1; see also \\citet{firmani2005} \\item Local massive galaxies, like our Milkyway, are not expected to host long GRBs. The last long GRB in our Galaxy should have occurred several gigayears ago. \\item The global and local GRB to SN ratios appear to be insensitive to the details of the cosmic star formation history, while the redshift of the peak GRB rate can vary appreciably (Tab.~1). \\item The local GRB/core-collapse ratio is much smaller than the one obtained from averaging over the universe; i.e., by one order of magnitude for $\\varepsilon =0.1$ (Tab.~1). \\item We obtain the expected result that the number of massive stars in the universe with a metallicity below a critical value $\\varepsilon$ does roughly scale with $\\varepsilon$. I.e., For $\\varepsilon =0.1$, we find a ratio of low-metallicity ($Z<Z_{\\odot}/10$) to total global supernova rates of 0.16. Further, the probability that a randomly chosen burst has a metallicity of $Z_{\\odot}/100$ is about 10\\%. \\item We derive the most likely redshift for gamma-ray bursts of specified metallicity for unbiased observations. For example, for a metallicity of $Z_{\\odot}/100$ -- as found for GRB050730 --, we find GRBs to occur most likely at redshifts of $z > 6$. Locally, the ratio of GRBs with a metallicity of $Z_{\\odot}/100$ to all GRBs is about 0.02 (Tab.~1). \\item The larger the low-metallicity bias of long GRBs, the less likely can binary scenarios explain the major fraction of them. \\item A confirmation of the low-metallicity bias of long GRBs to values of the order of $\\varepsilon =0.1$ would imply that fast rotation may be much more common at low metallicity among massive stars. \\end{itemize}"
    },
    "0512/astro-ph0512047_arXiv.txt": {
        "abstract": "We use new and archival \\emph{Chandra} and \\emph{ROSAT} data to study the time variability of the X-ray emission from the pulsar wind nebula (PWN) powered by PSR B1509$-$58 on timescales of one week to twelve years.  There is variability in the size, number, and brightness of compact knots appearing within 20$\\arcsec$ of the pulsar, with at least one knot showing a possible outflow velocity of $\\sim0.6c$ (assuming a distance to the source of 5.2 kpc).  The transient nature of these knots may indicate that they are produced by turbulence in the flows surrounding the pulsar.  A previously identified prominent jet extending 12 pc to the southeast of the pulsar increased in brightness by 30\\% over 9 years; apparent outflow of material along this jet is observed with a velocity of $\\sim0.5c$.  However, outflow alone cannot account for the changes in the jet on such short timescales. Magnetohydrodynamic sausage or kink instabilities are feasible explanations for the jet variability with timescale of $\\sim$ 1.3$-$2 years.  An arc structure, located 30$\\arcsec-45\\arcsec$ north of the pulsar, shows transverse structural variations and appears to have moved inward with a velocity of $\\sim0.03c$ over three years.  The overall structure and brightness of the diffuse PWN exterior to this arc and excluding the jet has remained the same over the twelve year span.  The photon indices of the diffuse PWN and possibly the jet steepen with increasing radius, likely indicating synchrotron cooling at X-ray energies. ",
        "introduction": "\\label{sec:int}  The pulsar B1509$-$58 and the supernova remnant (SNR) G320.4$-$1.2 (MSH 15$-$5\\emph{2}) represent one of approximately 20 known associations between a pulsar and a SNR.  This pulsar is one of the most energetic known, with a period ($P$) of 150~ms, a period derivative ($\\dot{P}$) of 1.2$\\times10^{-12}$~s~s$^{-1}$, a characteristic age $\\tau_c \\equiv P/2\\dot{P}\\approx1700$~yr, a spin-down luminosity $\\dot{E} \\equiv 4\\pi^2I\\dot{P}P^{-3} \\approx 1.8\\times10^{37}$~ergs~s$^{-1}$ (for a moment of inertia $I\\equiv10^{45}$~g~cm$^{2}$), and an inferred dipole surface magnetic field $B_p\\approx3.2\\times10^{19}(P\\dot{P})^{1/2}\\approx1.5\\times10^{13}$~G \\citep{kms94,lkg05}.  SNR G320.4$-$1.2 has been well studied at radio, optical, and X-ray wavelengths.  The radio morphology consists of a partial shell to the southeast and a series of bright clumps $\\sim25\\arcmin$ to the northwest \\citep{gbm99} that coincide with the optical nebula RCW~89 \\citep{rcw60,shm83}.  The X-ray morphology consists of a bright, elongated pulsar wind nebula (PWN) with a collimated jet extending $\\sim4\\arcmin$ to the southeast \\citep[hereafter G02]{shm83,gcm95,tmc96,bb97,gak02}.  To the northwest are thermal clumps associated with the radio clumps and RCW~89 \\citep{shm83}.  In addition, \\citetalias{gak02} identified several compact knots close to the pulsar, plus two semicircular arcs at a distance of 17$\\arcsec$ and 30$\\arcsec$ to the north of the pulsar.  The toroidal morphology and collimated jet are reminiscent of structures found in the PWNe powered by the Crab and Vela pulsars \\citep{hss95,wht00,hgh01}. The arcs have been proposed as being due to ion-compression in the particle-dominated equatorial flow from the pulsar and were interpreted by \\citetalias{gak02} as analogs of the ``wisps'' found in the Crab Nebula. SNR G320.4$-$1.2 has recently been observed in very high energy $\\gamma$-rays by HESS.  The emission is elongated along the PWN axis possibly indicating inverse Compton scattering of relativistic electrons \\citep{aaa05}.  The structures in the Crab and Vela PWNe are known to vary in brightness and position over short timescales (days to months).  In the case of the Crab Nebula, the outward moving X-ray and optical wisps (with velocity $v\\sim0.5c$) are thought to mark the PWN termination shock, while the small-scale X-ray and optical knots are thought to identify unstable, quasi-stationary shocks in the pulsar wind \\citep{hmb02}.  Radio wisps, which rarely correspond to optical wisps, develop and move outward at slower velocities ($v\\sim0.3c$) and there are even more slowly moving radio features ($v\\sim10^4$ km s$^{-1}$) farther away from the pulsar \\citep{bhf04}.  Recent observations of the wisps and polar knots in the near infrared indicate brightness variations on time scales as short as 20 minutes \\citep{msw05}. In the case of the Vela PWN, the X-ray arcs also move outward and vary in brightness by up to 30\\% \\citep{pks01}.  The variability of the Vela PWN jet observed in X-rays has been attributed to both kink instabilities, to account for the dramatic shape and brightness changes over the course of days, and sausage instabilities, to account for the relativistically moving ``blobs'' \\citep[$v\\sim0.5c$;][]{ptk03}.  The Crab nebula jet, on the other hand, shows only weak X-ray morphological variations on year-long time scales \\citep{mbp04} and relativistic outflow identified in the optical with $v\\sim0.4c$ \\citep{hmb02}.  The PWN powered by PSR B1509$-$58 represents a unique opportunity to study variability in PWNe.  The physical size of the PWN is approximately 10 and 100 times larger than the Crab and Vela PWNe, respectively.  Given the observed variability in the Crab and Vela PWNe and the distance to PSR B1509$-$58 \\citep[5.2$\\pm1.4$~kpc;][]{gbm99}, \\citetalias{gak02} predicted measurable variability in the PWN of PSR B1509$-$58 on time scales of a few years.  To that end, we obtained new \\emph{Chandra X-ray Observatory} observations of PSR B1509$-$58 and its PWN.  In this paper we compare our new images to existing \\emph{ROSAT} PSPC and HRI images and \\emph{Chandra} ACIS-I images and report on the variations observed over time scales of 1 week to 12 years. ",
        "conclusions": "\\label{sec:con}  Variability is observed in the X-ray PWN of PSR B1509$-$58 on time scales possibly as short as one week and up to twelve years.  Our primary results are as follows:  \\begin{enumerate} \\item{The compact, small-scale knots appearing within 20$\\arcsec$ of the pulsar exhibit transient behavior which may be attributed to turbulence in the flows surrounding the pulsar.  Possible knot motion is indicated with a velocity of $0.6c$.} \\item{Apparent outflow along the jet is observed with velocities of $\\sim0.5c$.  This outflow alone cannot account for the $\\sim$30\\% brightening of the jet between 1991/1992 and 2000.  The Alfv\\'{e}n crossing time for the jet is 1.3$-$2 years, therefore, MHD kink or sausage instabilities can account for the rapid morphological variations and perhaps the partial jet brightening.} \\item{The outer arc has possibly moved inward with a velocity of $0.03c$, however the transverse structural changes seen in the outer arc may account for the apparent motion.  We cannot determine at this time if the outer arc is truly quasi-stationary or if we are witnessing aliasing.} \\item{The diffuse PWN has not evolved significantly in structure or brightness over the 12-year time span.  Using the summed \\emph{Chandra} images, we identify two possible arc structures exterior to the outer arc.} \\item{The photon indices of the diffuse PWN and possibly the jet steepen with increasing radius indicating synchrotron cooling at X-ray energies.} \\end{enumerate}  Although our imaging capabilities have improved substantially since the first optical observations of time variability in the Crab Nebula \\citep{sca69}, our understanding of these variations in PWNe is still quite limited.  For instance, while we expect magnetic fields to play an important dynamical role in jets, and indeed we do see variations on the appropriate Alfv\\'{e}n crossing times, we do not know for certain if MHD instabilities are the root cause of the observed variations.  The arc structures we observe in G320.4$-$1.2 are equally enigmatic.  We do not yet know if they are in steady motion or are quasi-stationary wave phenomena.  The striking changes in the small-scale knots near the pulsar may simply be ``weather,'' diagnosing unimportant details in the PWN flow, or they may indicate important flow structure which is essential to understanding, for instance, diffusion of particles from the equatorial flow to higher latitudes, a loss essential to a post pair shock second-order Fermi acceleration model.  Certainly deeper, and appropriately spaced, X-ray observations will help resolve some issues such as the possible spatial aliasing of the outer arc.  Also, the higher signal-to-noise will provide better constraints on the spatial spectral index variations and allow us to determine if and how much mixing has occurred in the diffuse PWN.  Finally, we are excited by the recent development of relativistic MHD models and we hope that some of the variability we observe here can eventually be observed in those simulations."
    },
    "0512/astro-ph0512053_arXiv.txt": {
        "abstract": "The search for life on extrasolar planets is based on the assumption that one can screen extrasolar planets for habitability spectroscopically. The first space born instruments able to detect as well as characterize extrasolar planets, Darwin and terrestrial planet finder (TPF-I and TPF-C) are scheduled to launch before the end of the next decade. The composition of the planetary surface, atmosphere, and its temperature-pressure profile influence a detectable spectroscopic signal considerably. For future space-based missions it will be crucial to know this influence to interpret the observed signals and detect signatures of life in remotely observed atmospheres. We give an overview of biomarkers in the visible and IR range, corresponding to the TPF-C and TPF-I/DARWIN concepts, respectively. We also give an overview of the evolution of biomarkers over time and its implication for the search for life on extrasolar Earth-like planets. We show that atmospheric features on Earth can provide clues of biological activities for at least 2 billion years. ",
        "introduction": "Indirect methods to detect extrasolar terrestrial planets such as radial velocity, astrometry, and transits only allow getting morphological information. The planet is deduced from its effect on the parent star under survey. To understand the properties of a planet we need to collect photons from the planet directly to determine its physical and chemical properties. Spectral analysis of extrasolar planets is a very young science that only recently achieved first results for hot Jupiters. Some atmospheric components like Na were detected in absorption in the upper atmosphere of HD209458b (\\cite[Charbonneau, Brown, Noyes \\etal\\ (2002)]{Charbonneau02}). The Spitzer telescope recently detected the infrared emission from two transiting hot Jupiters by subtracting the stellar flux when the transiting planet is eclipsed by its star from the combined flux of the star and planet (see \\cite[Deming, Seager, Richardson \\etal\\ (2005)]{Deming05},  \\cite[Charbonneau, Allen, Megeath \\etal\\ (2005)]{Charbonneau05}).  The first space born instruments able to detect as well as characterize extrasolar planets are DARWIN (see e.g.  \\cite[Kaltenegger, Fridlund \\& Karlsson (2005)]{Kaltenegger05proc}), TPF-I and TPF-C (see e.g. \\cite[Borde \\& Traub (2005)]{Borde05}). DARWIN/TPF-I are both based on the concept of a free flyer IR interferometer architecture, and TPF-C is envisioned as a coronagraph in the visible. The strategy to search for biological activity on terrestrial planets is based on the assumption that one can use spectroscopy to screen extrasolar planets for habitability. Basing the search for life on the carbon chemistry assumption allows establishing criteria for habitable planets in terms of their size and distance from their stars ( \\cite[Kasting, Whitmire \\& Reynolds (1993)]{Kasting93}). Note that observations over interstellar distances can only distinguish those planets whose habitability and biological activity is apparent from observations of the reflected or emitted radiation. A crucial factor in interpreting planetary spectra is the point in the evolution of the atmosphere and its biomarkers over time. Concentrating on the evolution of our planet we establish a model for its atmosphere and the detectable biomarkers over its evolution history. Figure~\\ref{fig2} shows that atmospheric features on Earth can provide clues of possible life forms for at least 2 billion years (see \\cite[Kaltenegger, Traub \\& Jucks  (2005)]{Kaltenegger05}).   \\subsection{Atmosphere evolution on Earth} The Earth formed about 4.5 billion years ago. The primitive atmosphere was formed by the release of volatiles from the interior, and/or volatiles delivered during the late bombardment period. This atmosphere was most likely dominated by carbon dioxide, with nitrogen being the second most abundant gas and trace amounts of methane, ammonia, sulphur dioxide, hydrochloric acid and oxygen (see \\cite[Kasting \\& Siefert (2002)]{Kasting02}). Carbon dioxide and/or methane played a crucial role in the development of an early greenhouse effect that counteracted the lower solar output. Two major processes have changed the primitive atmosphere, the reduction of $CO_2$ and the increase of $O_2$. A huge amount of $CO_2$ must have been removed from the atmosphere, most likely by the burial of carbon into carbonate rocks, though the process is still debated. Primitive cyanobacteria are believed to have produced oxygen. After most of the reduced minerals were oxidized, about two billion years ago, atmospheric oxygen could accumulate. Different schemes have been suggested to quantify the rise of oxygen and the evolution of life by anchoring the points in time to fossil finds (see e.g. \\cite[Owen (1980)]{Owen80} \\cite[Schopf (1993)]{Schopf93} \\cite[Ehrenfreund \\& Charnley (2003)]{Ehrenfreund03}). There are still many open questions.  We use a climate model based on work by \\cite[Kasting \\& Catling (2003)]{Kasting03}, \\cite[Pavlov, Hurtgen, Kasting, Arthur (2003)]{Pavlov03}, \\cite[Segura, Krelove, Kasting \\etal\\ (2003)]{Segura03} and \\cite[Traub \\& Jucks (2002)]{Traub02} to create a schematic atmospheric model of our Earth over geological timescales. We find a surface temperature above the freezing point for all of Earth's history (not considering short term atmospheric changes during glaciation periods). Our radiative transfer model is based on the model by \\cite[Traub \\& Jucks (2002)]{Traub02} that was developed to successfully analyse data from stratospheric balloon flight campaigns ongoing during the last few years. The same radiative transfer code was used to successfully analyse data from Earthshine measurements in the visible (400 to 900 nm) (see \\cite[Woolf, Smith, Traub \\etal\\ ]{Woolf02}) and mid-infrared (700 to 2400 nm) (see \\cite[Turnbull, Traub, Jucks \\etal\\ (2005)]{Turnbull05}). It shows excellent agreement with measurements in both the visible and IR wavelength band. We present the spectra of Earth over geological timescale in anticipation of proposed space based terrestrial planet search missions to operate between 500 to 1100 nm in the visible to mid-IR range and 6 to 20$\\mu$m in the IR. Our model shows a considerable difference in biosignatures that can be detected with those missions. That should be able to constrain an Earth-like planet to its time in evolution.   \\subsection{Atmospheres and Biomarkers} The spectral characteristics of a planet are a very important guide to identifying the best wavelength region to probe for a planet as well as characterize the spectra once it has been detected. The TPF Science Working Group (TPF-SWG) identified the waveband between 8.5$\\mu$m to 20$\\mu$m and preferably 7$\\mu$m to 25$\\mu$m, respectively, for the search for biomarkers in the mid-IR region and 0.7$\\mu$m to 1.0$\\mu$m and preferably 0.5$\\mu$m to 1.1$\\mu$m for the visible to near IR region (see \\cite[Des Marais, Harwit, Jucks, \\etal\\ (2002)]{DesMarais02}). The DARWIN project concentrates on the waveband between 6$\\mu$m to 20$\\mu$m (see \\cite[Fridlund, Kaltenegger, \\etal\\ (2006)]{Fridlund06}). In the thermal part of the spectrum, the shape gives a measure of the temperature of the object examined. Observations from 8$\\mu$m to 12$\\mu$m of the $H_2O$ continuum allow estimations of the surface temperature of Earth-like planets. For higher concentrations of water vapor in the atmosphere, the continuum gets lowered due to water absorption.  Biomarkers are features whose presence or abundance requires a biological origin (see \\cite[Des Marais, Harwit, Jucks, \\etal\\ (2002)]{DesMarais02},  \\cite[Meadows \\etal\\ (2005)]{Meadows05}). They are created either during the acquisition of energy and/or the chemical ingredients necessary for biosynthesis (e.g. atmospheric oxygen and methane) or are products of the biosynthesis (e.g. complex organic molecules and cells). As signs of life in themselves $H_2O$ and $CO_2$ are secondary in importance because although they are raw materials for life, they are not unambiguous indicators of its presence. Taken together with molecular oxygen, abundant $CH_4$ can indicate biological processes. Depending on the degree of oxidation of a planet's crust and upper mantel non-biological mechanisms can also produce large amounts of $CH_4$ under certain circumstances. Oxygen is a chemically reactive gas. Reduced gases and oxygen have to be produced concurrently to be detectable in the atmosphere, as they react rapidly with each other. The 9.6$\\mu$m $O_3$ band is highly saturated and is thus a poor quantitative indicator, but an excellent qualitative indicator for the existence of even traces of $O_2$. Ozone is a very nonlinear indicator of $O_2$ because the ozone column depth remains nearly constant as $O_2$ increases from 0.01 present atmosphere level (PAL) to 1 PAL.  $N_2O$ is a very interesting molecule because it is produced in abundance by life but only in trace amounts by natural processes. There are no $N_2O$ features in the visible and two weak $N_2O$ features in the IR. In the IR it can only be detected in regions strongly overlapped by $CH_4$, $CO_2$ and $H_2O$, so it is unlikely to become a prime target for the first generation of space-based missions searching for extrasolar planets that will work with low resolution, but it is an excellent target for follow up missions. Spectral features of $N_2O$ become apparent in atmospheres with less $H_2O$ vapor. In the mid-IR $N_2O$ has a band at 7.9$\\mu$m, comparable in strength to the adjacent $CH_4$ band but weak compared with the overlapping $H_2O$ band. The absorption bands of those three species are different. It is not readily separable for low resolution spectroscopy for current Earth, but the methane feature is easily detectable for early type Earth planets according to our models. To detect and study surface properties we can only use wavelengths that penetrate to the planetary surface. On a cloud-free Earth, the diurnal flux variation caused by different surface features rotating in and out of view could be high. When the planet is only partially illuminated, more concentrated signal from surface features could be detected as they rotate in and out of view. Most surface features like ice or sand show very small or very smooth continuous opacity changes with wavelength. Figure~\\ref{fig1} shows that the signal detected from a cloud free planet is lower than that from an Earth-like planet. Thus the integration time needed to detect such planets will be longer.   \\subsection{The red edge of land planets} An interesting example for surface biomarkers on Earth is the red edge signature from photosynthetic plants at about 750 nm. Photosynthetic plants have developed strong infrared reflection as a cooling mechanism to prevent overheating and chlorophyll degradation. The primary molecules that absorb the energy and convert it to drive photosynthesis ($H_2O$ and $CO_2$ into sugars and $O_2$) are chlorophyll A (0.450$\\mu$m) and B (0.680$\\mu$m). Several groups have measured the integrated Earth spectrum via the technique of Earthshine, sunlight reflected from the ``dark'' side of the moon. Earthshine measurements have shown that detection of Earth's vegetation-red edge (VRE) is feasible but made difficult owing to its broad, essentially featureless spectrum and cloud coverage. Trying to identify such weak, continuum-like features at unknown wavelengths in an extrasolar planet spectrum requires models for different planetary conditions. Our knowledge of different surface reflectivities on Earth, like deserts, ocean and ice, help assigning the VRE of the Earthshine spectrum to terrestrial vegetation (see Figure~\\ref{fig2}). Earth's hemispherically integrated vegetation red-edge signature is weak, but planets with different rotation rates, obliquities, land-ocean fraction, and continental arrangement may have lower cloud-cover and higher vegetated fraction (see e.g. \\cite[Seager (2002)]{Seager02}).  \\subsection{Model results}  \\begin{figure} \\centering \\resizebox{4.6cm}{!}{\\includegraphics{fig01BMa.ps} } \\resizebox{4.75cm}{!}{\\includegraphics{fig01BMb.ps} } \\caption{(left) Signal of a current Earth atmosphere considering surface features assuming a clear atmosphere without clouds; Different surface composition (a) snow (b) sea (c) trees (d) sand. (right) Signal of current Earth with realistic surface features and (a) no clouds, (b) 100\\% cumulus cloud coverage at 4 km and (c) 100\\% cirrus cloud coverage at 12 km and (d) a mixture of clouds resembling the present Earth. }\\label{fig1} \\end{figure}  Our calculations (see \\cite[Kaltenegger, Traub \\& Jucks  (2005)]{Kaltenegger05}) show that the composition of the surface (especially in the visible), the atmospheric composition and temperature-pressure profile can all have a significant influence on the detectability of a signal. The reflectivity of the Earth has not been static throughout the past 4.5Ga (Ga = $10^9$ years ago). Oxygen and ozone became abundant about roughly 2.3Ga, affecting the atmospheric absorption component of the reflection spectrum. About 2Ga, a green phytoplankton signal developed in the oceans and about 0.44Ga, an extensive land plant cover developed, generating the red chlorophyll edge in the reflection spectrum. The oxygen and ozone absorption features could have been used to derive the presence of biological activity on Earth anytime during the past 50\\% of the age of the solar system, while the chlorophyll red-edge reflection feature evolved during the most recent 10\\% of the age of the solar system. Figure~\\ref{fig1} shows the spectral signature of an Earth atmosphere assuming different surface composition and the big impact clouds can have on the spectrum. From the 4 surfaces plotted on the left one sees that an Earth-size ocean planet will be more difficult to detect than a snow covered planet due to its low albedo. The smooth lines in the left panel of Figure~\\ref{fig1} show the albedo of the surfaces without overlayed atmosphere. The panel on the right in Figure~\\ref{fig1} shows the strong impact clouds can have on the detected signal. Present Earth has about 60\\% cloud coverage. 100\\% cumulus (low) and 100\\% cirrus (high) cloud coverage, respectively, are shown for reference. A planet without clouds or snow has a lower overall albedo and thus requires a longer integration time for detection than a planet with clouds. Figure~\\ref{fig2} shows Earth's reflected and thermal emission spectra respectively with its major molecular species ($H_2O$, $O_3$, $O_2$, $CH_4$, $CO_2$, $N_2O$) as well as minor contributors ($H_2S$, S$O_2$, $NH_3$, $SF_6$, CFC-11, CFC-12). The dark lines show a resolution of 70 in the visible and 25 in the IR, as proposed for the TPF-C and Darwin/TPF-I mission, respectively. The atmospheric features on an Earth-like planet change considerably over its evolution from a current day atmosphere (epoch5 = present Earth) over a $CO_2$/$CH_4$-rich (epoch 3 around 2Ga) to a $CO_2$ rich (3.9Ga = epoch 0) atmosphere. $O_3$ shows a strong feature in the IR for the last 2Ga of Earth's history even at a low resolution of 25. $O_2$ shows a weaker feature in the visible for the last 2Ga of Earth's history at a resolution of 70 (note that no noise level is assumed for these calculations).  In the mid-IR the classical signatures of biological activity are the 9.6$\\mu$m $O_3$ band, the 15$\\mu$m $CO_2$ band and the 6.3$\\mu$m $H_2O$ band or its rotational band that extends from 12$\\mu$m out into the microwave region (\\cite[Selsis, Despois, Parisot (2002)]{Selsis02}). For the visible, a similiar investigation has not been conducted yet. In the same spectral region, the 7.7$\\mu$m band of $CH_4$ is a potential biomarker for early-Earth type planets. In the visible to near-IR one can see a strong $O_2$ absorption feature at 0.76$\\mu$m, a broadband $O_3$ absorption at 0.45$\\mu$m to 0.75$\\mu$m and a strong $H_2O$ band at 0.94$\\mu$m. The strongest $O_2$ feature is the saturated Fraunhofer A-band at 0.76$\\mu$m that is still relatively strong for significantly smaller mixing ratios than present Earth's. $O_3$ has two broad features, the extremely strong Huggins band in the UV shortward of 0.33$\\mu$m and the Chappius band which shows as a broad triangular dip in the middle of the visible spectrum from about 0.45$\\mu$m to 0.74$\\mu$m. Methane at present terrestrial abundance (1.65ppm) has no significant visible absorption feature, but at high abundance it has strong visible bands at 0.88$\\mu$m and 1.0$\\mu$m, readily detectable in early Earth's history. $CO_2$ has negligible visible features at present abundance, but in a high $CO_2$-atmosphere of 10\\% it would have a significant band at 1.2$\\mu$m and even stronger ones at longer wavelengths. Especially in the early evolution stage, the weak 1.06$\\mu$m band can be observed (epoch 0).   \\begin{figure} \\centering \\resizebox{5.7cm}{!}{\\includegraphics{fig02BMa.ps} } \\resizebox{5.7cm}{!}{\\includegraphics{fig02BMb.ps} } \\caption{The atmospheric features on an Earth-like planet change considerably over its evolution from a current day atmosphere (epoch 5 = present Earth) over a $CO_2$/$CH_4$-rich atmosphere (epoch 3) to a $CO_2$ rich (3.9Ga ago = epoch 0) in both the visible (left) and the IR range (right). The black lines show spectral resolution of 70 for the visible and 25 for the IR, comparable to the proposed mission concepts. }\\label{fig2} \\end{figure} ",
        "conclusions": "\\label{sec:concl} Concentrating on the evolution of our planet we established a model for its atmosphere and the detectable biomarkers over its evolution history and showed the results for TPF-I/Darwin and TPF-C. These missions should be able to constrain an Earth-like planet in its evolution. Our calculations show that the atmospheric and the surface composition (especially in the visible) and cloud coverage influence the detected signal considerably (see Figure~\\ref{fig1}).  From the 4 surfaces plotted on the left one sees that e.g. an Earth-size ocean planet will be more difficult to detect than a snow covered planet due to its low albedo (here we consider no clouds). The atmospheric features on an Earth-like planet change considerably over its evolution from a current day atmosphere (epoch 5 = present Earth) over a $CO_2$/$CH_4$-rich atmosphere (epoch 3 around 2Ga) to a $CO_2$ rich (3.9Ga = epoch 0) see Figure~\\ref{fig2}. Atmospheric features on Earth can provide clues of biological activities for at least 2 billion years. $O_3$ shows a strong feature in the IR for the last 2Ga of Earth's history even at a low resolution of 25. $O_2$ shows a weaker feature in the visible for the last 2Ga of Earth's history at a resolution of 70 (note that no noise level is assumed for these calculations). Methane (a potential bioindicator for early Earth) and nitrous oxide (a biomarker with a weak signature in the IR) have features nearly overlapping in the 7$\\mu$m region, lying in the red wing of the 6$\\mu$m water band. The absorption bands of those three species are different. It is not readily separable for current Earth at low spectral resolution. $N_2O$ is a very interesting molecule because it is produced in abundance by life but only in trace amounts by natural processes. It is unlikely to become a prime target for the first generation of space-based missions searching for extrasolar planets that will work with low resolution, but it is an excellent target for follow up missions. But if methane is abundant, as it probably was for early Earth, then it is readily detectable, even at low spectral resolution, functioning as an early bio-indicator for young Earth-like planets.  \\vspace{-5pt}"
    },
    "0512/astro-ph0512579_arXiv.txt": {
        "abstract": "It is estimated that the Weibel instability is not generally an effective mechanism for generating ultrarelativistic astrophysical shocks. Even if the upstream magnetic field is as low as in the interstellar medium, the shock is mediated not by the Weibel instability but by the Larmor rotation of protons in the background magnetic field. Future simulations should be able to verify or falsify our conclusion. ",
        "introduction": "There is large literature on gamma-ray burst (GRB) afterglows based on the assumption that the X-ray, optical and radio afterglows are the synchrotron emission from relativistic electrons Fermi accelerated at the forward shock of the blast wave (see recent reviews by M\\'{e}sz\\'{a}ros 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2004). A longstanding difficulty with this assumption has been the inferred magnetic field needed to fit the afterglow data typically requires that the magnetic energy density exceeds by many orders of magnitude that which would be expected from the shock compression of the interstellar magnetic field of the host galaxy (Gruzinov \\& Waxman 1999; Gruzinov 2001). Many authors have therefore assumed that the shock somehow manufactures field energy to meet this requirement, but no convincing mechanism has been proposed to date. One mechanism discussed is the Weibel instability (Medvedev \\& Loeb 1999, Silva et al. 2003; Frederiksen et al. 2004; Jaroschek et al. 2005; Medvedev et al. 2005; Kato 2005; Nishikawa et al. 2005), which has the fastest growth rate and produces relatively strong small scale magnetic field even in an initially non-magnetized plasma. It is expected that the thermalization of the upstream flow could occur via scattering of particles on the magnetic fluctuations. In the electron-positron plasma, the instability does generate the magnetic field at about 10\\% of the equipartition level and does provide the shock transition at the scale of a dozen of electron inertial length (Spitkovsky 2005). However, simulations of colliding electron-proton flows show that while the electrons are readily isotropized, the protons acquire only small scattering in angles after passing the simulation box (Frederiksen et al. 2004). How long the field persists after the shock is also an important question (Gruzinov 2001) but here we discuss whether the Weibel instability can even cause the shock in the first place.  Moiseev \\& Sagdeev (1963, see also Sagdeev 1966) analyzed the structure of the nonrelativistic Weibel driven shock and found that the width of the shock transition should be very large because electrons easily screen proton currents thus suppressing development of the instability. Failure of the Weibel instability to preempt other shock mechanisms, except for very large Alfven Mach numbers, has been discussed in the context of nonrelativistic shocks by Blandford and Eichler (1987).  Here we estimate the width of the Weibel driven shock in the ultra-relativistic electron-proton plasma. Although based on a number of physical assumptions  about the behavior of the plasma parameters at the non-linear stage of the Weibel instability, such analytical scalings are necessary in any case, because, by evident reasons, simulations of plasmas are possible only with artificially low proton-to-electron mass ratios (e.g., Frederiksen et al. (2004) took $m_p/m_e=16$). In this paper, we present the parameters in physically motivated dimensionless form and we believe that our assumptions could be checked by numerical simulations. Only by combining numerical simulations with analytical scalings can we achieve reliable conclusions about the properties of real shocks.  As a model for the shock formation, we consider collision of two oppositely directed plasma flows. Eventually two diverging shocks should be formed with plasma at rest between them. However at the initial stage, the two flows interpenetrate each other exciting turbulent electro-magnetic fields. Particles are eventually thermalized by scattering off these turbulent fields. As electrons are thermalized relatively rapidly, we consider development of the Weibel instability in two proton beams propagating through relativistically hot isotropic electron gas. We estimate the proton isotropization length in such a system and conclude that the Weibel-mediated shocks are so wide that even in the interstellar medium, the shock should be formed at the scale of the Larmor radius of the proton in the background magnetic field.  The article is organized as follows. In sect.2, we find the growth rate of the proton Weibel instability. Saturation of the instability is considered in sect.3. In sect.4, we exploit the obtained results in order to estimate the width of the Weibel-mediated shock transition. Section 5 contains the discussion. ",
        "conclusions": "The estimate (\\ref{shockwidth}) was obtained under the assumption that there is no magnetic field in the upstream flow. If the flow is magnetized, a shock transition may be formed at the scale of the proton Larmor radius; therefore the above estimates are valid only if $eB_0L/m_p\\gamma<1$, where $B_0$ is the magnetic field in the upstream flow. One can conveniently characterize the magnetization of the flow by the parameter $\\sigma =B_0^2/(4\\pi m_p\\gamma n)$. Making use of Eq.(\\ref{shockwidth}), one finds that the shock may be driven by the Weibel instability if \\begin{equation} \\sigma<\\xi^4\\zeta^2\\left(\\frac{4\\tau m_e}{3\\pi m_p}\\right)^3=1.5\\times 10^{-11}\\xi^4\\tau^3\\zeta^2.\\label{magnetization} \\end{equation} If the shock propagates through the interstellar medium, the magnetization exceeds the right-hand side of Eq.(\\ref{magnetization}) by factor about 30: $\\sigma=5\\cdot 10^{-10}B^2_{-5.5}n^{-1}$, where $B=10^{-5.5}B_{-5.5}$ G is the interstellar magnetic field, $n$ cm$^{-3}$ the gas number density. This suggests that forward shocks that presumably produce GRB afterglows are mediated not by the Weibel instability but by the Larmor rotation of protons in the background magnetic field. This does not mean that the Weibel instability does not work at all. On the contrary, it does develop and may well create some small scale magnetic field  that is much stronger than the background field. The scattering off these magnetic fluctuations results in diffusion of the protons in angles. However, because the scattering does not manage  to isotropize the protons at the scale less than the Larmor radius it is hard to see how the Weibel instability could convert the kinetic energy to another forms. The strong dependence of the estimate (\\ref{maznetization}) on the parameters $\\xi$, $\\tau$, and $\\zeta$ makes accurate determination of their values, presumably by simulations, crucial to solidify this conclusion.  The estimate (\\ref{magnetization}) shows that a fraction $\\epsilon_B\\sim 10^{-4}$ of the total energy is converted into the magnetic energy unless the electrons are heated additionally within the shock structure. The generated small-scale field should decay (Gruzinov 2001; Milosavlevi\\'c \\& Nakar 2005) so that $\\epsilon_B$ may be even lower. According to Panaitescu \\& Kumar (2002) and Yost et al. (2003), the observed spectra and light curves of the GRB afterglows imply $\\epsilon_B\\sim 10^{-3}\\div 10^{-1}$ in most cases. Eichler \\& Waxman (2005) demonstrated that the above estimates may be rescaled such that the observations are fitted with values of $\\epsilon_B$ that are smaller by an arbitrary factor $f$, $m_e/m_p<f<1$. Taking this into account one can see that the Weibel instability could provide the necessary field unless the field decay is too strong. On the other hand, as the global structure of the shock transition is dictated by the Larmor rotation of the protons in the background field, some new physics could come into play.  A presumably important physical mechanism is the synchrotron maser instability at the shock front. This instability generates semi-coherent, low-frequency electromagnetic waves (Gallant et al. 1992; Hoshino et al. 1992; Lyubarsky in preparation). In low magnetized flows, the amplitude of these waves exceeds the amplitude of the shock compressed background field. In this case, relativistic particles radiate in the field of the waves via nonlinear Compton scattering (e.g., Melrose 1980, pp. 136-141). The power and characteristic frequencies of this emission are similar to those for synchrotron emission in the magnetic field of the same strength therefore the observed gamma ray bursts afterglows may be attributed to the nonlinear Compton scattering off the electromagnetic waves generated by the synchrotron maser instability at the shock front. It is beyond the scope of this paper to redo afterglow theory with the spectrum of low frequency electromagnetic waves that is expected from this instability.  We thank J.Frederiksen, A. Spitkovsky and R.Blandford for helpful discussions. Y.L. acknowledges support from the German-Israeli Foundation for Scientific Research and Development. % D.E. acknowledges the hospitality of the Kavli Institute for Particle Astrophysics and support from the Israel-U.S. Binational Science Foundation, the Israeli Science Foundation, and the Arnow Chair of Physics."
    },
    "0512/astro-ph0512265_arXiv.txt": {
        "abstract": "We report new limits on cosmic neutrino fluxes from the test flight of the Antarctic Impulsive Transient Antenna (ANITA) experiment, which completed an 18.4 day flight of a prototype long-duration balloon payload, called {\\em ANITA-lite}, in early 2004. We search for impulsive events that could be associated with ultra-high energy neutrino interactions in the ice, and derive limits that constrain several models for ultra-high energy neutrino fluxes and rule out the long-standing Z-burst model. We set a 90\\% CL integral flux limit on a pure $E^{-2}$ spectrum for the energy range $10^{18.5}$~eV$\\leq E_{\\nu} \\leq 10^{23.5}$~eV at $E_{\\nu}^2 F \\leq 1.6 \\times 10^{-6}$ GeV cm$^{-2}$ s$^{-1}$ sr$^{-1}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512323_arXiv.txt": {
        "abstract": "A beginning end-point for galaxy motions within the 10-Mpc Local Volume is constructed by requiring a smooth distribution of (luminous) matter at the time of recombination, which is shown to be equivalent to a smooth Hubble flow at early times.  It is found, by this purely kinematical method, that present peculiar motions are too small by a factor of at least several (and largely in the wrong direction) to have produced the observed structures within the age of the universe. Known dynamical effects are inadequate to remove the discrepancy. This result is different in origin from previously known ``cold flow'' problems. The simple dynamical picture often used within the Local Volume (for instance, in deriving masses through calculation of a zero-velocity surface) is thus called into question. The most straightforward explanation (though not the only possible) is that there exists a large quantity of baryonic matter in this region so far undetected, and unassociated with galaxies or groups. ",
        "introduction": "Matter was born smooth, and is everywhere in clumps.  Departures from homogeneity in the CMB are tiny, on the order of a few dozen parts per million; yet the Local Volume, the region within about ten megaparsecs, is full of structure (see Figure (\\ref{volume})).  Apart from the galaxies themselves, there are groups of galaxies as well as an overall flattening into the Supergalactic Plane.  \\begin{figure} \\plotone{f1.ps} \\caption{A representation of the Local Volume, looking in the Supergalactic Plane toward $L = 90\\arcdeg$.  Galaxies are shown as spheres with volume proportional to luminosity.  In the center is the Local Group; to the left Centaurus, to the right IC 342 and M81/2. The Supergalactic Plane itself is evident, as is the tendency for galaxies to collect in groups. Dotted lines are at 5 Mpc intervals.} \\label{volume} \\end{figure}  The process by which a (nearly) smooth distribution of matter develops present-day structure is complicated, engaging the efforts of many people (and computers). But endpoints, if they can be usefully fixed, by themselves tell something about the process on the average.  This study constructs a beginning endpoint for the galaxies in the Local Volume and investigates the consequences of it. (Structure formation by gravitational instability of small perturbations actually starts at recombination, many years after the Big Bang itself; but for our purposes it's close enough that it may be placed at $t = 0$.) ",
        "conclusions": "Imposing an end condition on the visible matter in the Local Volume has revealed a problem with the currently-visible structures: they are moving too slowly and often in the wrong direction for the observed peculiar velocities to have produced them from a smooth distribution at the time of recombination. Setting limits (rather than performing detailed calculations), it has been found that no known dynamical process explains the deficit.  The most straightforward explanation is that much or most of the baryonic matter in the Volume is in the form of ionized gas, or for some other reason not now presently detectable; and that it is distributed largely in the places where galaxies are not.  In that case it might be observed soon, as X-ray technology improves.  This is not the only possible explanation, however.  Some previously unsuspected way of modifying peculiar velocities might be at work, slowing everything down greatly.  This solution, though, is only speculative.  Complicating the interpretation of this result is the fact that the Local Volume is both dynamically unusual and the only region which may be examined in this detail.  It is possible that any volume in which the peculiar velocity field can be determined to an accuracy of, say, 20 km s$^{-1}$ and sampled to scales of well under one Mpc, will show this kind of anomaly.  It is also possible that something is happening here that does not occur in most of the universe.  It is certain, however, that the assumption of simple trajectories for galaxies within the Volume, freely employed up to this time, must be reexamined.  It might indeed be a reasonable approximation on some scales and in some cases, but subsequent use of it must be justified."
    },
    "0512/astro-ph0512359_arXiv.txt": {
        "abstract": "The detectors designed for Gaia, the next ESA space astrometry mission to be launched in 2011, will allow to observe repeatedly stars very close to Jupiter's limb. This will open a unique opportunity to test General Relativity by performing many Eddington-like experiments through the comparison between the pattern of a starfield observed with or without Jupiter. We have derived the main formulas relevant for the monopole and quadrupole light deflection by an oblate planet and developed a simulator to investigate the processing of the Gaia astrometric observation in the vicinity of the planet. The results show that such an experiment carried out with the Gaia data will provide a new fully independent determination of $\\gamma$ by means of differential astrometric measurements and, more importantly, for the first time will evidence the bending effect due to the quadrupole moment with a 3$\\sigma$ confidence level. Given the accuracy of the experiment for the monopole deflection, this will permit to test alternative modelling of the light bending by moving masses. ",
        "introduction": "\\label{section-1}   Several key questions of modern astrophysics regarding the formation and evolution of the Milky Way will be clarified with the next space astrometry  mission Gaia, approved in 2000 as a cornerstone within the European Space Agency science program~\\cite{gaia-book}. The mission is now funded and will enter in the C/D phase in 2006 for a launch scheduled in late-2011 and nominal operations over the next five years. Thanks to its high precision (10 $\\mu$arcsecond in angular measurements) and multi-epoch astrometry, Gaia will be able to detect the relative positional change of a star resulting for the tiny curvature of the light ray brought about by the gravitational pull of the Sun, and to a lesser extent by the giant planets.  As far as the light deflection produced by the solar mass is concerned, the science case has included from the outset an investigation of the value of the PPN parameter $\\gamma$, considered as a global unknown in the astrometric model. This parameter, equal to one in General Relativity (GR), indicates the amount of space-time curvature produced by a unit rest-mass and it is a fundamental part of the so called parameterized Post-Newtonian (PPN) formalism~\\cite{will}, originally developed by Eddington for the famous experiment during the solar eclipse in 1919. The PPN method aims to quantify the violations of the Equivalence Principle by parameterizing deviations which modify both local physical laws and large-scale gravitational phenomena, including the constancy of the constants. PPN formalism is valid for a broad class of metric theories and includes GR as a reference case, {\\it i.e.} as a current standard theory of gravitation. The slow motion and weak field limit allows to use PPN metric expansion as a function of several parameters, varying from theory to theory, useful to discern gravitational experiments in the Solar System.  The PPN parameter $\\gamma$ is the most important, since all the other parameters to all relativistic order within the alternative theories to gravity, converge to their GR value in proportion to $|1-\\gamma|$~\\cite{tur}.  The most general formulation of these theories contains an arbitrary function of a scalar field coupled to the stress-energy tensor in order to merge quantum mechanics with gravity.  Recent works on scalar-tensorial theories consistent with various cosmological scenarios suggest the search for discrepancies from unity of $\\gamma$ at the levels of $10^{-5}$ to $10^{-7}$; more precisely, Damour and Nordtvedt~\\cite{dam-nor} assume that there exist a cosmological attractor mechanism towards GR quantified from deviation of $\\gamma$ from unity within the above level of accuracy depending on the total mass density of the Universe. Then, an experiment purposely designed  could reveal the coexistence of the two mentioned fields throughout the cosmological evolution up to the present where pure tensor gravity explains the gravitational interaction from small to large scale. In addition, some dynamical models predict scalar fields dominating the current energy density of the universe, which should contribute to the value of $\\gamma$ at the level of $10^{-7}$ to $10^{-9}$, thus providing a universe evolution without the need of dark matter~\\cite{tur}. All this issues make precision tests of gravity in space a fascinating challenge for the next decades, as in the past the Eddington's observations in 1919 of star line-of-sight did, confirming the amount of 1\\hbox{$.\\!\\!^{\\prime\\prime}$}72 as deflection angle predicted by GR.  Nowadays, all present experimental tests are compatible with the predictions of GR. In particular, experiments conducted in Solar System have tested all the weak-field predictions of GR theory at better than the $10^{-3}$ level and down to $2\\times 10^{-5}$ level for $\\gamma$ as seen from the additional Doppler shift in radio-wave beams connecting the Earth to the Cassini spacecraft when they passed near the Sun~\\cite{bertotti-gamma}.  New space projects will go deeper in experimental gravitation. The NASA Gravity Probe B mission~\\cite{gpb}, whose first results are expected very soon, should directly measure $\\gamma$ to better than the $10^{-5}$ level, although this is not the main objective of the experiment and remains subjected to efficient instrumental calibrations. An early analysis of the Gaia's capabilities, indicates that its measurements should provide a precision of $5\\times 10^{-7}$ for this parameter~\\cite{mignard-gamma,ppn-ramod}, an improvement of more than two orders of magnitude to the current best estimate mentioned above.  This is quite comparable to the expectations of the LATOR project~\\cite{lator}, which is designed to use two optical interferometers between two micro-spacecrafts and aims to determine $|1-\\gamma|$ at the level of $10^{-8}$.  Besides the  determination of $\\gamma$  based on the global astrometry capabilities of Gaia, the mission will also allow to carry out dedicated small field experiments from the observations performed close to the surface of a giant planet like Jupiter or Saturn. The light bending due to their gravitational field will be detectable in the visible and we show in this paper that even the quadrupole field of Jupiter will be evidenced in such an experiment. It is therefore possible to carry out the same kind of eclipse experiment made by Dyson, Eddington, and Davidson in 1920, but now by comparing stellar positions in the immediate vicinity of the planets. Not only will this experiment be a further confirmation of a GR prediction, but it will also help understand the very difficult question of the light-bending by a moving source and that of the propagation of gravitation.  Using Jupiter to test GR, Kopeikin and Fomalont claimed they have determined the speed of gravity trough the VLBI~\\footnote{VLBI is the acronym for Very Long Baseline Interferometry} astrometric measurement of Shapiro time delay of light from the Quasar J0842+1835, as its image passed within 3.7 arcmin of the planet (about 10 Jupiter radii, a far cry from what Gaia can achieve)~\\cite{fom-kop}. According to the post-Minkowskian formalism assumed~\\cite{kop}, the gravitational field due to a moving body propagates with finite speed; consequently, it influences a photon with some retardation.  The speed of gravity enters as an extra tiny velocity-dependent corrections to the Shapiro time delay formula of the order of 4.8 ps. At present there is no general consensus on the results, especially upon the method to extend GR to the case where the velocity of gravity differs from the velocity of light~\\cite{carlip,will-c}. In particular, given the observed limits on $\\gamma$, the Fomalont-Kopeikin measurements are not accurate enough yet to determine the speed of gravity~\\cite{carlip}. Some authors interpret the result within the Lorentz transformation properties of the weak gravitational field and consider it as a violation not only of the Lorentz invariance but also of the Galilean one~\\cite{samuel,asada}. Paper~\\cite{frittelli} settles the significance of the Jupiter experiment as a simply standard aberration of light which propagates across an active medium with an effective index of refraction induced by the gravitational field of a lens in motion: the tiny corrections to the Shapiro time delay should be included into the fitting of microlensing events. However, even if there is still a controversy about the question of whether the results depend on the speed of gravity~\\cite{kop-cqg} or on the abberation of light~\\cite{asada,frittelli,sanchez}, there is a general agreement that this measurement is a remarkable precise VLBI detection of a post-Newtonian effect due to a planet.  The possibility to apply methods of relativistic astrometry to test GR has  prompted  several groups to model highly accurate angular observations in the relativistic framework. New formalisms have been proposed (~\\cite{klio, kop, ram-3,lp} and reference therein) to tackle the relativistic problem of the light path reconstruction, which, at least at the $\\mu$as accuracy, implies taking into account all the contributions not only due to the bulk mass of the Solar System bodies, but also to their quadrupole moment and to the time-dependent gravitational field generated by their motions (see \\tref{table-1}).  \\begin{table} \\setlength{\\tabcolsep}{8mm} \\begin{center} \\begin{tabular}{cr@{} lr@{}  lc} \\hline \\\\[-4pt] & \\multicolumn{2}{c}{$\\delta\\Phi_\\textrm{pN}$} & \\multicolumn{2}{c}{$\\delta\\Phi_{Q}$} & $\\delta\\Phi_\\textrm{max}$ \\\\[5mm] \\hline \\\\[-4pt] Sun      & $1.\\!\\!''75$ &  & $\\sim 1$~$\\mu$as  &   &$(180^\\circ)$  \\\\[3mm] &   $\\mu$as   &  & $\\mu$as            &    &              \\\\ Mercury  &         $83$ &  &   --              &   &   $9'$ \\\\ Venus    &        $493$ &  &   --              &   &   $4.5^\\circ $ \\\\ Earth    &        $574$ &  &   0.6             &   &$178^\\circ$ \\\\ Moon     &         $26$ &  &   --              &   &$9^\\circ$ \\\\ Mars     &        $116$ &  &  0.2              &   &$25'$ \\\\ Jupiter  &      $16280$ &  &        239        &   &$90^\\circ/3'$ \\\\ Saturn   &       $5772$ &  &           94      &   &$18^\\circ/51''$ \\\\ Uranus   &       $2081$ &  &           25      &   &$72'/6''$ \\\\ Neptune  &       $2535$ &  &            9      &  & $51'/3''$ \\\\ Pluto    &          $7$ &  &      --           &   &$8''$ \\\\ \\hline \\end{tabular} \\caption{\\label{table-1} Light deflection amounts at 1~$\\mu$arcsec due to the planets for a photon crossing the Solar System: pN is for post-Newtonian order and Q for the quadrupolar moment. The values are computed for grazing rays in the Gaia observing geometry; figures in the last column give the angular distances between the perturbing body and the star at which the effect is still 1~$\\mu$arcsecond, with Gaia at the Sun-Earth L2. Where two values are reported they refer, respectively, to pN and quadrupole effect.} \\end{center} \\end{table}  In particular, two papers~\\cite{kop-mon} and~\\cite{lp-mon} have gone through the post-Newtonian treatment of the light propagation in the case of an isolated axisymmetric body, in order to take into account the influence of the  multipole moments. The former is based on a mathematical solution of geodesic equation whereby mathematical expressions for the light deflection in any order of multipole perturbations are obtained. In the second paper, this contribution stems from the prior computation of the time transfer function between two point located at a finite distance. However neither approach has provided hints on how to apply the intricate formulas to a practical case like the one faced in the Gaia modelling.  Having in mind the capabilities of Gaia to observe stellar sources very close to Jupiter's limb (at a fraction of the planet radius), we have evaluated the light deflection produced by an oblate planet on grazing photons coming from distant stars. This study is part of a  wider project called GAia Relativistic EXperiment (GAREX), which aims to investigate all possibilities to test GR with Gaia measurements. In the solar system, this mission will mainly carry out: \\begin{itemize} \\item light deflection experiments, divided into (i) global astrometry, in particular highly accurate determinations of $\\gamma$  by observing the change in  star position at different angular distance to the Sun; (ii) small field experiments, with the examination of the light propagation by  means of differential measurements of stellar positions near the planets; \\item perihelion precession effect, related to the determinations of PPN parameter $\\beta$ \\footnote{The PPN parameter $\\beta$ measures the non-linearity in the superposition of gravitational fields.} from the orbit fitting of several thousands  minor bodies~\\cite{hestro}. \\end{itemize} In this paper we concentrate on the case of the light bending by Jupiter at the microarcsecond level and its detection by the small field astrometric observations with Gaia. In~\\cite{poster-paris} we have already  explored the favorable circumstances for Gaia in five years of continuous observations, approximately from 2012 to 2017, to detect the quadrupole light bending effect predicted by GR, but never observed.  For Jupiter the magnitude of the monopole deflection for a grazing ray is $\\sim 16$ milliarcsecond, to which a component from the quadrupole moment is superimposed with an amplitude of $\\sim 240 $ micro-arcsecond, as showed in \\tref{table-1}. This secondary deflection has a very specific pattern as a function of (i) the position of the star with respect to the oblate deflector and (ii) the orientation of its spin axis.  In \\sref{section-2} we derive the relevant formulas to express this light bending effect using the PPN formalism at large and small angle from the planet. In \\sref{section-3} we describe all the steps needed to simulate the small field experiment when observing stellar background in vicinity of Jupiter. Finally, in \\sref{section-4}, we discuss the data processing of a set of Montecarlo runs and the results in the determination of the PPN Parameter $\\gamma$ and the quadrupole effect from Jupiter. ",
        "conclusions": "\\label{section-5} The objectives of this paper has been to study a new way to determine the PPN parameter $\\gamma$ by exploiting a method of differential positional measurements around Jupiter and to assess the detection of the light bending from the quadrupole moment of Jupiter.  Thanks to the high astrometric accuracy to be achieved by the ESA astrometry mission Gaia and the repeated observations over five years this will be feasible as the new design of Gaia's optical instrument allows to process stellar observations very close to the surface of giant bodies.  We have shown that the deviation to GR with the monopole deflection can be assessed to $10^{-3}$ with Jupiter only, a result in the visible better than the optical accuracy already achieved by Hipparcos with the stars and the solar light bending. With the same observations, the quadrupole light deflection will be detectable for the first time with a 3-$\\sigma$ confidence level.\\\\ Although we have designed an ideal and simplified experiment it includes realistic observations of Jupiter feasible with Gaia, taking into account the astrometric accuracy with the star's magnitude deduced from the current error budget analysis. Whereas the Gaia concept and its design rest entirely on the global astrometry we have given for the first time a realistic figure on the true strength of Gaia to carry out also relativity testing with small field astrometry and propose some strategies to carry out this experiment in the best conditions. When the initial conditions of the Gaia scanning law are known, the same principles will be applied using a real stellar distribution. The method will be also extended to the observations of Saturn, at least for the monopole effect.  \\ack MTC acknowledges financial support from the Henri Poincar{\\'e} Fellowship during her stay at the Observatoire de la C{\\^o}te d'Azur (Cassiop{\\'e}e Department). We wish to thank the members of the \\textit{Relativistic and Reference Frame Working Group} of Gaia for their valuable comments and collaboration.  \\appendix"
    },
    "0512/astro-ph0512390_arXiv.txt": {
        "abstract": "{ In this paper we report a new method for determining the thickness of the non-edge-on disk galaxies. Our method is based on the observation comparing with the theoretical researches of the distribution of the vertical velocity dispersion, which is obtained from the solutions of three dimensional Poisson's equation and the galactic dynamical equation. This method also allows us to investigate the mass-to-light ratio of the disk. As examples, the thickness and mass-to-light ratio of two disk galaxies, NGC 1566 and NGC 5247, which have been extensively studied in spectroscopy, have been calculated and the results are presented. ",
        "introduction": "The thickness of disk galaxies is a very important parameter for understanding these objects, although all of which are very thin comparing with the length of the disk. Sancisi and Allen (1979) estimated the thickness of the edge-on Sb galaxy NGC 891, on the basis of observation of neutral hydrogen. Van der Kruit \\& Searle (1981a) proposed a model for light distribution in the disks of edge-on spiral galaxies, assuming that a galaxy has a locally isothermal, self-gravitating, truncated and exponential disk. The model has the feature of being isothermal in $z$-direction at all radii with a scale parameter of $z_0$ and has an exponential dependence of surface brightness on $r$ with scale length of $r_d$. The space luminosity of this model can be described by \\begin{equation} L(r,z) = L_0 e^{ - r/r_d } \\sec h^2 (z/z_0 ) \\end{equation} With this model, van der Kruit \\& Searle (1981a, 1981b, 1982a, 1982b, named KS hereafter) determined $r_d$ and $z_0$ for several edge-on galaxies without an appreciable bulge. Unfortunately, this method may not be suitable for general disk galaxies, e.g. face-on galaxies.  Peng (1988) put forward a method to measure the thickness of face-on galaxies on the basis of asymptotic formula of the disturbed gravitational potential. And a revised method based on the exact integral expression for disturbed gravitational potential has been presented by Zhao et al. (2004) to estimate the thickness of face-on disk galaxies. But the results obtained by these two correlative methods should be checked with the results gotten by other independent ways.  In this paper, we present the solution of the Jeans equation along $z$-direction first and then describe how to use this solution to determine the thickness of face-on disk galaxies. In Section 3, we show the applications of our method to two near face-on galaxies, NGC 1566 and NGC 5247, which have been extensively studied in spectroscopy. The results are also presented. And a brief discussion is given in the last section of this paper. ",
        "conclusions": "\\begin{table*} \\centering \\caption{Parameters of the near face-on galaxies from our study compare with those of edge-on galaxies from the study by van der Kruit \\& Searle.} \\begin{tabular}{lcccccc} \\hline\\hline NGC &RC3 Type &$M_H$$^a$ &$r_d$$^b$ (kpc) &$H$$^c$ (kpc)& $r_d/H$ &ref \\\\ \\hline 1566 &Sbc & -24.0 & 3.3  &1.1      &3.0  &our study\\\\ 5247 &Sbc & -23.2 & 4.8  &1.5      &3.2  &our study\\\\ 891  &Sb  & -23.3 & 4.9  &0.99     &5.0     &(1)\\\\ 4013 &Sb  & -22.4 & 3.4  &1.1      &3.1     &(2)\\\\ 4217 &Sb  & -22.9 & 3.5  &1.7      &2.1     &(2)\\\\ 4244 &Scd & -20.7 & 2.6  &0.58     &4.5     &(3)\\\\ 4565 &Sb  & -23.7 & 5.5  &0.79     &7.0     &(3)\\\\ 5023 &Scd & -19.5 & 2.0  &0.46     &4.3     &(2)\\\\ 5907 &Sc  & -23.1 & 5.7  &0.83     &6.9     &(2)\\\\ 7814 &Sab & -23.5 & 8.4  &2.0      &4.2     &(4)\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] The magnitudes are all gotten from Jarrett et al. 2003; \\item[$^{\\mathrm{b}}$]$r_d$ in KS is represented by $h$; \\item[$^{\\mathrm{c}}$]$H$ is represented by $z_0$, which is obtained by fitting the light distribution with $\\mu(R,z)=~\\mu(0,0)(R/h) K_1(R/h)$sech$^2(z/z_0)$, $K_1$ is the modified Bessel function, in KS. \\end{list} References { (1) van der Kruit \\& Searle 1981b; (2) van der Kruit \\& Searle 1982a; (3) van der Kruit \\& Searle 1981a; (4)van der Kruit \\& Searle 1982b.} \\end{table*}  \\begin{figure*} \\centering \\epsfig{file=fig5.eps}\\caption{Thickness plotted versus Hubble type(1, 2, 3, 4, 5 and 6 for Sa, Sab, Sb, Sbc, Sc and Scd, respectively.) The continuous line is the fitted result of the median value of each type (represented by the crosses).  } \\label{all} \\end{figure*}  So far, we have proposed an method to determine the thickness of the non-edge on disk galaxies $H$, based on the photometric result, scale length of disk ($r_d$) and extrapolated central surface brightness ($I_0$); and spectroscopic result, the vertical stellar velocity dispersion ($\\left\\langle {V_z^2 } \\right\\rangle$). The fundamental assumption required to derive the conclusions presented here is that the disk mass profile is exponential, both along the vertical direction (i.e., $z$-direction, see equation (3)), which is different from the isothermal model used by KS, and along the radial direction (i.e., $r$-direction, see equation (1)). It has been shown by de Grijs \\& van der Kruit (1996) that the best fitting vertical model is more peaked than expected for an isothermal sheet distribution and the simple exponential fits turned out to be good approximations of the vertical surface brightness profile. Peletier \\& de Grijs (1997) also find that the $K$-band light (often assumed to trace mass) indeed has a very peaky, almost exponential, vertical distribution in edge-on spirals. By contrast, Gould, Bahcall \\& Flynn (1996) find that the vertical distribution of Galactic M dwarfs (which probably trace the stellar disk) is somewhat less peaky.  Because our method is intended to work for a pure disk galaxy, the photometric decomposition is another important ingredient affecting the fitted result. As described in Section 3.1, photometric decomposition allows us to identify the real disk region and the range that the bulge/nucleus (if the galaxy has) affects. We also use the data \\emph{just} outside the effective radius $r_e$ to fit $\\alpha$ for NGC 1566, and we obtain a much bigger $\\chi^2$ of 1.8 and a smaller $\\alpha$ of 0.96 kpc$^{-1}$. For NGC 5247, we use the data out to the center since there are no data just near the effective radius. This also gives a bigger $\\chi^2$ of 0.35 and a smaller $\\alpha$ of 1.0 kpc$^{-1}$. This suggests that we need to select the radial range according to the total mass of the bulge/nucleus. The inclination is another ingredient affecting the fitting results, particularly true for the mass-to-light ratio. Applied corrections for different inclinations will result in different mass-to-light ratio, but almost the same thickness parameter. For the same velocity dispersion data, a smaller inclination gives a larger mass-to-light ratio.  The parameter $H$ should be regarded as about twice the scale height of the best fit of an exponential to the actual distribution. We've compared our some relevant results, e.g. $r_d/H$ (the scale length to the effective height ratio), $H$ (the effective height) of these three galaxies with those of eight edge-on spirals of KS, which are directly measured by fitting the photometric data. As shown in Table 4, our results are very consistent with KS. This make us believing that our new method for determining the thickness of non-edge-on disk galaxies is reasonable and feasible. In Figure 5 we plot the thickness versus the Hubble type. The continuous line is the fitted result for the crosses, which represent the median value of the thickness for each type. The trend that the thickness of disks decreases along the Hubble sequence is also found by Ma (2002) and Zhao, Peng \\& Wang (2004). However, it needs further study on this correlation because the sample used here is rather small.  The method presented in this paper is rather difficult to be used for measurements of large galactic samples because the distributions of the vertical stellar velocity dispersions are inconvenient to be obtained in practice. Fortunately, the most essential purpose of the present paper is to test the availability of the method proposed by Zhao et al. (2004), which is on the basis of the density wave theory. Hence, we measured the thickness of the disks of NGC 1566 and NGC 5247 used the method put forward by Zhao et al (2004) and obtained $0.6\\pm0.4$ kpc and $0.7\\pm0.5$ kpc, for NGC 1566 and NGC 5247, respectively. These results are somewhat smaller than the results obtained by using the method shown in this paper. However, these values are also in the range of the measurement error. Therefore, the method presented in Zhao, Peng \\& Wang (2004), which is more easily to be carried out, might be available and might be used to measure the thickness of the disks of nearly face-on galaxies for large samples."
    },
    "0512/astro-ph0512029_arXiv.txt": {
        "abstract": "After a polemical introduction about the proper activity of an astrophysicist facing a dominant theoretical model and many Tb of highly informative data, I review a few recent results on the properties of galaxies in the nearby (redshift one-tenth) Universe that directly bear on physical cosmology.  In one example, I show that there are a number of ways of measuring, or strongly constraining, massive galaxy--galaxy major merger rates, which are predicted with limited uncertainties in the current generation of models.  In another, I show that we can go beyond ``correlations'' between individual galaxy properties and ``environment''.  Our results show that it is galaxy star-formation histories---not their morphologies---that are sensitive to environmental density.  I look forward to a future, perhaps not that far away, in which these results guide a fundamental modification to our theoretical assumptions, though I fear that the dominant paradigm may not require subversion. ",
        "introduction": "\\subsection{There's a lot of information out there}  Figure~\\ref{fig:m51} shows a (sorry, black-and-white) image of M51 taken from the Sloan Digital Sky Survey \\cite{york00a} imaging data. It positively \\emph{teems} with \\emph{information}.  You can see not just a bulge and spiral arms, but (in the color version) that star forming regions and dusty regions are systematically displaced within the spiral arms, that the spiral galaxy is interacting with a smaller bulge-dominated galaxy, that there are tidal tails and loops at very low surface brightnesses, that the spiral structure is related to the interaction, and even that the bulge-dominated galaxy contains an active nucleus.  One of the ironies of my life is that these fantastically beautiful, detailed, and informative pictures are almost \\emph{completely useless} for constraining physical models for the formation and evolution of galaxies in a cosmological context.  We simply do not have theoretical models---simulations---that produce anything that looks even remotely like Figure~\\ref{fig:m51}.  As you read the contributions to this volume, many of which speak to the great success of our cosmological model and the tremendous advances made in the last ten years in understanding galaxy evolution, I would like you to keep in mind this important point: \\emph{These great theories do not make M51.}  Of course, that is not to say they won't, someday. \\begin{figure}[htb] \\begin{center} \\epsfig{figure=m51.eps,width=\\figurewidth} \\end{center} \\caption{A black-and-white representation of SDSS visible imaging of M51.  Note the tidal tails, relationship of the spiral structure to the interaction, and the regularity of the dust.  This looks much better in color, of course.\\label{fig:m51}} \\end{figure}  If that is one of the ironies of my life, then one of the \\emph{crusades} of my life is to promote an understanding of the tremendous amount of astronomical information available to us, and the best ways to harness that information in the service of physics.  By ``information'' here I don't just mean the touchy--feely thing you get from reading this volume, but I mean the scalar quantity, measured in bits, first worked out for the world by Shannon \\cite{shannon49a} (in what must be one of the most important developments in the entire history of science).  I'll come back to this later, as it bears on questions we will encounter, such as ``should I measure the age--mass relation or the color--magnitude relation?'', and ``should I extract the mean color as a function of environment, or the mean environment as a function of color?''.  Philosophically, The primary goal of observational astrophysics---as distinct from, say, pure astronomy---ought to be to \\emph{rule out} theoretical physical models.  Note that all important physical experiments have the property that they ruled out---or could have ruled out---one of the fundamental, dominant theories of their day. In physical cosmology, we have a fundamental, dominant theory for the growth of structure: Cold Dark Matter (CDM).  Our role as observers is not to \\emph{bolster} this model, or find ad-hoc parameters we can add to the model that \\emph{make it consistent} with the data.  Our role is to find experiments that have the power, even in the face of uncertainties (about, \\eg, how galaxies form), to rule out or substantially modify the fundamental assumptions of this theory.  If an experiment does not have the power to rule out the theory, then it can hardly be said to provide substantial support for it if its results end up being in agreement!  That said, CDM may indeed turn out to be correct.  It is \\emph{certainly} correct on large scales, indeed with its successful prediction of the angular spectrum of CMB inhomogeneities (yes, it was truly a prediction, a parameterized prediction, and the results \\emph{do} live in that parameterized but nonetheless highly constraining prediction space), the amplitude of the galaxy power spectrum, and the baryon feature now detected at low redshift, CDM is one of the best-tested theories in all of the physical sciences.  But \\emph{these tests are all on large scales} ($>10~\\mathrm{Mpc}$).  On scales the size of galaxies and galaxy clusters, CDM is \\emph{not} well tested because on these scales the star-formation, black-hole-accretion, thermodynamic, and dissipative-evolution uncertainties come in.  So lets harness the enormous power (the information) of our observations, and the rapidly improving power of our modeling, to make rock-hard tests of CDM on galaxy scales.  \\subsection{Why work at a redshift of one-tenth?}  The principal disadvantage of working at low redshift is that the redshift \\emph{range} is small.  Although we \\emph{do} see evolution (in, \\eg, the galaxy luminosity function, or mean galaxy specific star-formation rates) \\emph{within} low-redshift samples like SDSS \\cite{york00a} and 2MASS \\cite{skrutskie97a}, we can measure evolution directly with much greater precision by comparing samples at very different redshifts.  On the other hand, the advantages of working at low redshift are many, and they all relate to the enormous amount of information we have about this epoch in the Universe.  With SDSS, 2MASS, and GALEX \\cite{martin05a} we have large enough solid angle and high enough sensitivity to see a large fraction of the galaxy population in a large fraction of the total volume of the Universe out to a redshift of one-tenth.  In this review, I will focus on things we have learned from the SDSS data, simply because that is what I know best.  The SDSS is enormously \\emph{over-designed.}  The spectra in the SDSS have far higher signal-to-noise than is necessary to obtain a redshift; indeed we measure not just redshifts but star-formation rates \\cite{kauffmann03a, quintero04a}, stellar population mixtures \\cite{quintero04a}, and dust attenuations (both those affecting the lines \\cite{tremonti04a} and those affecting the stars \\cite{kauffmann03a}).  But that's not all!  The imaging in which these galaxies were selected for spectroscopy is far deeper than necessary for object selection; a typical Main Sample \\cite{strauss02a} galaxy target is detected in the SDSS imaging at a $S/N$ of many hundreds.  This means that we have very high-precision galaxy colors and magnitudes, sizes, concentrations, and surface brightnesses \\cite{hogg02a, blanton03d}.  Finally, and importantly for what follows, the large contiguous area of the SDSS (and other surveys) permits the study of clustering and galaxy environments; it also allows us to identify large, gravitationally bound systems, such as groups, clusters, and superclusters.  Survey edges do not technically \\emph{prevent} the measurement of clustering and environments, but in practice, they make it difficult.  So with the huge ``volume-to-boundary'' ratio of the SDSS, we have good-quality measures of galaxy environments, the mean galaxy density, and the correlation function on all scales out to the scale of large-scale homogeneity \\cite{hogg05a, eisenstein05b} (these studies also demonstrate that the photometric calibration of the survey is very stable on large angular scales). ",
        "conclusions": ""
    },
    "0512/astro-ph0512503_arXiv.txt": {
        "abstract": "We construct orbit-based axisymmetric dynamical models for the globular cluster M15 which fit groundbased line-of-sight velocities and Hubble Space Telescope line-of-sight velocities and proper motions. This allows us to constrain the variation of the mass-to-light ratio $M/L$ as a function of radius in the cluster, and to measure the distance and inclination of the cluster. We obtain a best-fitting inclination of $60^\\circ\\pm15^\\circ$, a dynamical distance of $10.3\\pm0.4$ kpc and an $M/L$ profile with a central peak. The inferred mass in the central 0.05 parsec is 3400 $\\Msun$, implying a central density of at least $7.4\\times10^6\\Msun$ pc$^{-3}$. We cannot distinguish the nature of the central mass concentration. It could be an IMBH or it could be large number of compact objects, or it could be a combination. The central 4 arcsec of M15 appears to contain a rapidly spinning core, and we speculate on its origin. ",
        "introduction": "\\label{sec:intro}  M15 is a well-studied globular cluster. It has a very steep central luminosity profile, and may be in the post-core-collapse stage (e.g., Phinney 1993; Trager, King \\& Djorgovski 1995).  Measurements of nearly two thousand line-of-sight velocities (from the ground and with HST) and proper motions (with HST) have recently become available (Gebhardt et al.\\ 2000, hereafter G00; McNamara, Harrison \\& Anderson 2003, hereafter M03).  McNamara, Harrison \\& Baumgardt (2004, hereafter M04) restricted themselves to the subset of 237 stars inside $0\\farcm3$ of the center of M15 for which both Fabry--Perot radial velocities and HST proper motions were measured, and computed the mean dispersions in these measurements. Assuming the cluster is an isotropic sphere, and the observed stars are representative, the ratio of these dispersions (one in \\kms, the other in \\masyr) provides the distance (Cudworth 1979, Binney \\& Tremaine 1987). M04 find a distance of $9.98\\pm0.47$ kpc, which is consistent with the canonical value of 10.4 kpc (Durrell \\& Harris 1993), but is smaller than, e.g., the recent determination of 11.2 kpc by Kraft \\& Ivans (2003) who used a globular cluster metallicity scale, based upon Fe~II lines.  Here we extend the M04 study by using a larger fraction of the line-of-sight velocity and proper motion samples, and comparing these with more general dynamical models to study the internal structure of the cluster as a function of radius. We follow the approach taken by van de Ven et al.\\ (2005, herafter V05), who constructed axisymmetric dynamical models for the globular cluster $\\omega$ Centauri and fitted these to groundbased proper motions and line-of-sight velocities. This technique provides the internal dynamical structure as well as the inclination of the cluster, an unbiased and accurate dynamical distance, and the $M/L$-profile. Our aim is to derive similar information for M15. We are particularly interested in the $M/L$ profile, as significant mass segregation is believed to have occurred in the cluster (Dull et al.\\ 1997). The HST proper motions have sufficient spatial resolution to study the dynamical structure and mass concentration in the center.  In Section~\\ref{sec:data}, we summarize the observational data. In Section~\\ref{sec:second-moments}, we consider the influence of measurement errors on the data, select the stars to be used for the dynamical modeling and also study the possible residual systematic effects in the observed mean motions. We construct dynamical models in Section~\\ref{sec:models}, derive a distance, and investigate the effect of the unknown inclination and of radial $M/L$ variations in the cluster. We discuss the dynamics of the central 0.2 parsec of M15 in Section~\\ref{sec:spinning-core}, and summarize our conclusions in Section~\\ref{sec:discussion}.  \\begin{figure*} \\epsscale{1.0} \\plotone{f1.eps} \\figcaption{\\figdisp} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  We studied the globular cluster M15 by fitting line-of-sight velocities, HST proper motions and surface brightness profiles with orbit-based axisymmetric dynamical models.  The observations used for the modeling consisted of a luminosity profile from Noyola \\& Gebhardt (2005), 1264 line-of-sight velocity measurements from G00 and a sample of 703 HST proper motions from M03. The line-of-sight data extends out to $7'$, while the proper motions cover the inner $0\\farcm25$. The models provide a good fit to the observations and allow us to measure the distance and inclination of the cluster, the orbital structure, and the mass-to-light ratio $M/L$ as a function of radius. We obtain a best-fit value for the inclination of $i=60^\\circ\\pm15$ and a dynamical distance of $D=10.3\\pm0.4$ kpc, in good agreement with the canonical value.  Our best-fit model has a $500^{+2500}_{-500}$ $\\Msun$ dark central mass and the $M/L$ profile shown in Figure~\\ref{figmlprof}, which has a central peak and a minimum at $0\\farcm1$. The overall shape of the profile resembles the shape expected for an expanded core globular cluster (Dull et al.\\ 2003). The central $M/L$ peak and the dark central mass together represent a mass of 3400 $M_\\odot$ inside the inner $1\\farcs0$ (0.05 parsec at a distance of 10.3 kpc).  This suggests that the center harbors a large amount of dark mass. We cannot distinguish the nature of the central mass concentration. It could be an IMBH or it could be large number of compact objects, or it could be a combination.  We found that a heavily smoothed image of M15 shows a flattened structure inside $4''$, with a different PA from the outer part of the cluster. The line-of-sight and proper motion data inside this radius can be fitted using relation (1). The fit gives a PA similar to that of the flattened structure and a rotational velocity of 10 \\kms. This suggests the structure is real, and constitutes a fast-spinning decoupled core at the center of the cluster.  A significant improvement in the accuracy of the dynamical models for M15 is possible by increasing the accuracy of the proper motions and radial velocities.  This appears possible with the ACS onboard HST, and with high-resolution spectrographs on 8m class telescopes.\\looseness=1"
    },
    "0512/astro-ph0512288.txt": {
        "abstract": "We determine the characteristics of the 7\\,mm to 20\\,cm wavelength radio variability in Sgr A* on time scales from days to three decades.  The amplitude of the intensity modulation is between 30 and 39\\% at all wavelengths.  Analysis of uniformly sampled data with proper accounting of the sampling errors associated with the lightcurves shows that Sgr A* exhibits no 57- or 106-day quasi-periodic oscillations, contrary to previous claims.  The cause of the variability is investigated by examining a number of plausible scintillation models, enabling those variations which could be attributed to interstellar scintillation to be isolated from those that must be intrinsic to the source.  Thin-screen scattering models do not account for the variability amplitude on most time scales.  However, models in which the scattering region is extended out to a radius of 50-500\\,pc from the Galactic Center account well for the broad characteristics of the variability on $>4$-day time scales.  The $\\sim 10$\\% variability on $<4$-day time scales at $0.7-3\\,$cm appears to be intrinsic to the source. % Centimeter-wavelength variations on time scales less than $5\\,$days appear to be intrinsic to the source. The degree of scintillation variability expected at millimeter wavelengths depends sensitively on the intrinsic source size; the variations, if due to scintillation, would require an intrinsic source size smaller than that expected. ",
        "introduction": "% INTRODUCTION %-------------------------------------------------------------------------------------------- The compact radio source associated with the black hole at the Galactic Center, Sgr A*, is known to vary at millimeter and centimeter wavelengths on time scales from hours to years (Brown \\& Lo 1982; %Genzel \\et\\ ???? ; Zhao \\et\\ 1992; Bower \\et\\ 2002; Herrnstein \\et\\ 2004).  The origin of these variations remains unclear, with strong arguments for both extrinsic and intrinsic mechanisms having been advanced (e.g. Zhao \\et\\ 1989 compared with Zhao \\et\\ 2001).  Interstellar scintillation is the primary mechanism which may cause any extrinsic variability. The same plasma that is responsible for the scatter-broadening of Sgr A* at millimeter and centimeter wavelengths (e.g. Lo \\et\\ 1998, Bower \\et\\ 2004) is also expected to cause the source to exhibit refractive intensity variations.  It has been argued that much of the monthly to yearly variability in Sgr A* at wavelengths longer than 6\\,cm can be explained in terms of refractive interstellar scintillation provided that scattering material moves across our line of sight at a relatively high speed of $\\sim 1000-2000\\,$km\\,s$^{-1}$ (Zhao \\et\\ 1989).   However, the variability amplitude is not so easily accounted for: a scattering medium modeled as a single thin screen underpredicts the observed variability amplitude ({\\it ibid.}) while extended medium models, which are in principle capable of explaining higher refractive modulation amplitudes for the same degree of scatter broadening, have not been investigated in the context of the Galactic Center.  Recent interpretations favor an intrinsic origin for much of the centimeter wavelength variability.  These center around claims of 106-day quasi-periodic variations at wavelengths shorter than $\\sim 3\\,$cm (Zhao \\et\\ 2001) and of 57-day quasi-periodic behavior at 2.3\\,GHz (Falcke 1999).  The oscillations possess only a modest spectral purity, with the highest purity $\\nu/\\Delta \\nu=2.2 \\pm 0.3$ reported at 1.3\\,cm.  Zhao \\et\\ (2001) discuss the origin of these oscillations in terms of periodic flares from a jet nozzle or an instability in the accretion disk triggering, for example, quasi-periodic production of convection bubbles.  It is widely supposed that the oscillations must reflect a process intrinsic to Sgr A* itself because scintillation is incapable of producing such regular oscillations.  Yet observations of certain intra-day variable quasars, whose variations are proven to be scintillation-induced,  invalidate this argument because their fluctuations often exhibit even higher degrees of spectral purity (Kedziora-Chudczer \\et\\ 1997; Rickett, Kedziora-Chudczer \\& Jauncey 2002).  Nonetheless there is little dispute that at least some of the variability is intrinsic. Detections of flares at millimeter, IR and X-ray wavelengths (Wright \\& Backer 1993; Tsuboi, Miyazaki \\& Tsutsumi 1999, Eckart \\et\\ 2004; Baganoff \\et\\ 2001) conclusively demonstrate that the source is intrinsically variable.  A possible connection between X-ray flaring and unusually large flux density excursions at 7\\,mm is also reported (Zhao \\et\\ 2004).  However, it is difficult to ascertain how much variability observed at centimeter wavelengths could be attributed to flaring since neither the duty cycle nor the energy distribution of mm or X-ray flares is well-constrained, much less the physical connection between centimeter and mm or X-ray behavior.  %The literature is replete with pithy remarks about whether scintillation is a viable explanation for any of the variations observed in Sgr\\,A*.  Despite the many recent observational results concerning the properties of Sgr A*'s variability, a dearth of corresponding theoretical efforts has failed to place these results in context, leaving us none the wiser as to their cause.  For instance, while it is acknowledged that scintillation variability is likely to be important at centimeter wavelengths, no variations have been specifically attributed to it, and no realistic modeling has been applied to investigate what contribution it could conceivably make.  This paper aims to redress the balance by investigating two outstanding issues: (i) what exactly does a model of Sgr A*'s variability need to explain and (ii) can one deduce which variations {\\it must} be intrinsic to the source by eliminating the variations that can be explained by interstellar scintillation?  The next section of this paper is devoted to the former question, including a critical examination of the $\\sim 100\\,$day quasi-periodic oscillations reported in Sgr A* (Zhao \\et\\ 2001), while \\S\\ref{Scint} addresses the latter question.  We compare the models to the observations in \\S\\ref{Comparison}, and summarize our findings and briefly detail their implications in \\S\\ref{Conclusions}.  %-------------------------------------------------------------------------------------------- ",
        "conclusions": "\\label{Conclusions} % CONCLUSIONS %--------------------------------------------------------------------------------------------  Our analysis of multi-frequency monitoring data presented by Zhao \\et\\ (1992, 2001), Herrnstein \\et\\ (2004) and Falcke (1999) indicates that Sgr A* exhibits no quasi-periodic oscillatory behavior on any time scale between one week and 200 days.   The variability amplitudes are remarkably constant with frequency, varying between 30 and 39\\%, but the time scale on which they saturate increases with wavelength.  Several structure functions show evidence for variability on multiple time scales.  If the errors associated with the flux density measurements are correct, the data of Herrnstein \\et\\ (2004) indicate that the source exhibits unresolved 6-10\\% inter-day ($<4$-day) variations between 7\\,mm and 2\\,cm. No structure functions exhibit, within the errors, any evidence for appreciable variability with time scales longer than 1000\\,days.  The long term stability of the radio flux implies there is very little long-term variation in the accretion rate. % Short-term fluctuations can be driven by excitation of high energy electrons rather than accretion rate changes.  We sought to explain the general features of the variability by reproducing the shape and amplitudes of the observed structure functions using several scintillation models.  Both thin-screen and extended-medium models were considered.  No thin screen model accounts for the properties of the variations.  The structure functions of the observed lightcurves rise less steeply with time lag and saturate at higher amplitudes than predicted.  They underpredict the amplitude of the variability by at least a factor of two at the saturation time scale and often by a more than an order of magnitude on shorter time scales.  If the medium responsible for the scattering of Sgr A* lies on a thin screen all of the observed flux variability must be intrinsic to the source itself.  Certain extended-medium models, on the other hand, do explain the amplitude of the fluctuations over a large range of time scales.  Models in which the electron density fluctuations follow a Kolmogorov power spectrum, corresponding to $\\beta=11/3$, and a slightly steeper, $\\beta=3.9$, power spectrum were investigated.  Of the two models, only those with a $\\beta=3.9$ spectrum account for the amplitude of the fluctuations at all wavelengths.  If scintillation is to precisely predict the amplitude of the flux density variability of Sgr A*, this suggests that the power spectrum of the Galactic Center turbulence is slightly less steep, with an index lying in the range $3.8 \\la \\beta < 3.9$.   The most successful extended-medium model examined was used to predict the maximum contribution that scintillation could make to future observations of Sgr A* at millimeter wavelengths.  The expected variability amplitude depends strongly on the intrinsic source size.  A $1-10\\,\\mu$as object at 3\\,mm would undergo fractional root-mean-square fluctuations of $\\sim 25$\\%, but a $100\\,\\mu$as source would exhibit only 3\\% variations. Scintillation is even more sensitive to source size at 1\\,mm, with a $1\\,\\mu$as source expected to display 26\\% variations, but a $10\\,\\mu$as source would display only 3\\% variability.   % [From Geoff: We should discuss or at least present physical conditions appropriate for our favored scintillating medium: density, or integrated column density.  Is there a known medium that we can associate this with?  The HIM?  What does it imply that we favor $\\beta=3.9$?  With what physical structures can we associate the extended scattering medium?  Given only the approximate match between our model structure functions and the actual structure functions, we do not expect that the extended scattering medium must in fact be parameterized as discussed in the text. In fact, the extended scattering medium might consist of only a few thin scattering media distributed over a range of distances from Sgr A*. These thin media would have characteristics similar to the thin medium discussed by Lazio \\& Cordes (1998), with densities $>10^2 {\\rm\\ cm^{-3}}$. The details of the structure functions are not sufficient to constrain this result further.  For comparison purposes, we do compute the mean density associated with the extended scattering medium as presented. In this case, the mean density is  $\\sim 10\\,$cm$\\,^{-3}$, assuming the outer scale of the turbulence to be 1\\,pc.   This density is substantially greater than the density of the diffuse hot ionized gas ($T_e \\sim 10^7\\,$K, $n_e \\sim 0.05 {\\rm\\ cm^{-3}}$) detected in the Galactic Center.  As Lazio \\& Cordes (1998) discuss, the scattering medium may be the interface between this hot medium and molecular clouds in the central 100 pc.  Our results are consistent with the picture reached by Lazio \\& Cordes (1998) for the scattering medium with the modification that the scattering may take place at a range of distances from Sgr A*."
    },
    "0512/astro-ph0512244_arXiv.txt": {
        "abstract": "The vertical profiles of disc galaxies are built by the material trapped around stable periodic orbits, which form their ``skeletons''. According to this, the knowledge of the stability of the main families of periodic orbits in appropriate 3D models, can predict possible morphologies for edge-on disc galaxies. In a pilot survey we compare the orbital structures which lead to the appearance of ``peanuts'' and ``X''-like features with the edge-on profiles of three disc galaxies (IC~2531, NGC~4013 and UGC~2048).  The subtraction from the images of a model representing the axisymmetric component of the galaxies reveals the contribution of the non-axisymmetric terms. We find a direct correspondence between the orbital profiles of 3D bars in models and the observed main morphological features of the residuals. We also apply a simple unsharp masking technique in order to study the sharpest features of the images. Our basic conclusion is that the morphology of the boxy ``bulges'' of these galaxies can be explained by considering disc material trapped around stable 3D periodic orbits. In most models these building-blocks periodic orbits are bifurcated from the planar central family of a non-axisymmetric component, usually a bar, at {\\em low order} vertical resonances. In such a case the boxy ``bulges'' are {\\em parts} of bars seen edge-on. For the three galaxies we study the families associated with the ``peanut'' or ``X''-shape morphology are most probably bifurcations at the vertical 2/1 or 4/1 resonance. ",
        "introduction": "In Patsis \\& Grosb{\\o}l (1996), Patsis, Athanassoula, Grosb{\\o}l et al. (2002a), and Patsis, Skokos \\& Athanassoula (2002b, hereafter PSA), have been presented boxy orbital profiles of 3D, time-independent, models of disc galaxies. They have been built by combinations of stable, 3D, orbits belonging in most cases to families of periodic orbits bifurcated from the planar x1 family (see e.g.  Contopoulos 2002, or Contopoulos \\& Grosb{\\o}l 1989) at the vertical $n/1$ resonances.  The most efficient dynamical mechanism for introducing vertical resonances in a system is the presence of non-axisymmetric components, especially bars.  Vertical profiles of 3D bars have been constructed by PSA, based on the orbital analysis of 3D Ferrers bars (Pfenniger 1984, 1985; Skokos, Patsis \\& Athanassoula 2002a,b). A profile of a single family includes stable orbits in a range of ``energies'' (Jacobi constants - $E_j$). In order to construct a family profile one needs to know its stable orbits. Out of this library of stable periodic orbits any subgroup that helps in matching the observed morphological structure can be chosen. Every family has its own characteristic orbital profile, and the profiles of the same family in several models are similar. Combinations of the profiles of several families of a model comprise a model profile. Particular useful representations of the orbital profiles that allow comparisons with real galaxies, or with snapshots of $N$-body simulations, are their {\\em weighted images} (PSA). The advantage of these images is that they show the relative importance of every orbit in a profile. They also show the details of the morphology in locations where more than one orbit contribute. The images are composed by periodic orbits weighted by the mean density of the model at the points visited by the orbit. Along each orbit one picks points at equal time steps. The density of the model is calculated at each of these points and the mean density is taken as the weight of the orbit. Having constructed images for every orbit (normalized over its total intensity), we can combine them to built a profile for a family of orbits. For the profile of a family are used orbits equally spaced in their mean radius.  The basic conclusions of the orbital analysis relevant to the present study are: \\begin{enumerate} \\item The morphology supported by the composite orbital profile of a family may differ from the morphology of the individual periodic orbits calculated at a particular $E_j$. Thus, in order to get the backbone of a structure to be compared with the morphology of a galaxy, one needs to know the evolution of the shape of the stable orbits of a family as a function of $E_j$.  \\item The radial extent of the profile of an orbital 3D family is usually confined within a radius corresponding to a certain $E_j$ value. Orbits with a Jacobi constant larger than this $E_j$, increase their size by increasing practically only in the vertical dimension, and this leads to models with stair-type edge-on profiles (for more details see Patsis et al. 2002a). \\end{enumerate}  To these two, one should add that it is the presence of the vertical resonances and not the detailed kind of perturbation in the models that shapes the boxiness of the orbital profiles. However, ``X''-like features as these we discuss here are typical of strong bar components.  Recently, unsharp masking techniques applied to images of galaxies (Aronica, Athanassoula, Bureau et al. 2003; Aronica, Bureau, Athanassoula et al. 2004; Bureau, Athanassoula, Chung et al. 2004) as well as to images of snapshots of $N$-body simulations (Athanassoula 2005a,b) have shown excellent agreement between the image morphologies and what is predicted by the orbital theory in PSA.  In particular, Aronica et al. (2003), compared the image of ESO~597-036 after unsharp masking with a model in PSA. They found besides a conspicuous ``X''-shape feature in the central part, surface brightness enhancements along the equatorial plane of the galaxy. Both features have their counterparts in the orbital models and can be explained by material trapped around stable periodic orbits. Similar features are indicated by Aronica et al. (2004) also for ESO-443G042 and by Bureau et al. (2004) for the case of NGC~128. The comparison of features between orbital and $N$-body models on the other hand show agreement in even finer details. In both papers by Athanassoula (2005a,b) one can find in the edge-on views of the models besides ``X''-like features, density enhancements on the equatorial plane and features like ``parentheses'' (see e.g. Fig.~6 in Athanassoula 2005a). This is a strong indication that a large percentage of the material in the $N$-body simulation follows orbits around the families bifurcated at the vertical $n/1$ resonances with small $n$.  In the present paper we apply two image processing techniques on the images of three edge-on disc galaxies with rather boxy ``bulges'' in order to detect structures similar to those predicted by the orbital theory. We focus our attention to the structures observed in the central regions of IC~2531, NGC~4013 and UGC~2048 (NGC~973). First we subtract from the I-band images of the galaxies an axisymmetric model that was found by Xilouris, Kylafis, Papamastorakis et al. (1997) to describe best the smooth distribution of stars and dust in these objects. The models are used to isolate the non-axisymmetric term in the profiles of the galaxies. This term is expected to reflect in a straightforward way the presence of vertical resonances (PSA). A second technique we apply is just a gaussian filtering (unsharp masking) on the images. In Section \\ref{observations} we give information about the observations of the galaxies, in Section \\ref{model} we describe the different image processing techniques, in Section \\ref{results} we give the results of the image processing we performed and finally we discuss our results and present our conclusions in Section \\ref{conclusions}. ",
        "conclusions": "Both image processing techniques revealed a kind of ``X'' structure in the central components of the three edge-on galaxies with boxy profiles we studied.  In Fig.~1, where we give the residuals after subtracting a Xilouris et al. (1997) model we see the overall shape of the non-axisymmetric component. The unsharp masking in Fig.~2 shows more clearly the high intensity features of the ``bulges''. They emerge out of the equatorial plane as distinct branches.  In most cases of the models in PSA an ``X'' is supported by material trapped around x1v1 orbits (Skokos et al. 2002a), i.e. by stable 3D orbits introduced in the system at the vertical 2/1 resonance. However, there are other 3D families as well, which could support straight line segments emerging out of the equatorial plane in the edge-on projections of the orbital models. Such is the case of x1v5, according to the nomenclature of Skokos et al. (2002a), which is a family born at the vertical 4/1 resonance. This family has been associated with boxy central components already by Pfenniger (1985). A typical orbital profile is given in Fig.~7a in PSA.  In models (orbital or $N$-body), a visual difference can be found if we compare the side-on views. The distance between the wings of the ``X'' feature in the x1v5 case is larger than the distance of the corresponding wings in x1v1 profiles. The difference can be seen in Fig.~3, where we give in (a) a side-on typical x1v1 profile and in (b) a profile dominated by the presence of the x1v5 family. They correspond to models D and B in PSA.  As we mentioned in the introduction, there is a certain value of $E_J$, beyond which the orbits of a family grow practically only in the $z$ direction. So, for each family, there is a maximum radius on the equatorial plane, within which the projections of its orbits are confined.  In the side-on views we can measure the length of the projection along the major axis of the bar.  In the two specific examples we give here, the projection of the x1v1 orbits on the major axis of the bar in Fig.~3a corresponds roughly to a length 45\\% of the longest bar supporting orbits, i.e. the length of the bar in our orbital model.  On the other hand in Fig.~3b, this percentage reaches almost 90\\%. \\begin{figure} \\centerline{\\includegraphics[scale=0.58]{fig3.eps}} \\caption{(a) A typical x1v1 side-on profile based on orbits of model D in PSA. Arrows indicate the inward bending of the ``X'' wings at large distances above the equatorial plane. (b) A typical side-on x1v5 profile corresponding to model B in PSA. It extends to larger radii than the profile in (a).} \\end{figure} Unfortunately this criterion cannot in general apply on the images of real galaxies since we are missing the essential information about the orientation of the boxy structure with respect to the line-of-sight. Thus, at the level of the current analysis, we cannot point to a single vertical resonance and attribute to it the observed morphological feature. A galaxy may lack a vertical 2/1 (or even 3/1) resonance, in which case material may be trapped by stable orbits bifurcated at the vertical 4/1 resonance. The corresponding profile in such a case will look like what we see in Fig.~3b. It consists of stable orbits of the following families: (a) x1v5, (b) a bifurcation of it, and (c) z3.1s (PSA).  Another point that has to be mentioned, is the inwards bending or break of the wings of ``X'' close to their maximum height above the equatorial plane. According to the orbital models this can be explained by the fact that the orbits of a family increase their size practically only in the vertical direction beyond a certain $E_j$ value. In Fig.~3a we point with arrows to these breaks at the x1v1 profile. However, this is expected to happen to the profiles of other families as well if they are populated with orbits with high energies ($E_j$) (see a characteristic case in PSA, Fig.18). This tendency can be observed at the wings of the ``X'' features in Fig.~2a (NGC~4013) and Fig.~2c (UGC~2048). We indicate these breaks with arrows.  The main conclusion of the present study is that the careful subtraction of axisymmetric components from the profiles of the three edge-on disk galaxies (NGC~4013, IC~2531 and UGC~2048) reveals an ``X''-like morphology, which indicates the presence of a bar. With this method we could isolate the light coming from the non-axisymmetric components. The similarity of these components with the morphology of the orbital profiles is a strong indication that the boxy ``bulges'' of these three galaxies are parts of bar structures observed edge-on.  If the ``X''-like structure of the residuals is due to stars trapped around stable periodic orbits, then we expect the presence of sharp features (the wings of the ``X'') {\\em inside} the bar region. By applying the unsharp masking on the images we find, in the foreseen by the orbital models regions, sharp features. The two methods are complementary. The unsharp masking does not give any information about the overall shape of the non-axisymmetric component. It shows however that in the regions we see the residuals we have the expected sharp features. These features show a nice correspondence with the dense parts of published orbital models (see PSA, Fig.1a, 3b, 7a, 9a). These techniques, and especially the study of the residuals after subtracting axisymmetric models, should be applied in a large sample of galaxies in order to estimate the frequency of these features in the profiles of boxy edge-on galaxies."
    },
    "0512/astro-ph0512072_arXiv.txt": {
        "abstract": "We present spectroscopic and photometric observations of the eclipsing system \\V10, previously thought to be a member of the rare class of ``cool Algols\". We show that it is instead a hierarchical triple system in which the inner eclipsing pair (with $P = 2.35$ days) is composed of main-sequence stars and is well detached, and the third star is also visible in the spectrum. We combine the radial velocities for the three stars, times of eclipse, and intermediate astrometric data from the HIPPARCOS mission (abscissae residuals) to establish the elements of the outer orbit, which is eccentric and has a period of 15.8 yr. We determine accurate values for the masses, radii, and effective temperatures of the binary components: $M_{\\rm Aa} = 1.282 \\pm 0.015$~M$_{\\sun}$, $R_{\\rm Aa} = 1.615 \\pm 0.017$~R$_{\\sun}$, and $T_{\\rm eff}^{\\rm Aa} = 6180 \\pm 100$~K for the primary (star Aa), and $M_{\\rm Ab} = 0.9315 \\pm 0.0068$~M$_{\\sun}$, $R_{\\rm Ab} = 0.974 \\pm 0.020$~R$_{\\sun}$, and $T_{\\rm eff}^{\\rm Ab} = 5300 \\pm 150$~K for the secondary (Ab). The masses and radii have relative errors of only 1--2\\%.  Both stars are rotating rapidly ($v \\sin i$ values are $36 \\pm 2$~\\kms\\ and $20 \\pm 3$~\\kms) and have their rotation synchronized with the orbital motion.  There are signs of activity including strong X-ray emission and possibly spots.  The mass of the tertiary is determined to be $M_{\\rm B} = 0.925 \\pm 0.036$~M$_{\\sun}$ and its effective temperature is $T_{\\rm eff}^{\\rm B} = 5670 \\pm 150$~K. The system is at a distance of $166.9 \\pm 5.6$ pc.  Current stellar evolution models that use a mixing length parameter $\\alpha_{\\rm ML}$ appropriate for the Sun agree well with the properties of the primary, but show a very large discrepancy in the radius of the secondary, in the sense that the predicted values are $\\sim$10\\% smaller than observed (a $\\sim$5$\\sigma$ effect). In addition, the temperature is cooler than predicted by some 200~K. These discrepancies are quite remarkable given that the star is only 7\\% less massive than the Sun, the calibration point of all stellar models. Similar differences have been seen before for later-type stars, but the source of the problem has remained unclear. A comparison with the properties of other stars of similar mass as the secondary in \\V10 has allowed us to identify the chromospheric activity as the likely cause of the effect. Inactive stars agree very well with the models, while active ones such as \\V10 Ab appear systematically too large and too cool. Theory provides an understanding of this in terms of the strong magnetic fields commonly associated with stellar activity, which tend to inhibit convective heat transport.  The reduced convection explains why fits to models with a smaller mixing length parameter of $\\alpha_{\\rm ML} = 1.0$ seem to give better agreement with the observations for \\V10~Ab. ",
        "introduction": "\\label{sec:introduction}  Accurately determined properties of stars in detached eclipsing binaries provide fundamental data for testing models of stellar structure and stellar evolution \\citep[see, e.g.,][]{Andersen:91, Andersen:97}. For stars less massive than the Sun properties such as the stellar radius and the effective temperature have occasionally been found to disagree with model predictions \\citep[see, e.g.,][]{Lacy:77, Popper:97, Clausen:99a, Torres:02, Ribas:03}. Directed efforts to find additional systems in this regime suitable for testing theory \\citep{Popper:96, Clausen:99b} have produced a few cases, while other examples have been found serendipitously \\citep[e.g.,][]{Creevey:05, Lopez-Morales:05}. The present binary system is in the second category, since it was originally thought to be of a completely different nature.  The photometric variability of V1061 Cygni (also known as HD~235444, HIP~104263, RX\\,J2107.3+5202, $\\alpha = 21^{\\rm h} 07^{\\rm m} 20\\fs52$, $\\delta = +52\\arcdeg 02\\arcmin 58\\farcs4$, J2000, SpT F9, $V = 9.24$) was discovered photographically by \\cite{Strohmeier:59}, and the object was classified by \\cite{Strohmeier:62} as an Algol-type binary with a period of 2.346656 days.  Other than occasional measurements of the time of primary eclipse, the system received very little attention until the spectroscopic work by \\cite{Popper:96}, who observed it as part of his program to search for eclipsing binaries containing at least one lower main-sequence star (late F to K).  On the basis of two high-resolution spectra and other information Popper concluded that \\V10 was most likely a semi-detached system of the rare ``cool Algol\" class, and dropped it from his program. Unlike the classical Algols, which are composed of a cool giant or subgiant and an early-type star, the mass-gainer in the cool Algols is also of late spectral type \\citep[see][]{Popper:92}.  Since less than a dozen of these systems are known, \\V10 was placed on the observing list at the Harvard-Smithsonian Center for Astrophysics (CfA) in 1998 for spectroscopic monitoring, and photometric observations began later. Not only did we discover that it is not a cool Algol \\citep[it is well detached, as reported by][]{Sheets:03}, but we also found that: \\emph{i)} it is triple-lined (and a hierarchical triple); \\emph{ii)} the secondary in the eclipsing pair is less massive than the Sun and therefore potentially interesting for constraining models of stellar structure and evolution (Popper's original motivation for observing it); and \\emph{iii)} the mass ratio of the binary is quite different from unity, which makes it a favorable case for such tests. Furthermore, the comparison with theory shows a significant discrepancy in the radius of the secondary, corroborating similar evidence from other systems and providing some insight into the problem.  We describe below our observations and complete analysis of this system, including a discussion of the possible nature of the deviations from the models for low-mass stars. ",
        "conclusions": "The results of our spectroscopic, photometric, and astrometric analyses of the \\V10 system have taken us in a rather different direction than we anticipated when we began this study. The possible status of the object as an example of the rare class of cool Algols is now clearly ruled out.  Instead, we have shown here that it is a hierarchical triple system with an outer period of 15.8~yr, in which the eclipsing inner pair is well detached, is composed of main-sequence stars, has a mass ratio quite different from unity, and has a secondary that is slightly below a solar mass. The absolute masses and radii for the binary components are determined with a relative precision of 2\\% or better, and the mass of the third star is good to 4\\%.  While the primary star is well fit in mass, radius, temperature, and luminosity by standard stellar evolution models with a metallicity near solar and a mixing length parameter set by the calibration to the Sun, the secondary appears $\\sim$10\\% too large. This discrepancy is 5 times the size of the observational errors, and quite surprising for a star that differs by only 7\\% in mass from the Sun. There are also indications that it is cooler than predicted by some 200 K. \\V10 is yet another example highlighting our incomplete understanding of the structure and evolution of stars in the lower main sequence.  Similar differences in size and temperature have been noticed previously for lower-mass stars with accurately measured properties, but the source of the problem has remained unclear. By comparing \\V10 Ab to several other objects of nearly the same mass we have identified the activity level as a key factor distinguishing cases that show the radius discrepancy from those that do not. This link between activity and increased radius has been mentioned in the literature before, but is shown here for the first time for stars with accurately known masses, radii, and temperatures.  It is often stated that the structure and evolution of stars (particularly those close to a solar mass) are completely determined once the chemical composition and mass are specified. It is quite clear now that for stars of order 1~M$_{\\sun}$ or less an additional parameter must be taken into account, which has to do with the level of chromospheric activity.  Whether this parameter is directly the rotational speed (e.g., $v \\sin i$) or period, the magnetic field strength, the Rossby number \\citep[e.g.,][]{Noyes:84, Basri:87}, or some other more complicated activity indicator remains to be determined.  To first order it appears that the effective mixing length may be a useful proxy, but with one exception current stellar evolution models are not publicly available for more than one value of $\\alpha_{\\rm ML}$, so testing this is somewhat difficult in practice. If the predictions from the models are to reach an accuracy matching current observations of low-mass stars ($\\sim$1--2\\% relative errors in the masses and radii), this effect can no longer be ignored. Further progress will require more examples of binary systems with well determined physical parameters and different levels of activity in the relevant mass regime in order to help calibrate any such parameter.  Testing models of single-star evolution by means of eclipsing binaries, as astronomers have done for decades, might perhaps be seen as part of the problem since the enhanced activity displayed by many of these systems is a direct result of tidal synchronization that occurs only in close binaries.  Although wider eclipsing pairs with inactive components under 1~M$_{\\sun}$ certainly do exist, they are less common and in some respects more difficult to study. Stars near the bottom of the main sequence, M dwarfs in particular, are found to be active more often than not.  The effects of magnetic fields on the evolution of stars have already begun to be explored by theorists \\citep[e.g.,][although the latter authors focus on more massive stars]{D'Antona:00, Mullan:01, Maeder:03}, and initial comparisons with observations are encouraging. One area where this is likely to have a significant impact is the study of T~Tauri stars, which are typically very active. Even in substellar objects such as brown dwarfs activity appears to be quite common, and might be expected to have similar consequences on their structure. Some evidence for this has already been reported \\citep{Mohanty:04}."
    },
    "0512/astro-ph0512591_arXiv.txt": {
        "abstract": "{ Recent observations with the XMM-Newton satellite confirmed the existence of the {\\it soft excess} phenomenon in galaxy clusters, earlier discovered in several EUVE, ROSAT and BeppoSAX observations. Among the clusters for which XMM has reported detection of soft excess emission are MKW03s and A2052, two clusters in the Hercules concentration.  The Hercules supercluster lies along the southern extension of the North Polar Spur, a region of bright soft X-ray emission clearly visible in ROSAT All-Sky Survey images. We analyze 11 pointed ROSAT PSPC observations toward 3 clusters in the Hercule concentration, MKW03s, A2052 and A2063, and 8 neighboring fields in order to investigate the soft X-ray emission in that region of the sky. We find that the soft X-ray emission varies by a factor of few on scales of few degrees, rendering the background subtraction a complex task. If the Noth Polar Spur emission is of local origin, we find that only A2052 and A2063 have evidence of cluster soft excess emission, and that the OVII emission lines detected in XMM observations of A2052 and MKW03s are not associated with the cluster. If part or all of the North Polar Spur soft X-ray enhancement is of extragalactic nature, the three clusters feature strong soft excess emission, and the OVII emission lines observed by XMM are genuinely associated with the clusters.  We interpret the soft excess emission with the presence of warm gas, either intermixed with the hot intra-cluster medium or in filamentary structures located around the clusters, and estimate that the warm gas is approximately as massive as the hot intra-cluster medium. ",
        "introduction": "}  Excess of soft X-ray emission from galaxy clusters, above the contribution from the hot (T $\\sim 10^{7}-10^8$ K) intra-cluster medium (ICM), is commonplace. Analyses of large samples of galaxy clusters (Bonamente et al. 2002; Kaastra et al. 2003a,b; Nevalainen et al. 2003) indicate that the phenomenon is detected in about 30-50\\% of the clusters. Among the clusters that have reported soft excess emission to date are Coma (Lieu et al. 1996a; Bonamente, Joy and Lieu 2003), Virgo (Lieu et al. 1996b), A1795 (Mittaz, Lieu and Lockman 1999; Bonamente, Lieu and Mittaz 2001a), A2199 (Lieu, Bonamente and Mittaz 1999; Bonamente, Lieu and Mittaz 2001a; Kaastra et al. 2001), S\\'ersic 159-03 (Bonamente, Lieu and Mittaz 2001b; Kaastra et al. 2003a,b), a few clusters in the Shapley concentration (Bonamente et al. 2001), two clusters in the Hercules concentration, MKW03s and A2052 (Kaastra et al. 2003a,b), A3112 (Nevalainen et al. 2003) and a few more clusters detected at lower significance (Bonamente et al. 2002).  Detection of soft excess radiation from clusters relies on a correct background subtraction. Inhomogenieties in the diffuse soft X-ray background on a scale of few degrees are common across the whole sky. In addition, variations in the 1/4 keV band are not necessarily correlated with variations in the neighboring 3/4 keV band, indicating that the soft X-ray background originates from a multi-phase plasma (Snowden et al. 2000). The background subtraction task is particularly complex in the case of the Hercules supercluster due to the presence of an extended emission feature along the line of sight, the North Polar Spur.  In this paper we analyze the ROSAT PSPC data of 3 clusters in the Hercules supercluster, MKW03s, A2052 and A2063 and of 8 neighboring background fields. Our PSPC data covers larger areas than the available XMM-Newton observations. With these data, we quantify the level of soft X-ray background fluctuations near the Hercules supercluster, and obtain several {\\it in situ} background measurements in order to investigate the effect of background subtraction on the determination of soft excess fluxes. The ROSAT data presented in this paper reveals important new information concerning the reality and the nature of the excess emission in the Hercules supercluster.  The paper is structured in the following way. In section \\ref{xmm} we provide a summary of the earlier XMM-Newton results for the clusters  MKW03s and A2052, in section \\ref{data} we present the PSPC data and in section \\ref{nps} we present our results on the soft X-ray emission from the North Polar Spur. In section \\ref{cse} we discuss the presence of soft excess emission from the three clusters in the Hercules  concentration, in section \\ref{inter} we provide an interpretation of the excess emission and section \\ref{concl} contains our conclusions. In this paper we use a $\\Omega_b$=0.04, $\\Omega_m$=0.3 and $\\Omega_{\\Lambda}$=0.7 cosmology, and a Hubble constant of $H_0=72$ km s$^{-1}$ Mpc$^{-1}$. The redshifts of the 3 clusters are z=0.045 (MKW03s) and z=0.035 (A2052 and A2063). ",
        "conclusions": "} PSPC observations of three clusters in the Hercules concentration, MKW03s, A2052 and A2063, and of 8 background fields in their vicinity were used to study the presence of soft excess emission. We showed that the  Hercules supercluster lies in a region of the sky -- the southern extension of the North Polar Spur -- that has strong and inhomogenous soft X-ray emission.  The soft excess fluxes of MKW03s, A2052 and A2063 depend on the nature of the North Polar Spur. If the soft X-ray emission from the NPS is of local origin, only A2052 and A2063 show evidence of cluster soft excess. In this case, the soft excess fluxes we measure with PSPC are lower than those reported in the XMM study of Kaastra et al. (2003a), and the OVII emission lines they detected towards MMKW03s and A2052 are not associated with the clusters. If part or all of the NPS emission is of extragalactic origin, the 3 clusters feature strong soft excess emission, in accord with the XMM results.  The soft excess emission is interpreted with the presence of warm gas that can be either intermixed with the hot gas, or located in filamentary structures around the clusters. The PSPC data show that the warm gas feature a temperature of $kT\\leq0.3$ keV. If the warm gas is intra-cluster, it is approximately as massive as the hot ICM, and significantly more massive if located in filamentary structures around the clusters."
    },
    "0512/astro-ph0512628_arXiv.txt": {
        "abstract": "We report on a multiwavelength campaign on the TeV \\gray blazar Markarian (Mrk) 421 performed during December 2002 and January 2003. These target of opportunity observations were initiated by the detection of X-ray and TeV \\gray flares with the All Sky Monitor (ASM) on board the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) and the 10~m Whipple \\gray telescope. The campaign included observational coverage in the radio (University of Michigan Radio Astronomy Observatory), optical (Boltwood, La Palma KVA 0.6m, WIYN 0.9m), X-ray (RXTE pointed telescopes), and TeV \\gray (Whipple and HEGRA) bands.  At TeV energies, the observations revealed several flares at intermediate flux levels, peaking between 1 and 1.5 times the flux from the Crab Nebula. While the time averaged spectrum can be fitted with a single power law of photon index $\\Gamma =2.8$ from $dN_\\gamma/dE\\propto E^{-\\Gamma}$, we find some evidence for spectral variability. Confirming earlier results, the campaign reveals a rather loose correlation between the X-ray and TeV \\gray fluxes. In one case, a very strong X-ray flare is not accompanied by a comparable TeV \\gray flare. Although the source flux was variable in the optical and radio bands, the sparse sampling of the optical and radio light curves does not allow us to study the correlation properties in detail.  We present a simple analysis of the data with a synchrotron-self Compton model, emphasizing that models with very high Doppler factors and low magnetic fields can describe the data. ",
        "introduction": "\\paragraph{} The space-borne EGRET {\\it (Energetic Gamma Ray Experiment Telescope)} detector on board the {\\it Compton Gamma-Ray Observatory} discovered strong MeV and GeV \\gray emission from 66 blazars, mainly from Flat Spectrum Radio Quasars and Unidentified Flat Spectrum Radio Sources  \\citep{Hart:92,Hart:99}. Ground-based Cherenkov telescopes discovered TeV \\gray emission from seven blazars, five of which were not detected by EGRET \\citep{Week:04,Ahar:05}. Although \\gray emission from blazars have been studied for more than a decade now, it is still unclear where and how the emission originates. According to the most common paradigm, the emission originates close to a mass-accreting supermassive black hole, in a relativistically moving collimated plasma outflow (jet) that is aligned with the line of sight to within a few degrees. The relativistic Doppler effect can explain the intensity of the blazar emission, and its rapid variability at X-rays and \\gray energies on hour time scales: for emission originating from synchrotron or synchrotron self-Compton (SSC) models, the apparent luminosity increases approximately as the fourth power of the relativistic Doppler factor\\footnote{The relativistic Doppler factor is given by $\\delta_{\\rm j}\\,=$ $\\left[\\Gamma(1-\\beta\\,\\cos{\\theta})\\right]^{-1}$ with $\\Gamma$, the bulk Lorentz factor of the jet plasma, $\\theta$, the angle between jet axis and the line of sight, and $\\beta$, the plasma velocity, in units of the speed of light.} $\\delta_{\\rm j}$, and the observed flux variability timescale is inversely proportional to $\\delta_{\\rm j}$.  Blazars are powerful sources across the electromagnetic spectrum. Typical spectral energy distributions (SEDs) for the high energy peaked TeV blazars show two broad peaks, one at infrared to X-ray energies and the other at X-ray to \\gray energies. The low-energy peak is commonly believed to originate as synchrotron emission from a population of relativistic electrons gyrating in the magnetic field of the jet plasma. The origin of the high-energy peak is unproven. The commonly adopted and best studied models assume that the \\grays are produced in inverse Compton processes by the same electrons that emit the synchrotron radiation at longer wavelengths (for a recent review of observations and models, see \\citet{Kraw:04r}). In so-called hadronic models, \\grays are emitted as synchrotron radiation of extremely energetic protons \\citep{aharonian,muecke}, as inverse Compton and synchrotron emission from a Proton Induced Cascade (PIC) \\citep{manheim}, or from $\\pi^{0}\\rightarrow$ $\\gamma\\gamma$ decays following the interaction of high energy protons with some target material \\citep{Pohl:00}. Recent reviews on observations of blazars with TeV emission and models developed to describe the data can be found in various review articles and books \\citep{Kraw:04r,Kraw:05,Tave:04,Trev:03,Feli:04}. Reviews focussing on observations and models of sources with MeV/GeV emission can be found in \\citet{Siko:01b,Copp:99}. Broader overviews of the field of TeV \\gray astronomy are given in \\citet{Buck:02,Ong:03,Trev:03,Feli:04}.  Multiwavelength observations are key for understanding the blazar phenomenon. The acquisition of good multiwavelength data sets has encountered substantial difficulties as the sensitivities of current TeV observatories require flares for sampling the TeV light curves on a time scale of hours. Some sources were observed with excellent multiwavelength coverage but during relatively unspectacular quiescent phases; in other cases, the sources were flaring, but the fluxes were only poorly sampled in frequency space and in time. The most remarkable result from the multiwavelength campaigns is that there is good evidence for a correlation between the X-ray fluxes and the TeV \\gray fluxes for the two sources Mrk 421 \\citep{Buck:96,Taka:96,Taka:00,Blaz:05} and Mrk 501 \\citep{Djan:99,Samb:00,Kraw:02}.  In this paper we present results from a multiwavelength campaign on the TeV blazar Markarian (Mrk) 421. The source is a nearby (z = 0.031) high energy-peaked BL Lac object, and was the first extragalactic source detected in the TeV \\gray band \\citep{Punc:92}. % In November 2002, \\gray observations with the Whipple 10~m telescope revealed several Mrk 421 flares with fluxes exceeding three times the steady flux from the Crab Nebula. The All Sky Monitor instrument aboard {\\it RXTE} also showed extremely strong 2-12 keV fluxes reaching 100 milli-Crab. Collectively these triggered a coordinated campaign. We invoked radio, optical, and X-Ray ({\\it RXTE}) observations to commence as soon as the waning Moon would allow the Cherenkov telescope to take data once more. Although the X-ray and TeV \\gray fluxes had decreased substantially when the campaign started on December 4th, we acquired a data set with a high signal-to-noise-ratio X-ray light curve and X-ray energy spectra, and good signal-to-noise-ratio TeV light curves and TeV \\gray energy spectra. The combined data allowed us to study the X-ray/TeV \\gray flux correlation over several weeks. Following our previous study \\citep{Blaz:05}, this is the second campaign that measures the X-ray/TeV \\gray flux correlation over several weekes. For the first time, we use here simulated lightcurves to address the statistical significance of the X-ray/TeV \\gray flux correlation and to constrain the time lag between the two light curves. Simulations are necessary as subsequent data points in the light curve are not independent of each other (see e.g., the discussion by \\citet{edelson}). The rest of the paper is organized as follows.% After describing the data sets in Section 2, we explain the method that we used to reconstruct TeV \\gray energy spectra from the Whipple data, and give the results obtained with the method in Section 3. Subsequently, we present the results of the campaign in Section 4, and conclude with a summary and a discussion in Section 5.  If not mentioned otherwise, errors are quoted on the level of one standard deviation, and upper limits are given on 90\\% confidence level. ",
        "conclusions": "\\label{discussion} \\paragraph{} The multiwavelength campaign showed Mrk 421 in a level of intermediate activity. During the observational campaign, Mrk 421 showed significant flux variability in the radio, optical, X-ray and \\gray bands and significant spectral variability at X-rays and TeV \\grays. While we measured an average TeV \\gray photon index of $\\Gamma\\,=$ 2.8, the observations revealed evidence for spectral variability on a time scale of days. In particular, the data suggest very soft energy spectra with $\\Gamma \\approx 4$ during the first half of the observation campaign. One of the most interesting results from this campaign is that the X-ray and TeV \\gray fluxes are correlated on the $\\sim $97\\% confidence level, but that we find widely different TeV \\gray fluxes for a single X-ray flux and vice versa. The most extreme example is the ``orphan X-ray flare'', seen on January 13, 2003 (MJD 52653). The loose X-ray/TeV correlation may suggest that the model parameters (e.g. the volume of the emission zone) change with time, or that the commonly made assumption of a single synchrotron self-Compton (SSC) emission zone is an over-simplification. Our previous observations of Mrk 421 had already shown a rather loose X-Ray/TeV \\gray correlation \\citep{Blaz:05} and the same applied for 1ES 1959+650 \\citep{Kraw:04}. In the case of Mrk 501, a rather tight correlation has been reported \\citep{Kraw:00,Kraw:02}. \\paragraph{} The analysis of the correlation between the X-ray flux and photon index during a flare indicated a ``clockwise'' hysteresis. For Mrk 421, \\citet{Taka:96} also reported clockwise evolution. However, \\citet{Taka:00} reported evidence for both, clockwise evolution during some flares and anti-clockwise evolution during other flares. If the SSC model indeed applies, these results may imply that the characteristic times scales of the most important processes (acceleration time, radiative cooling time, escape time) change from flare to flare and thus yield the different observed hysteresis behaviours. Recently, \\citet{Soko:04} emphasized that the geometry of the emitting region and its orientation relative to the line of sight influences the observed flux and spectral evolution and might thus further complicate the interpretation of the results.  The X-ray and TeV \\gray emission from Mrk 421 data have been modeled with synchrotron Self-Compton codes by many groups, see e.g. \\citep{Inou:96,Bed:97,Bed:99,Boet:97,Mast:97,Taka:01,Kraw:01,Kono:03,Kino:02,Blaz:05}. A crucial model parameter is the jet Doppler factor $\\delta_{\\rm j}$. The published models with Doppler factors $\\delta_{\\rm j}$ of 20 or lower generally predict TeV energy spectra that are softer than the observed ones, especially if a correction for extragalactic absorption would be applied which steepens the energy spectra considerably. Models with Doppler factors $\\delta_{\\rm j}$ on the order of 50 give satisfactory model fits to both the X-ray (synchrotron) and the TeV (Inverse Compton) emission (see the self-consistent modeling of \\citet{Kraw:01,Kono:03} and the discussions by \\citet{Tave:04,Pine:05}). \\citet{Pine:05} observed the Mrk 421 parsec-scale radio jet with the Very Large Baseline Array (VLBA). Remarkably, they find apparent pattern speeds of only $\\sim 0.1 \\mathrm{c}$. As discussed by the authors, the highly relativistic motion inferred from TeV observations can be reconciled with the modestly relativistic flow calculated from VLBA observations by postulating that the jet slows down between the sub-parsec (TeV) and parsec (VLBA) regimes. It may be possible to describe the multi-wavelength data with a synchrotron-Compton model and lower Doppler factors by invoking additional seed photons. Two new model variants that combine ingredients of SSC and external Compton models have been proposed by \\citet{Geor:03} and by \\citet{Ghis:05}. While the first authors assume that downstream emission regions provide seed photons, the second authors speculate that the jet is composed of a fast spine with a slow-moving envelope. In this model, the fast spine emits the X-ray and \\gray radiation. Modeling of the data with these two inhomogeneous models is outside the scope of this paper.  In this discussion we do not want to embark on comprehensive modeling of the data from the entire campaign. Our main aim is to draw the attention of the reader to a single remarkable fact: while it is difficult to model the data with Doppler factors on the order of 20 and lower, much higher Doppler factors cannot be excluded right away. Fig.\\ \\ref{nufnu} shows two synchrotron self-Compton models based on the simple snapshot code of  \\citet{Kraw:04}. The code assumes a single spherical emission volume of radius $R$ relativistically approaching the observer with a jet Doppler factor $\\delta_{\\rm j}$. The emission volume is homogeneously filled with a tangled magnetic field of strength $B$ and a non-thermal electron population. The electron energy spectrum follows $dN/d\\gamma\\propto$ $\\gamma^{-p}$ with $p\\,=$ 2 for electron Lorentz factor $\\gamma$ between $\\gamma_{\\rm min}$ and $\\gamma_{\\rm b}$ and $p\\,=3$ for Lorentz factors between $\\gamma_{\\rm b}$ and $\\gamma_{\\rm max}$. The code models extragalactic absorption owing to the $\\gamma_{\\rm TeV} + \\gamma_{\\rm CIB} \\rightarrow e^+ e^-$ pair-production processes of TeV photons on photons from the Cosmic Infrared Background (CIB) using the CIB model of \\citet{Knei:02}.  We discuss two models. We show the first model for illustrative purposes only. It uses the \"conventional\" model parameters ($\\delta_{\\rm j}\\,= 50$) derived from the time-dependent self-consistent analysis of a different but similar data set \\citep{Kraw:01}. The second model uses a very high Doppler factor ($\\delta_{\\rm j}\\,= 1000$). All the model parameters are given in Table \\ref{fits}. For both models, we assured that the model parameters were chosen self-consistently. Causality arguments require that the radius $R$ satisfies $R<\\delta_{\\rm j}c\\Delta T_{\\rm obs}\\,=\\,2.7\\times 10^{15}$~cm for $\\delta_{\\rm j}\\,=50$ and $R<5.4\\times10^{16}$~cm for $\\delta_{\\rm j}\\,=1000$ for a flux variability time scale of $\\Delta T_{\\rm obs}\\,=$ 30~min. Note that the flux variability time scale sets a lower limit on $R$ but no upper limit, if the flux variability time scale is not dominated by light travel time effects but by other effects (e.g.\\ by the stability of a strong shock front). We checked that the SED (i) fits the X-ray and TeV \\gray data, and (ii) is consistent with the expected spectral shape owing to radiative cooling. In the first model, the latter self-consistency is assured by our previous self-consistent time-dependent modeling. In the second model, we construct an electron spectrum based on the general results for electrons suffering synchrotron and Inverse Compton losses \\citep{Syro:59,Kard:62,Ginz:64,Pach:70,Inou:96}). We assume that the electron energy spectrum breaks at $\\gamma_{\\rm b}\\,=$ $1.8\\times 10^3$. The laboratory-frame synchrotron cooling time for electrons at the break is $t_{\\rm s}\\,=$ $[\\frac{4}{3}$ $\\sigma_{\\mbox{\\small T}}\\,$ $c\\,$ $\\delta_{\\rm j}\\,$ $\\frac{B^2}{8\\pi\\,m_{\\rm e}\\,c^{2}}\\,$ $\\gamma_{\\rm b}$ $]^{-1}~$ $\\approx\\, 28min $ ($\\sigma_{\\mbox{\\tiny T}}$ is the Thomson cross section and $m_{\\rm e}$ is the electron mass). An electron spectrum as the one used here could result from the radiative cooling of a $p\\,=2$ electron energy spectrum that extends from $\\gamma_{\\rm min}$ to $\\gamma_{\\rm max}$, resulting in a spectrum with $p\\,=2$ and $p\\,=3$  below and above $\\gamma_{\\rm b}$, respectively.  As shown in Fig.\\ \\ref{nufnu}, both models fail to predict the observed radio fluxes as a consequence of synchrotron self-absorption. We would like to emphasize that we do not regard the discrepancy as a shortcoming of the model. Electrons producing the radio emission cool on much longer time scales, and the radio emission will be dominated by an accumulation of downstream plasma which stopped contributing to the X-ray and TeV emission long time ago. We could model the radio emission with another emission component (see e.g.\\ \\citet{Blaz:05}). However, doing so is arbitrary: for a small number of data points we would add an additional model component with many free model parameters.  In Table \\ref{fits} we list for both models the magnetic field energy density $u_{\\rm B}$, the energy density in electrons $u_{\\rm e}$, the ratio $r = u_{\\rm e}/u_{\\rm B}$, and the kinetic luminosity $L_{\\rm k} = \\pi R^{2} c \\Gamma^{2} (u_{\\rm e} + u_{\\rm B})$ for $\\Gamma\\,=$ $\\delta_{\\rm j}$ \\citep{Bege:94}. The model with a low Doppler factor is closer to equipartition between magnetic field and particles. The kinetic luminosity is similar for the two models, with a high radiative efficiency of the low-$\\delta_{\\rm j}$ model making up for the stronger boosting of the high-$\\delta_{\\rm j}$ model.  If taken seriously, the model with $\\delta_{\\rm j}\\,=$ 1000 would imply that the X-ray and TeV \\gray emission is produced by an ultra-relativistic particle dominated wind, very close to the supermassive black hole. The fact that seven blazars have been detected at TeV energies seems to argue against extremely relativistic outflows with bulk Lorentz factors $\\Gamma$ on the order of 1000, as isotropic emission would be beamed into an opening angle of $\\Gamma^{-1}$ and would make the observation of the emission unlikely. However, the argument only applies if the jet opening angle is equal or smaller than $\\Gamma^{-1}$. Having a larger jet opening angle would require a higher total jet-luminosity as some jet segments would not contribute to the observed emission. However, the jet-luminosities listed in the table are several orders of magnitude below the Eddington luminosity of a $\\sim 10^{8.4}$ solar mass black hole that is suspected to be at the core of Mrk 421 \\citep{Bart:03,Falo:02}."
    },
    "0512/astro-ph0512302_arXiv.txt": {
        "abstract": "\\  Most star complexes are in fact  complexes of stars, clusters and  gas clouds; term \"star complexes\" was introduced as general one disregarding the preferential content of a complex. Generally the high rate of star formation in a complex is accompanied by the high number of bound clusters, including massive ones, what was explained by the high gas pressure in such regions. However, there are also complexes, where clusters seems to be more numerous in relation to stars than in a common complex. The high rate of clusters - but not isolated stars - formation seems to be typical for many isolated bursts of star formation, but  deficit of stars might be still explained by the observational selection. The latter cannot, however, explain the complexes or the dwarf galaxies, where the high formation rate of only stars is observed.  The possibility of the very fast dissolution of parental clusters just in such regions should itself be explained. Some difference in the physical conditions (turbulence parameters  first of all) within the initial gas supercloud might be a reason for  the high or low stars/clusters  number ratio in a complex. ",
        "introduction": "The common opinion is that  formation of star clusters is the only way of star  formation. Since 1950s we know that the stars forms by groups, in clusters and associations.  The presently isolated stars might be formed  in groups, which have been dissolved till the present time.  Most star complexes are in fact  complexes of stars, clusters and  gas clouds; term \"star complexes\" was initially suggested for all of them, disregarding what is their preferential or more evident population [1, 2]. However, later on we have presented evidences for existence of complexes, where   clusters seems to be more numerous in relation to stars than in a common complex [3].  Some difference in the physical conditions within the initial gas supercloud might be a reason for such high C/S number.  We present here more examples of locations where the rate of cluster formation seems do not correlate well with the rate of the isolated star formation. These cases are to be studied in more details. The high C/S ratio might be explained by the observational selection, but it is more difficult for the low C/S ratio. The plausible explanation of low C/S ratios by the very fast dissolution of parental clusters just in such regions should itself be explained. ",
        "conclusions": "The superassociatons or the localized bursts of star formation are nothing but the star/cluster complexes, involved  totally in the intensive star formation process; most of them are the supergiant HII regions. However, it is not the case for NGC 205 = OB 78 superassociation in M31.  This ~1 kpc in size region is  filled with OB-stars; a few small HII regions are near, but outside, of it (fig. 4). This was explained long ago by the lack of the gas inside OB78 - the gas was therefore expulced??? from the superassociation  by the pressure from O-stars/SNe earlier than  it was ionized. This plausibly  was due to the high density of these stars; this suggestion might be checked by the observational data on stars and gas in this  region - and also in M51, where the very various interrelations between HII gas and star/cluster complexes are seen and are to be studied.  Anyway, our present goal is to stress the absence of star clusters inside OB 78. This is  an evident contradiction to the high pressure explanation  of the formation of the bound massive clusters. Moreover, this case is  dissimilar to many other localized  bursts of star formation, where just the YMCs, if not SSCs, dominate, like it is the known case in the Antennae galaxies, and, to lesser extension, inside the M101 supergiant HII regions.  Somewhat similar case is IC 10, the dwarf galaxy of the Local group.  The  current  rate of star formation there is the highest of the all Local Group galaxies, but no YMCs are known there  (Grebel, [18]). As S.Larsen has commented, \"may be there is someting peculiar going on there\" ([18], p. 427).  The very low  cluster formation rate is known in the irregular Local Group galaxy IC 1613. Being  normalized to the same star-formation rate, it is 600 times lower than in the LMC, which is a galaxy of the same type  [19]. The quite low number of star clusters in IC 1613 has been known since W.Baade's investigations in Fifties. This led Hodge (Galaxies, 1986) to conclusion that the unrecognised agent should exist, determining whether the formation of the populous clusters is possible in a galaxy or not. Anyway, 13 OB-associations and even a superassociation, noted by W.Baade, are  known in IC 1613.  Therefore, it looks like this agent is not always the dierect consequence of the general SFR.  There may exist not only local, but also temporal difference in the clusters and stars rate of formation. At least in the LMC it is the case. The break in the formation of (at least massive) clusters, which  lasted in the LMC for 4--14 Gyr, was not matched by a decrease in the star-formation rate (van den Bergh [20]). The conclusion by van den Bergh was: \"The dramatic contrast between the history of the cluster formation and that of field stars suggests that star clusters cannot be used as proxies for star formation\".  The steep IMF found by Massey [21] for the seemingly isolated stars of the LMC field might be considered the indication of different mode of formation for isolated stars, but recently  Elmegreen and Scalo [22] concluded that this steeper IMF resulted from unappropriate assumption of the constant SFR, which is hardly is justified for the field stars.  All in all, these data demonstrate  the important and often unrecognised problem does exist  and should be carefully investigated. The best approach to this issue  seems to be studying the populations of the same age within isolated star/cluster complexes in resolvable galaxies. (The data on the MW galaxy suffered from many selection effects and uncertainty of distances; anyway, they suggest also  the existence of complexes of clusters,  like the compact group in Cassiopeia, described in [3], p. 204).  The objects within a complex have arised from the (initially) bound gas supercloud and the ISM properties within  it must be reflected in stars/clusters number ratio  over all the complex. This ratio for the objects of the same age should be determined for  the large number of the best outlined complexes  in different galaxies and at different location within a galaxy.  The Cepheid investigations would be valuable, yet considering these are time-consuming and age-limited, the data on the  bright stars could be used up to the same magnitude limit  in a galaxy.  The position of a complex in a galaxy   might be  connected with the value of stars/clusters ratio inside the complex. There are indeed some guesses that this might be the  case. OB78 in M31 seems to locate at the cross  of two spiral arms going to opposite directions. In the LMC, the NGC 2164 group of clusters is in isolated position  far from the LMC center, whereas  the NGC 2058/2065 group is near the tip of the bar and the dense group of Cepheids  is to SW of it.   The  bright complexes are often at the tips of  the spiral arms, whereas  in the grand design galaxies complexes are often  located at equal spacing along the arm,  suggesting the  formation under action of Parker instability. The different mechanisms of formation may be reflected in the different clusters/stars ratio and parameters of turbulence in ISM might depend on a position inside a galaxy.  The peculiar shape may also be connected with the certain mechanism of formation and C/S ratio.  The arc-shaped complexes might be formed in result of the dynamical pressure. The arc-like shape has  been observed too at the all-galaxy scale, for the leading edges of a number of galaxies moving through the IGM, i.g. for DDO 265 in M81 group (fig. 5). More examples are given in  [23]. Note that the  arcs of Sextant in the LMC and that of the Western complex in M83 consist  of 5 - 7 clusters each,  whereas the older arc of Quadrant  consists from  both clusters and stars [24]. The star formation triggered by the high (especially dynamical?) pressure may result mostly in the bound clusters, whereas gravitational fragmentation of a supercloud probably lead to the normal (most often observed) C/S ratio. This is surely the most frequent case of a star/cluster complex formation. However, the low C/S ratio in regions presumably  formed under the high pressure conditions, which  demonstrate the high SFR, seems to be not only rare, but also enigmatic cases.  The issue of the mode of star formation might have the deep cosmological implications. For example, the formation of the YMC (and more so of the SSC) in the BCD galaxies, many of which are isolated, miight be triggered by gas infalling in dark-matter haloes; this gas may experience sloshing  and oscillation favoring compressiion and instabilities [25].   This work has been done with the support of Russian Foundation for the Basic Investigations, project RFFI 03-02-16288."
    },
    "0512/astro-ph0512134_arXiv.txt": {
        "abstract": "Planetary and stellar dynamos likely result from turbulent motions in magnetofluids with kinematic viscosities that are small compared to their magnetic diffusivities. Laboratory experiments are in progress to produce similar dynamos in liquid metals. This work reviews recent computations of thresholds in critical magnetic Reynolds number above which dynamo amplification can be expected for mechanically-forced turbulence (helical and non-helical, short wavelength and long wavelength) as a function of the magnetic Prandtl number $P_M$. New results for helical forcing are discussed, for which a dynamo is obtained at $P_M=5\\times10^{-3}$. The fact that the kinetic turbulent spectrum is much broader in wavenumber space than the magnetic spectrum leads to numerical difficulties which are bridged by a combination of overlapping direct numerical simulations and subgrid models of magnetohydrodynamic turbulence. Typically, the critical magnetic Reynolds number increases steeply as the magnetic Prandtl number decreases, and then reaches an asymptotic plateau at values of at most a few hundred. In the turbulent regime and for magnetic Reynolds numbers large enough, both small and large scale magnetic fields are excited. The interactions between different scales in the flow are also discussed. ",
        "introduction": "INTRODUCTION}  Plasmas in stellar interiors and conducting fluids in planetary cores are characterized by a magnetic Prandtl number $P_M$ (the ratio of the kinematic viscosity $\\nu$ to the magnetic diffusivity $\\eta$) much smaller than one. As a few examples, the magnetic Prandtl number in the solar convective region is estimated to be $P_M \\approx 10^{-5}-10^{-6}$ \\cite{Parker79}, and in the Earth's core $P_M \\approx 10^{-5}$. Liquid sodium experiments are also characterized by small values of $P_M$.  While numerical simulations of dynamo action in these objects are available, the large values of the kinetic ($R_V$) and magnetic ($R_M$) Reynolds numbers forbid a study using realistic values of $P_M$. Simulations of the geodynamo \\cite{Glatzmaier99} or the solar convective region \\cite{Cattaneo99} are often done for $P_M \\sim 1$. While the proper separation of the kinetic and magnetic dissipation scales cannot be achieved in these simulations, values of $P_M$ much smaller than one can be reached under more idealized conditions. Pseudospectral methods in periodic boxes give an excellent tool to study the behavior of magnetohydrodynamic (MHD) turbulence in the regime $P_M<1$. The further assumption of incompressibility allows for an extra gain in computer power. Paralellized pseudospectral codes have reached for hydrodynamic turbulence resolutions of $4096^3$ grid points, and Taylor Reynolds numbers of $R_\\lambda \\approx 1200$ \\cite{Kaneda03}. These methods are conservative and nondispersive, being well suited for the exploration of turbulent flows in regimes hard to explore in the laboratory. And under some circumstances, the range of values of $P_M$ can be extended using subgrid scale (SGS) models.  In this work, we review recent results from simulations of helical and non-helical dynamos at $P_M<1$ using pseudospectral codes in periodic boxes \\cite{Ponty05,Mininni05c,Mininni05e}. To extend the range in $P_M$ in the simulations, SGS models were used in these works. We discuss some of these models with particular emphasis in the Lagrangian Average MHD (LAMHD) equations \\cite{Holm02a,Holm02b,Montgomery02,Mininni05a,Mininni05b}. In addition, and to validate results from SGS models, new results from direct numerical simulations (DNS) with resolutions of $1024^3$ grid points are presented.  We focus on the properties of the magnetic field in the kinematic dynamo regime, when the intensity of the magnetic field is small and the effect of the Lorentz force can be neglected. In particular, we discuss the behavior of the threshold in $R_M$ for dynamo action as $P_M$ is decreased. That is to say, given a hydrodynamic state and an arbitrary small magnetic perturbation, what is the minimum value of $R_M$ (or maximum value of $\\eta$, in some convenient set of dimensionless units) to have a dynamo instability such that the system reaches after a finite time a magnetohydrodynamic steady state. Below this threshold, the magnetic perturbation is damped and the final state of the system is hydrodynamic.  Two flows have recently been studied in this context: the flow resulting from Taylor-Green forcing \\cite{Ponty05,Mininni05c}, and the result of Roberts forcing \\cite{Feudel03,Mininni05e}. The first case corresponds to a flow with no net helicity that gives large scale dynamo action, while the former studies are for a helical flow where only small scale dynamo action is permitted by introducing mechanical energy in the largest available scale. In a different context, for isotropic, homogeneous, and delta-correlated in time external forcing, the problem has also been studied in Refs. \\cite{Haugen04,Schekochihin05}. In addition to reviewing this results, we compare them against new simulations using Arn'old-Childress-Beltrami (ABC) forcing with energy injection at intermediate scales. This is a helical case where large scale magnetic amplification is allowed. For this forcing, values of $P_M$ down to $5\\times10^{-3}$ are reached. The results obtained for such a low value of $P_M$ are expected to be of relevance for astrophysical and geophysical applications, as well as for laboratory dynamos.  For all cases where a large scale flow is present, dynamo action is observed to persist at the smallest values of $P_M$ that can be reached. Moreover, for values of $P_M$ smaller than $\\sim 0.1$ an independence of the threshold with $P_M$ is observed. While for the Taylor-Green (non-helical) forcing and the Roberts forcing (helical, but with magnetic amplification only at small scales) a sharp increase in the critical parameter is observed before reaching the asymptotic regime, in the ABC case almost no such increase is found.  The structure of the paper is as follows. In Sec. \\ref{sec:equations} we present the equations, and the several forcing functions used. We also describe the code and the SGS models. Section \\ref{sec:thresholds} discusses the thresholds for dynamo action for the several flows, and Sec. \\ref{sec:scales} discusses the role played by the different scales in the problem. Finally, Sec. \\ref{sec:conclusions} presents the conclusions. ",
        "conclusions": "CONCLUSIONS}  We reviewed results of dynamo action at low magnetic Prandtl number \\cite{Ponty05,Mininni05c,Mininni05e} for several mechanical forcing functions, including helical and non-helical flows, as well as small and large scale dynamo action. For all cases where a large scale flow is present, a similar behavior is obtained in the threshold for dynamo action: as $R_V$ is increased (or $P_M$ decreased), the value of $R_M^c$ increases sharply as turbulence develops and finally reaches an asymptotic regime independent of the value of $P_M$. The large scale flow plays an important role in the establishment of the asymptotic behavior, as shown by the anti-correlation between the characteristic length of the flow and $R_M^c$. As turbulence develops, the laminar flow creates small scales through hydrodynamic instabilities, and the large scale laminar flow loses its infinite correlation time. Then, both the large and small scale velocity fields amplify the magnetic field, giving rise to the asymptotic regime.  New results from simulations with ABC forcing present a distinctive behavior. Having the flow maximum kinetic helicity and permitting large scale dynamo action, the critical magnetic Reynolds is smaller than for the other two flows by one order of magnitude. Also, only a twofold increase in $R_M^c$ is observed as turbulence develops (in contrast to a tenfold increase for the other flows). As a result, dynamo simulations close to the threshold do not show small scale amplification (down to $P_M = 5\\times 10^{-3}$), and only the large scale flow is responsible for the large scale dynamo action. However, as the value of $R_M$ is increased, as is expected for astrophysical and geophysical flows, small scale amplification in the kinematic regime is recovered.  Although the conditions in our simulations are idealized, we believe the existence of an asymptotic regime for $P_M<1$ has profound implications for laboratory experiments and modeling of astrophysical and geophysical dynamos. In most of these cases, a large scale flow (such as differential rotation) and turbulent fluctuations are known to be present."
    },
    "0512/astro-ph0512652_arXiv.txt": {
        "abstract": "We study the hydrodynamical transition from an hadronic star into a quark or a hybrid star. We discuss the possible mode of burning, using a fully relativistic formalism and realistic Equations of State in which hyperons can be present.  We take into account the possibility that quarks form a diquark condensate.  We also discuss the formation of a mixed phase of hadrons and quarks, and we indicate which region of the star can rapidly convert in various possible scenarios.  An estimate of the final temperature of the system is provided.  We find that the conversion process always corresponds to a deflagration and never to a detonation.  Hydrodynamical instabilities can develop on the front.  We estimate the increase in the conversion's velocity due to the formation of wrinkles and we find that, although the increase is significant, it is not sufficient to transform the deflagration into a detonation in essentially all realistic scenarios.  Concerning convection, it does not always develop. In particular the system does not develop convection if hyperons are not present in the initial phase and if the newly formed quark phase is made of ungapped (or weakly gapped) quarks. At the contrary, the process of conversion from ungapped quark matter to gapped quarks always allows the formation of a convective layer. Finally, we discuss possible astrophysical implications of our results. ",
        "introduction": "The possible existence of compact stars partially or totally made of quarks has been proposed many years ago \\citep{witten,Bodmer:1971we,Itoh:1970uw}.  A likely origin of these compact stellar objects is the conversion of a purely nucleonic (or hadronic) star into a star containing deconfined quark matter through a quark deconfinement phase transition.  Several works have discussed when, during the life of the compact star, the deconfinement transition should likely take place, either soon after the formation of the neutron star in the supernova explosion \\citep{Lugones97,Benvenuto99}, or during the cooling of the protoneutron star \\citep{Pons:2001ar}. The formation of quark matter could also be delayed, if the deconfinement process takes place through a first order transition so that the purely hadronic star can spend some time as a metastable object \\citep{Berezhiani:2002ks,Drago:2004vu,Bombaci:2004mt}.  A rather controversial question concerns the duration of the deconfinement process itself.  In the first paper discussing this problem \\citep{Olinto86} it was assumed that the conversion proceeds as a slow combustion and it was found that the velocities involved strongly depend on the temperature of the star and they can be small if the star is hot. Soon after the seminal work of Olinto, \\citet{Horvath88} studied the stability of the process and they found that in the presence of gravity the combustion front becomes instable.  These authors also stressed that since the appearance of instabilities increases the velocity of the combustion, it is possible that the slow combustion becomes a detonation.  On the other hand, no estimate of the velocity of the conversion front, taking into account the instability, was presented and the fluidodynamics equations were written in a non-relativistic frame.  The first relativistic calculation of the conversion process was done by \\citet{Cho93}.  To determine which type of conversion takes place, either a detonation or a deflagration, they studied the conservation of the energy-momentum tensor and of the baryonic flux through the conversion front.  Using very simple Equations of State (EoSs) for hadronic matter (the Bethe-Johnson and the Fermi-Dirac EoSs) and the MIT-bag model for quark matter (and without considering the possibility of a mixed phase) they found that the conversion is never a detonation and only for very special values of the parameters it is a slow combustion.  For nearly all the parameters' values they obtained an unstable front.  In the same relativistic scheme \\citet{Lugones94} studied the case in which the combustion front is preceded by a precompression wave. In this scheme they fixed the velocity of combustion and used the conservation law to determine the final temperature of the combusted phase. Also in this case (as for \\citet{Cho93}) the combusted phase was assumed to be pure strange matter in $\\beta$-equilibrium, described by the MIT-bag model and it was not considered the possibility of a mixed phase.  In an other work \\citet{Lugones02} found that Rayleigh-Taylor instability can significantly increase the velocity of the combustion; moreover, the presence of a magnetic field can generate a strong asymmetry in the propagation flame, with the maximum velocity in the polar direction.  In our investigation of the conversion process we assume that the formation of quark matter (QM) takes place inside a relatively cold and $\\beta$-stable compact star.  This scenario is compatible both with the works in which quark deconfinement takes place in a protoneutron star immediately after deleptonization \\citep{Pons:2001ar}, and also with the possibility of a delayed formation of QM, as discussed in \\citet{Berezhiani:2002ks,Drago:2004vu,Bombaci:2004mt,vidana:2005}.  We will study the conversion problem from the fluidodynamical point of view \\citep{landau,Cho93}, so that it will be possible to determine the type of conversion from nucleonic (or hadronic) matter to QM. Following \\citet{landau} there are four possibilities: detonation, weak detonation, deflagration (or slow combustion) and strong deflagration.  The fastest process is detonation, and in that case the front velocity is suprasonic. It is therefore interesting to explore the possibility that, at least in the case of detonation, the conversion process is so rapid that only strong interactions can take place near the front during the conversion, while weak interactions will restore $\\beta$-equilibrium only after the front. The idea to distinguish between fast processes (typically mediated by the strong interaction) and slow processes (due to weak interaction) is common to other physical situations, as, e.g. the computation of viscosity (see \\citet{Lindblom:2001hd,Drago:2003wg}).  In the literature it has been shown that detonation is difficult to achieve when realistic EoSs are used and matter is assumed to reach $\\beta$-equilibrium during the conversion process \\citep{Cho93}. Since the EoS of matter is stiffer if $\\beta$-processes are forbidden, then, in principle, it could be easier to obtain a detonation assuming that matter immediately after the front is not yet in $\\beta$-equilibrium.  In this paper we will consider both the possibility that slow weak interactions take place only after the conversion to QM (which is due to rapid and flavor conserving strong interactions) and also the situation in which $\\beta$-equilibrium is immediately restored. It is also important to note that neutrino trapping delays $\\beta$-stability. Although in this paper we will not discuss neutrino trapping, an estimate of its importance can be obtained comparing the $\\beta$-stable with the not $\\beta$-stable scenario. Concerning temperature, one also need to check both possibilities, namely that the detonation front is so rapid that matter immediately after the conversion front is still cold, and the possibility that strong interactions have enough time to thermalize the newly formed phase.  The situation is rather different when analyzing deflagration, which is a subsonic process. Since for deflagration heat transmission from the burned to the unburned zone can be crucial, one must also consider burned matter at finite temperature.  The main aim of our work is to study the deconfinement process using realistic EoSs, in particular we will take into account the possibility of forming a mixed phase of quarks and hadrons and we will also discuss the effect of the formation of a diquark condensate. Another important open question that we will investigate is the possibility for convection to develop. As we will show, although the conversion front turns out to be always unstable to gravity-induced Rayleigh-Taylor instability, convection can actually develop only for specific choices of the EoSs, for instance the presence of hyperons in the hadronic phase or the formation of a diquark condensate in the quark phase will help the formation of a convective layer.  The outline of our paper is the following: in Sec.~\\ref{fluid-eqs} we discuss the fluidodynamics of the conversion process; in Sec.~\\ref{eos} we present the EoSs used both for hadronic and for QM and the structure of the resulting compact star. In this section we also discuss the thermal formation of mixed-phase; in Sec.~\\ref{temperature} we estimate the temperature reached by the quark (or mixed) phase after deconfinement; in Sec.~\\ref{deflagration} we show the results of our analysis aiming at classifying the type of conversion; in Sec.~\\ref{instability} we discuss hydrodynamical instabilities and their effect on the conversion velocity; in Sec.~\\ref{convection} we discuss the conditions allowing convection to take place and, finally, in Sec.~\\ref{results} we summarize our findings and discuss the astrophysical implications of our work. ",
        "conclusions": "\\label{results}  To clarify the astrophysical implications of our work let us discuss two scenarios for the formation of QM in which our formalism can be applied.  In the first scenario the surface tension at the interface between hadrons and quarks is negligible or vanishes, i.e. the star can not become metastable respect to the formation of a drop of QM ($\\sigma$ smaller then a few MeV/fm$^2$).  In this case the formation of QM takes place soon after the supernova explosion. In agreement with the analysis of \\citet{Pons:2001ar}, QM starts forming when the proto-neutron star has deleptonized and its temperature drops down to a few MeV.  At the center of the proto-neutron star a drop of QM can form and, if $B$ is not too small, a finite strangeness content is needed for the deconfinement process to be exothermic. The strangeness content could be already present in the hadronic star due to the formation of hyperons.  When the drop starts expanding, the process of conversion can be extremely fast within the layer in which deconfinement is energetically convenient even in the absence of weak processes.  In this case the conversion front moves at the velocity of the deflagrative front $v_{\\mathrm{df}}$, which approaches the velocity of sound (see Sec.~\\ref{deflagration}). As shown in Figs.~\\ref{fig-eos155}, \\ref{fig-eos165}, only at very large densities the EoS of not $\\beta$-stable matter is composed of pure quarks. As the conversion layer moves outward, the front enters the region of mixed phase where $v_{\\mathrm{df}}$ decreases till it vanishes at the low density boundary of the mixed phase. Before reaching that point the conversion process involving weak reactions becomes first competitive and then dominant. The diffusion of strangeness is a relatively slow process whose velocity can nevertheless be significantly increased by hydrodynamical instabilities. In this scenario heat diffusion plays a marginal role.  In the alternative scenario the surface tension is larger and the hadronic star can therefore become metastable. In the region in which pure quark matter can form there is no substantial difference with the previous scenario although heat diffusion can be necessary to have a rapid expansion of the pure QM (see discussion at the end of Sec.~\\ref{tn}). When the region of mixed phase is reached, the only way of rapidly producing this new phase is via thermal nucleation. In this scenario the deflagrative velocity is therefore limited by the heat diffusion velocity if strangeness need not to diffuse. Again strangeness diffusion will be crucial to convert the outer layers of the star.  If the surface tension exceeds few ten MeV/fm$^2$ the process of formation of mixed phase is extremely slow and it cannot be described using a fluidodynamical scheme.  The main result of our analysis, based on realistic EoSs, is that the conversion from hadronic matter to QM, or to a mixed phase of hadrons and quarks and also the transition from unpaired QM to gapped QM always takes place as a deflagration.  This result does not change if a finite temperature of the system is taken into account. In our analysis we have shown that the maximum temperature obtained is $\\sim$ 50 MeV near the center of a massive star.  For such a relatively low temperature the system remains strongly degenerate and thermal effects are small. To estimate the increase of the conversion velocity due to hydrodynamical instabilities we have used a fractal scheme. Although the wrinkles which develop on the front surface can significantly increase the conversion velocity, in most realistic cases the process remains subsonic and the transformation from deflagration to detonation does not take place. Concerning the possibility of developing convection, this is possible if hyperons are present and if $B$ is not too large and the mass of the compact star not too small. Convection can also develop if quarks can form a condensate. In particular, in the conversion from ungapped to gapped QM convection always takes place.  Let us now discuss two astrophysical problems in which the type of conversion, either deflagrative or detonative, and the conversion velocity play a crucial role.  \\subsubsection*{Neutron star velocities} It has been proposed by \\citet{Bombaci:2004nu} that the high velocities displayed by some neutron stars can be attributed to an asymmetric neutrino emission associated with the formation of QM inside the hadronic star.  The origin of this asymmetry could be related to a process of deconfinement starting off the center of the star.  In our analysis we have shown that if hyperons are present or if a diquark condensate forms then convection can develop. The possibility of rapidly transporting hot material to the surface of the star via the formation of a convective layer can indeed be at the origin of strong asymmetries in the conversion process\\footnote{The model proposed in \\citet{Bombaci:2004nu} and here discussed has {\\it no connection} with models in which the kicks are explained as due to parity violating processes in the presence of a strong magnetic field, a mechanism which is known to provide almost no contribution to the neutron star velocity.}.  \\subsubsection*{Gamma Ray Bursts} It has been speculated several times \\citep{Cheng:1995am,Bombaci:2000cv,Wang:1999mf,Ouyed:2001cg, Berezhiani:2002ks,Bombaci:2004mt, vidana:2005} that the so-called long GRBs can be originated by the conversion of an hadronic star into a star containing deconfined QM, either as a pure phase or in phase in which quarks are mixed with hadrons. Moreover, as discussed in \\citet{Drago:2005qb} (see Sec.~\\ref{doublescenario} and the analysis of the time-structure of the light curves of GRBs presented in \\citet{Drago:2005cc}) the conversion process can take place in two steps, with a first transition from hadrons to ungapped (or 2SC) quarks and a second transition in which a CFL phase is produced. In order to associate an emission peak with each of the two transitions, the conversion process must be rapid enough to deposit in a few seconds (or less) a huge energy inside the star. Neutrinos will then transport the energy to the exterior on a time scale of order $(10\\div 20)$~s. Clearly, the result of our calculation provides these large velocities, since the conversion process occurs on a time scale of $(0.1\\div 1)$ s for the first transition in the case of a laminar front and it is much more rapid if the hydrodynamical instabilities are taken into account. The second transition lasts only some 10$^{-3}$ s due to the formation of a convective layer. If the two processes takes place one after the other it is even possible that the formation of diquark condensate accelerates the conversion process by developing a convective layer inside the hadronic phase.  \\noindent It is also important to recall that the way in which the conversion to quark matter takes place, either via a detonation or a deflagration, is crucial. It has been shown that the mechanical wave associated with a detonation would expel a relatively large amount of baryon from the star surface \\citep{fryer}. In the case of a detonation the region near the surface of the compact star where the electron-photon plasma forms (via neutrino-antineutrino annihilation) would be contaminated by the baryonic load and it would be impossible to accelerate the plasma up to the enormous Lorentz factors needed to explain the GRBs."
    },
    "0512/astro-ph0512187_arXiv.txt": {
        "abstract": "We present the first large-scale radiative transfer simulations of cosmic reionization, in a simulation volume of $(100\\,h^{-1}\\rm Mpc)^3$. This is more than a 2 orders of magnitude improvement over previous simulations. We achieve this by combining the results from extremely large, cosmological, N-body simulations with a new, fast and efficient code for 3D radiative transfer, $\\rm C^2$-Ray, which we have recently developed. These simulations allow us to do the first numerical studies of the large-scale structure of reionization which at the same time, and crucially, properly take account of the dwarf galaxy ionizing sources which are primarily responsible for reionization. In our realization, reionization starts around $z\\sim21$, and final overlap occurs by $z\\sim11$. The resulting electron-scattering optical depth is in good agreement with the first-year WMAP polarization data. We show that reionization clearly proceeded in an inside-out fashion, with the high-density regions being ionized earlier, on average, than the voids. Ionization histories of smaller-size (5 to 10 comoving Mpc) subregions exabit a large scatter about the mean and do not describe the global reionization history well. This is true even when these subregions are at the mean density of the universe, which shows that small-box simulations of reionization have little predictive power for the evolution of the mean ionized fraction. The minimum reliable volume size for such predictions is $\\sim30$ Mpc. We derive the power-spectra of the neutral, ionized and total gas density fields and show that there is a significant boost of the density fluctuations in both the neutral and the ionized components relative to the total at arcminute and larger scales. We find two populations of H~II regions according to their size, numerous, mid-sized ($\\sim10$ Mpc) regions and a few, rare, very large regions tens of Mpc in size. Thus, local overlap on fairly large scales of tens of Mpc is reached by $z\\sim13$, when our volume is only about 50\\% ionized, and well before the global overlap. We derive the statistical distributions of the ionized fraction and ionized gas density at various scales and for the first time show that both distributions are clearly non-Gaussian. All these quantities are critical for predicting and interpreting the observational signals from reionization from a variety of observations like 21-cm emission, Ly-$\\alpha$ emitter statistics, Gunn-Peterson optical depth and small-scale CMB secondary anisotropies due to patchy reionization. ",
        "introduction": "Understanding the large-scale geometry of reionization (sometimes also referred to as topology of reionization), i.e.\\ the size- and spatial distribution of the ionized and neutral patches, and their evolution in time is one of the most important problems in this fast-developing field.  Better understanding of this geometry is crucial for making detailed and realistic predictions for the observational features of reionization. Detecting these features is the goal of a number of upcoming meter-wavelength radio synthesis telescopes, such as PAST\\footnote{\\tt http://web.phys.cmu.edu/$\\!\\sim\\!$past/}, LOFAR\\footnote{\\tt http://www.lofar.org}, MWA\\footnote{\\tt http://web.haystack.mit.edu/arrays/MWA}, and SKA\\footnote{\\tt http://www.skatelescope.org}. Also being planned are observations of small-scale CMB anisotropies created by ionized patches \\citep[e.g.][]{2003ApJ...598..756S}, and direct observations of high-redshift Ly-$\\alpha$ emitters and their clustering using either James Webb Space Telescope or ground-based telescopes \\citep[e.g.][]{2004ApJ...617L...5M,2005ApJ...619...12S}. Such observations can in principle map the complete progress of reionization through time and space.   In the last few years a variety of different cosmological radiative transfer methods have been developed. These generally fall into two basic groups, moment methods \\citep{2001NewA....6..437G,2002ApJS..141..211C,2003ApJS..147..197H}, and ray-tracing methods \\citep{1998A&A...331..335M,1999MNRAS.309..287R, 1999ApJ...523...66A,2000MNRAS.314..611C,2001MNRAS.321..593N, 2001NewA....6..359S,2002ApJ...572..695R,2003A&A...405..189L, 2003MNRAS.345..379M,2004MNRAS.348..753S, 2004MNRAS.348L..43B,2005MNRAS...361..405I,methodpaper}. Several simulations of reionization have been performed using some of these codes, most often as a post-processing step  (i.e. the ``static limit'' which neglects the gasdynamical response to photoionization and heating) to cosmological N-body simulations with and without gas \\citep{2001MNRAS.321..593N,2001NewA....6..359S,2002ApJ...572..695R, 2003MNRAS.345..379M}, while others directly coupled the radiative transfer to the gas evolution and followed the evolution self-consistently \\citep{2001NewA....6..437G}. Despite these significant advances, all current reionization simulations are limited to fairly small volumes, with computational box sizes not exceeding $20\\,h^{-1}$ comoving Mpc, and usually much smaller than that. The main reason for this limitation is that the ionizing photon output during reionization is dominated by dwarf galaxies, which at early times are far more numerous than the larger galaxies, while the ionizing photon consumption (ionizations and recombinations) is dominated by even smaller structures, due to the hierarchical nature of $\\Lambda$CDM structure formation. The need to resolve such small galaxies imposes a severe limit on the computational box size. On the other hand, these ionizing sources were strongly clustered at high redshift and, as a consequence, the H~II regions they created are expected to quickly overlap and grow to very large sizes, reaching up to tens of Mpc \\citep[e.g.][]{2004ApJ...609..474B,2005MNRAS.363.1031F,2005astro.ph..7014C,2005pgqa.conf..369I}. The many orders of magnitude of difference between these scales demand extremely high resolution from any simulations designed to study early structure formation from the point of view of reionization. Further limitations are imposed by the low efficiency of the used radiative transfer methods. Most methods are based on ray-tracing and thus are quite accurate, but their computational expense grows roughly proportionally to the number of ionizing sources present. This generally renders them impractical when more than a few thousand ionizing sources are involved, severely limiting the computational volume that can be simulated effectively.   Over the years many analytical approaches to modelling reionization have been proposed \\citep[e.g.][]{1994ApJ...427...25S,2003ApJ...595....1H,2003ApJ...586..693W,2004ApJ...613....1F, 2005ApJ...624..491I}. However, these models are all statistically-averaged ones and can thus only make statistical predictions. Moreover, they have not been checked against simulations or observations and hence their level of reliability is currently unclear. In general, semi-analytical models are inevitably simplified in order to render the problem tractable and their prediction power depends on how well they can represent the relevant features of reionization. They could be quite useful in situations when simulations are too expensive, e.g. for exploring the parameter space, or studying very rare objects which requires huge volumes, well beyond the reach of current simulations. However, the correctness, reliability, and limitations of any semi-analytical model should still be established first by comparison with simulations.  \\begin{figure*} \\includegraphics[width=6.5in]{f1.ps} \\caption{Early structure formation in $\\Lambda$CDM, at $z=10$, from our N-body simulation: projection of the cloud-in-cell densities on the fine simulation grid ($3248\\times3248$ pixels) in a 20 comoving Mpc slice ($\\sim 6\\times10^8$ particles in the slice) of the $(100\\,h^{-1})^3$ Mpc$^3$ simulation volume. (See $\\tt http://www.cita.utoronto.ca/\\!\\sim\\!iliev/research.html$ for the full-resolution images and some movies of our simulations). \\label{simul_fig}} \\end{figure*}  The current development of novel, more efficient codes for both cosmological N-body and hydrodynamical simulations and for numerical radiative transfer finally allows reionization simulations in large volumes. In this work we present the first large-scale, in a volume of $(100\\,h^{-1}\\rm Mpc)^3$, radiative transfer simulations of this process. We achieve this by combining the results from an extremely large N-body simulation performed with the code PMFAST \\citep{2005NewA...10..393M} with a very fast and efficient cosmological radiative transfer code called $C^2$-Ray which we have recently developed \\citep{methodpaper}. Our underlying N-body simulation has a sufficient mass resolution to resolve all halos down to dwarf galaxies inside our volume reliably, as well as their clustering and the relevant large-scale fluctuations of the density field. Our new  ray-tracing radiative transfer method is based on explicit photon conservation in space and time which allows us to use large time steps and fairly coarse grids without loss of accuracy. Ionization fronts (I-fronts) are correctly tracked even for very optically-thick cells. These features make our code far more efficient than previous ones, allowing for faithful treatment (using parallel machines) of tens, even hundreds of thousands of ionizing sources on much larger grids than before.  We also note the very recent results of \\citet{2005ApJ...633..552K, 2005astro.ph.11627K} which simulate extremely  large cosmological volumes, up to $\\sim 1$ Gpc in size. However, these simulations have very coarse resolutions and do not resolve the individual ionizing sources. These are instead represented only in a mean, approximate way based on separate, much smaller scale radiative transfer simulations. Such approach may be appropriate for certain problems, like the one discussed in these papers, namely modelling the rare, bright high-redshift QSO's. However it cannot be used to answer the questions posed and addressed in the present work, due to its approximate nature and lack of both a source population resolution and a spatial one.  We assume that the gas is closely following the dark matter distribution, which is a quite accurate assumption at the large scales considered here \\citep{2004MNRAS.355..451Z}. Furthermore, the gas back reaction due to the ionization can be ignored to a good approximation at these scales since the global, large-scale I-fronts are highly supersonic, of weak R-type (from ``rarefied'', i.e. typically occurring in low-density gas), out-running any reaction of the gas \\citep{1987ApJ...321L.107S}. This approximation breaks down in dense gas inside halos, where the I-fronts slow down and transform to a D-type (from ``dense'', i.e. typically occurring in dense gas), generally preceded by a shock \\citep{2004MNRAS.348..753S,2005MNRAS...361..405I}. Thus, on smaller scales a fully-coupled hydrodynamic and radiative transfer treatment is required.  The general outlay of this paper is as follows. We describe our numerical methodology in \\S~\\ref{sim_sect}. In \\S~\\ref{results_sect} we present our results: on the globally-averaged quantities (e.g. ionized fraction, mean number of recombinations per atom, electron scattering optical depth) in \\S~\\ref{global_quantities_sect} and on the reionization geometry in \\S~\\ref{geometry_sect}. Finally, we discuss our results in \\S~\\ref{discuss_sect}. This paper concentrates on the geometry of reionization, the corresponding implications for the observation of the 21-cm signal will be presented in a companion paper \\citep{21cmreionpaper}.  Throughout this study we assume a flat $\\Lambda$CDM cosmology with parameters ($\\Omega_m,\\Omega_\\Lambda,\\Omega_b,h,\\sigma_8,n)=(0.27,0.73,0.044,0.7,0.9,1)$ \\citep{2003ApJS..148..175S}, where $\\Omega_m$, $\\Omega_\\Lambda$, and $\\Omega_b$ are the total matter, vacuum, and baryonic densities in units of the critical density, $h$ is the Hubble constant in units of 100 $\\rm km\\,s^{-1}Mpc^{-1}$, $\\sigma_8$ is the standard deviation of linear density fluctuations at present on the scale of $8 \\rm h^{-1}{\\rm Mpc}$, and $n$ is the index of the primordial power spectrum. We use the CMBfast transfer function \\citep{1996ApJ...469..437S}.  \\begin{figure*} \\begin{center} \\includegraphics[width=4in]{f2.eps} \\caption{Halo mass function from our simulation at various redshifts, as labeled, (thick, long-dashed) and analytical approximations: the standard Press-Schechter (PS; thin, solid) and Sheth-Tormen (ST; thick, dotted). \\label{PS_fig}} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{discuss_sect} We have presented the first large-scale radiative transfer simulation of the reionization of the universe. The total electron-scattering optical depth produced by our simulation agrees well with the results on CMB polarization from the first-year WMAP data, but the final overlap occurs at $z\\sim11$, somewhat too early to clearly agree with the current observations of high-z QSOs and galaxies, which point to reionization ending around $z=6-7$. However, we note that there are several effects which we do not include here which are expected to extend reionization without destroying the agreement with the WMAP results on the optical depth. These effects include small-scale (here sub-grid) gas clumping and lower ionizing efficiency of the sources, among others \\citep{2005ApJ...624..491I}. We are currently working on studying these effects in more detail with further simulations.  For a fixed distribution of sources the I-fronts escape from the high-density gas surrounding the sources and propagate faster into the lower-density gas in voids. This led to predictions in earlier work that reionization proceeded ``outside-in'', instead, with the preferential ionization of the lower-density regions \\citep[e.g.][]{2000ApJ...535..530G}. We have seen some such behaviour in the toy test runs we performed as a simple application in \\citet{methodpaper}. However, the full reionization simulation we have presented here points to the opposite conclusion. We demonstrated that in our simulation the process is ``inside-out'', i.e. with the high-density regions ionized earlier on average, and the large voids reionized last. The main reason for this is that the character of reionization is dictated by the interplay between H~II region expansion and evolving structure formation. The ionizing sources at high redshifts formed at the highest-density, rare peaks of the density field. As the cosmological structure formation progresses, more and more new sources form inside the density peaks, as these collapse. Thus, even though the I-fronts due to the earliest source clusters might escape into nearby voids and ionize them quickly, the numerous newly-formed sources ensure that on average there is always more mass ionized than volume. Earlier simulations which predicted ``outside-in'' reionization employed much smaller boxes, which resulted in faster reionization and less evolution in the source population in that time period.  This inside-out nature of reionization leads to an increased ionizing photon consumption since denser and clumpier gas has shorter recombination times, resulting in multiple recombinations per atom. In our particular simulation approximately 1.6 ionizing photons per atom were eventually required for completing the process.  Our conclusions are not affected by the relatively coarse resolution of the simulation presented here, with cell size $\\sim0.5\\,h^{-1}$~Mpc, which significantly filters the density fluctuations. We have now also run a simulation with the same underlying density field and sources, but with higher radiative transfer grid resolution ($406^3$ cells) \\citep{21cmreionpaper}, as well as multiple simulations with smaller box size ($35\\,h^{-1}$~Mpc), and thus higher spatial resolution \\citep{selfregulated}. Increased spatial resolution and the corresponding higher density contrasts emphasizes the inside-out nature of reionization even more and only strengthtens our current conclusions. In turn, such increased ionizing photon consumption would require fairly efficient production of ionizing photons at high-z, either due to massive, hot stars, high efficiency of star formation, high escape fractions, or a combination of all these.  Our current simulation does not include the effects of mechanical feedback (e.g. from supernovae) which can disrupt the smaller sources and modify their star formation rates. Currenly we also do not include the radiative input from ionizing sources with halo masses smaller than $2.5\\times10^9\\,M_\\odot$, which are below the resolution of our N-body simulation. Hence, our source population should be considered a low limit to the realistic one. These effects would be studied in a subsequent work \\citep{selfregulated}.  The scale of our simulation has allowed us for the first time to study numerically the large-scale geometry and statistics of reionization. We derived the correlation between the density of a region and its ionization state. We showed that the relation is complex and its mean and dispersion are significantly redshift- and scale-dependent. At late times and small scales the correlation essentially disappears. Furthermore, we demonstrated that the reionization history of sub-regions exhibits significant scatter at small scales and provides good description of the mean behaviour only at large scales, above 20-30 $h^{-1}$ Mpc.  The patchiness of reionization results in a significant increase of both the neutral- and ionized-gas density fluctuations, with important implications for statistical observations of reionization at, e.g., the 21-cm line of neutral hydrogen \\citep{21cmreionpaper} and small-scale CMB anisotropies. The power spectra of ionized and neutral density fluctuations at wavenumbers $k\\gtrsim1\\,\\rm h\\,Mpc^{-1}$ are boosted by up to factors of $\\sim2-3$ compared to the fluctuations of the total gas density. This roughly translates to a stronger fluctuation signal at arcminute and larger scales.  We derived the size distributions of the H~II regions in our simulations in both number counts and volume filling factors. We found two populations of bubbles, one with a large numbers of medium-sized ($\\sim10$ Mpc) and one with a few, rare and very large bubbles of size tens of Mpc. The first population provides most of the statistical fluctuations at arcminute scales discussed above, while the large ones should be the most prominent and most easily-detectable features of reionization. We also derived the distribution of Gunn-Peterson optical depths in our simulation volume as a first step to more detailed predictions for the observations of Ly-alpha emitters at high redshift by current and upcoming large surveys.  Finally, we demonstrated and for the first time quantified the non-Gaussian nature of the reionization statistics. The probability distribution functions for both the ionized mass fraction and the ionized mass are generally strongly non-Gaussian on all scales and at all times, especially at the beginning and end of reionization. At scales below the typical scales for the H~II regions the distribution even inverts, with the highest probability for a region to be either highly-neutral or highly-ionized. This is a feature of our reionization model, where no sources vary strongly in luminosity over time (or die). This is justified for our ionizing sources, which are relatively large galaxies above the ionized-gas Jeans mass, whose formation cannot, therefore, be easily suppressed. We will study the effects of the presence of smaller ionizing sources and more realistic source behaviour in a follow-up work."
    },
    "0512/gr-qc0512109_arXiv.txt": {
        "abstract": "\\noindent  We propose a class of actions for the spacetime metric that introduce corrections to the Einstein-Hilbert Lagrangian depending on the logarithm of some curvature scalars. We show that for some choices of these invariants the models are ghost free and modify Newtonian gravity below a characteristic acceleration scale given by $a_0 = c\\mu$, where $c$ is the speed of light and $\\mu$ is a parameter of the model that also determines the late-time Hubble constant: $H_0 \\sim \\mu$. In these models, besides the massless spin two graviton, there is a scalar excitation of the spacetime metric whose mass depends on the background curvature. This dependence is such that this scalar, although almost massless in vacuum, becomes massive and effectively decouples when one gets close to any source and we recover an acceptable weak field limit at short distances. There is also a  (classical) ``running'' of Newton's constant with the distance to the sources and gravity is easily enhanced at large distances by a large ratio. We comment on the possibility of building a model with a MOND-like Newtonian limit that could explain the rotation curves of galaxies without introducing Dark Matter using this kind of actions. We also explore briefly the characteristic gravitational phenomenology that these models imply: besides a long distance modification of gravity they also predict deviations from Newton's law at short distances. This short distance scale depends on the local background curvature of spacetime, and we find that for experiments on the Earth surface it is of order $\\sim 0.1mm$, while this distance would be bigger in space where the local curvature is significantly lower. ",
        "introduction": "The relative importance of the gravitational interaction increases as we consider larger scales, and it is at the largest scales that we can measure where the observed gravitational phenomena do not agree with our expectations. The Hubble constant measuring the rate of expansion of the Universe does not fall with time as predicted by General Relativity (GR) for a Universe that contains only known forms of matter, and the dynamics of galaxies seem to require much more matter than observed if explained in terms of GR. The most common approach to these problems is to assume the presence of unseen forms of energy that bring into agreement the observed phenomena with GR. The standard scenario to explain the dynamics of galaxies consists in the introduction of an extra weakly interacting massive particle, the so-called Cold Dark Matter (CDM), that clusters at the scales of galaxies and provides the required gravitational pull to hold them together. The explanation of the observed expansion of the universe requires however the introduction of a more exotic form of energy, not associated with any form of matter but associated with the existence of space-time itself: vacuum energy. And while CDM can be regarded as a natural possibility given our knowledge of elementary particle theory, the existence of a non-zero but very small vacuum energy remains an unsolved puzzle for our high-energy understanding of physics. However, the apparent naturalness of the CDM hypothesis finds also problems when one descends to the details of the observations. Increasingly precise simulations of galaxy formation and evolution, although relatively successful in broad terms, show well-known features that seem at odds with their real counterparts, the most prominent of which might be the ``cuspy core'' problem and the over-abundance of substructure seen in the simulations (see $e.g.$ \\cite{Ostriker:2003qj}). But, despite of this, the main problem that the CDM hypothesis faces is probably to explain the correlations of the relative abundances of dark and luminous matter that seem to hold in a very diverse set of astrophysical objects \\cite{McGaugh:2005er}. These correlations are exemplified in the Tully-Fisher law \\cite{Tully:1977fu} and can be interpreted as pointing to an underlying acceleration scale, below which the Newtonian potential changes and gravity becomes stronger. This is the basic idea of MOND (MOdified Newtonian Dynamics), a very successful phenomenological modification of Newton's potential proposed in 1983 \\cite{Milgrom:1983ca} whose predictions for the rotation curves of spiral galaxies have been realised with increasing accuracy as the quality of the data has improved \\cite{Sanders:2002pf}. Interestingly, the critical acceleration required by the data is of order $a_0 \\sim c H_0$ where $H_0$ is today's Hubble constant and $c$ the speed of light (that we will set to 1 from now on). The problem with this idea is that MOND is just a modification of Newton's potential so it remains silent in any situation in which relativistic effects are important. Efforts have been made to obtain MONDian phenomenology in a relativistic generally covariant theory by including other fields in the action with suitable couplings to the spacetime metric \\cite{Bekenstein:2004ne} (see also \\cite{Drummond:2001rj} for other approaches to galactic dynamics without Dark Matter). But these models do not address in a unified way the Dark Energy and Dark Matter problems, while a common origin is suggested by the observed coincidence between the critical acceleration scale and the Dark Energy density. In this paper we will propose a class of generally covariant actions, built only with the metric, that have the right properties to address these problems in a unified way, and where the relation $a_0 \\sim H_0$ finds a natural explanation. The theories we will consider modify gravity in the infrared, making it stronger below a characteristic acceleration scale, but this is not their only characteristic feature. When we are in a situation in which the dominant gravitational field is external, like in table-top experiments on Earth (that measure the gravitational field of some probes embedded in the dominant background gravitational field of the Earth), we can also expect short distance modifications of Newton's potential.  Regarding the long distance modifications, the source-dependent characteristic distance beyond which Newtonian gravity is modified in these theories, that we shall call $r_c$, is given by \\be \\f{G_NM}{r_c^2} = \\f{\\mu}{2}\\,, \\label{MOND} \\ee where $G_N$ is Newton's constant, $M$ the mass of the source and $\\mu$ is a parameter of the model that also determines the late-time Hubble constant: $H_0 \\sim \\mu$. This makes these models promising candidates to build a theory with a MOND-like Newtonian limit that could address the dynamics of galaxies without the need for Dark Matter. But when measuring the gravitational attraction between two probes in the external dominant gravitational field of a massive object of mass $M$, at a distance $r_d$ from its centre, we can also expect short distance modifications of Newton's law for distances smaller than \\be r_{SD}\\sim \\f{\\mu r_d^3}{G_N M}. \\ee If we plug in this expression the radius and mass of the Earth (with $\\mu \\sim H_0$), we get that for table top tests of Newton's law performed on the Earth surface $r_{SD} \\sim 10^{-2}cm$. This range is very interesting because it is the range currently being probed by experiments \\cite{Adelberger:2003zx}. But notice that the phenomenology of these theories is very different from the one expected from other theoretical considerations that also suggest a deviation from Newton's law at that scale motivated by the cosmological constant problem\\footnote{It is well known that naturalness arguments lead to the expectation that new physics associated with electro-weak symmetry breaking, besides the Higgs boson, should be seen in the LHC. Otherwise the electroweak scale becomes unstable under quantum corrections. Applying the same logic to the gravitational sector one would expect new gravitational phenomena to kick in at the vacuum energy scale that would cut-off the quantum divergences contributing to the vacuum energy. This hypothetical new physics should be seen in sub-mm measurements of Newton's potential \\cite{Beane:1997it}.} or the gauge hierarchy problem \\cite{Arkani-Hamed:1998rs}. The fact that this scale is the same in both cases is just a numerical coincidence. If we performed the same experiment on space, in the neighbourhood of the Earth's orbit for instance where the dominant gravitational field is that of the Sun, the relevant mass and distance we should use in the previous estimation is the Sun's mass and the Sun-Earth distance. In this case the ``short distance'' corrections would be expected in our theory at distances less than about $10^4 m$! This however does not mean that there should be big modifications to the motion of the planets or other celestial bodies. When the gravitational field we are measuring is that of the Sun, the corrections are suppressed at distances less than $r_c$ that is in this case of the order of $10^3 AU \\sim 10^{11}km$.  These characteristic experimental signatures arise in our theory because of the presence of an extra scalar excitation of the spacetime metric besides the massless spin two graviton. But while the graviton remains massless, the extra scalar has a mass that depends on the background curvature. This dependence is such that for the models that we will be interested on, those that modify gravity at large distances, this field becomes massive and effectively decouples when the background curvature is large, and in particular when we approach any source.  In the next section we will briefly review the results we obtained in \\cite{Navarro:2005gh,Navarro:2005da} studying models that involve inverse powers of the curvature in the action, giving some general expressions and discussing the generic features of the framework we will use for modifying gravity in the infrared. In particular we will focus on a class of models that had been proposed to address the acceleration of the Universe \\cite{Carroll:2004de} and also modify gravity at large distances \\cite{Navarro:2005gh,Navarro:2005da}. This discussion will enable us to motivate the class of actions that we will present in the third section, that depend on the logarithm of some curvature invariants. We will see that the theories that we propose in this section modify gravity at the MOND characteristic acceleration scale, and the gravitational interaction can easily become stronger at large distances. We will explore briefly the characteristic gravitational phenomenology expected in these models, discussing possible tests of these theories. In the fourth section we offer the conclusions. We will comment on further generalisations of the proposed actions and on the generic phenomenological features expected in the class of models that modify gravity at the MOND acceleration scale. We will also comment on the possibility of obtaining these theories as an effective action for the spacetime metric that takes into account strong renormalisation effects in the infrared that might appear in GR. ",
        "conclusions": "In this paper, motivated by the phenomenological success of MOND fitting the rotation curves of spiral galaxies without requiring Dark Matter, we have proposed a class of actions that modify gravity below the characteristic acceleration scale required by MOND: $a_0\\sim H_0$. There are two effects in these theories that are responsible for the infrared modification. First, there is an extra scalar excitation of the spacetime metric besides the massless graviton. The mass of this scalar field is of the order of the Hubble scale in vacuum, but its mass depends crucially on the background over which it propagates. This dependence is such that this excitation becomes more massive as we approach any source, and the extra degree of freedom decouples at short distances in the spacetime of a spherically symmetric mass. This feature makes this excitation to behave in a way that reminds of the chameleon field of \\cite{Khoury:2003aq}, but in our case this ``chameleon'' field is just a component of the spacetime metric coupled to the curvature. But there is a second effect in these theories: the Planck mass that controls the coupling strength of the massless graviton also undergoes a rescaling or ``running'' with the distance to the sources (or the background curvature). This phenomenon, although a purely classical one in our theory, is reminiscent of the quantum renormalisation group running of couplings. So one might wonder if actions of the type (\\ref{actionMOND}) could be an effective classical description of strong renormalisation effects in the infrared that might appear in GR (see $e.g.$ \\cite{Reuter:1996cp} and references therein), as happens in QCD. In fact, corrections depending on the logarithm of the renormalisation scale are ubiquitous in quantum field theory, and it appears natural to identify the renormalisation scale with a function of the curvature if we want to build an effective classical action for the spacetime metric that takes into account these quantum effects. Indeed, we have seen that these models offer a phenomenology that seems well suited to describe an infrared strongly coupled phase of gravity: at high energies/curvatures we can use the GR action or its linearisation as a good approximation, but when going to low energies/curvatures we find a non-perturbative regime. At even lower energies/curvatures perturbation theory is again applicable, but the relevant theory is of scalar-tensor type in a de Sitter space.  We would like also to emphasise that there are clearly many modifications of the proposed class of actions that would offer a similar phenomenology, such that gravity would be modified below a characteristic acceleration scale of the order of the one required in MOND. For instance if we consider the action \\be S=\\int \\!\\!d^4x\\sqrt{-g}\\frac{1}{16\\pi G_N}\\left\\{R-\\mu^{2}\\left({\\rm Log}[f]\\right)^n\\right\\}\\,, \\ee with the same assumptions on $f$ that we did before we get that now the critical acceleration $a_0$ also has a ``running'' with the mass as \\be a_0=\\f{G_NM}{r_c^2} \\sim \\mu\\left({\\rm Log}\\left[\\f{48a_0^3}{Q_0G_NM}\\right]\\right)^{(n-1)/2}. \\ee To introduce some kind of scale dependence of $a_0$ could be interesting since MOND typically gives an overestimation of the amount of visible matter at cluster scales.  Since the expansions we have used break down for some intermediate range of energies/distances one should still show that the dynamics in this ``non-perturbative'' regime are consistent for some choice of $f$ and that one recovers an acceptable matching between the high energy/short distance and low energy/large distance regimes. And then one should compute the predictions of these theories in many different situations for which there are experimental data to compare with before any of these models could be considered a viable alternative to a $\\Lambda$CDM cosmology. This is a non-trivial task but it is worth undertaking it because we have seen that there are reasons to believe that one might explain many aspects of the cosmological and astrophysical observations without introducing Dark Matter in this class of theories. And, as we have seen, these theories also offer the unique possibility of being tested not only through astrophysical observations, but also through well-controlled laboratory experiments where the outcome of such experiments is correlated with parameters that can be determined by means of cosmological and astrophysical measurements."
    },
    "0512/astro-ph0512522_arXiv.txt": {
        "abstract": "We report results of multiband optical photometric monitoring of two well known blazars, S5 0716+714 and BL Lacertae, carried out during 1996 and 2000$-$01 with an aim to study optical variations on time scales ranging from minutes to hours and longer. The light curves were derived relative to comparison stars present on the CCD frames. Night to night intensity variations of $\\ge$ 0.1 mag were observed in S5 0716+714 during a campaign of about 2 weeks in 1996. A good correlation between the lightcurves in different optical bands was found for both inter-night and intra-night observations. In all, two prominent events of intra-night optical variability (INOV) were detected in S5 0716+714. Each of these rapidly varying segments of the lightcurves can be well fitted with an exponential intensity profile whose rate of variation is essentially the same in both cases. Our long-term monitoring data of S5 0716+714 showed a distinct flare around JD 2451875 which can be identified in the {\\it BVRI} bands. This flare coincides with the brightest phase recorded during 1994$-$2001 in the long-term lightcurves reported by Raiteri et al.\\ (2003). No evidence for the object to become bluer when brighter was noticed on either inter-night or intra-night time scales. On the other hand, our essentially simultaneous multiband optical observations of BL Lacertae in October 2001 showed flux variations that were not achromatic. This blazar definitely was found to become bluer when brighter on intra-night time scales and there is a less significant trend of the same type on inter-night time scales. Based on five nights of observations during a week, BL Lacertae showed a peak night-to-night variability of $\\sim$ 0.6 mag in {\\it B}. Thus, we found that the present optical observations of the two prominent blazars, made with similarly high sensitivity, reveal a contrasting behaviour in terms of the dependence of spectral hardening with increasing brightness, at least on intra-night, and possibly also on inter-night, time scales. ",
        "introduction": "Active galactic nuclei (AGNs) are known to be variable on different time scales across the electromagnetic spectrum. Photometric studies of intensity variations in AGNs provide a uniquely powerful tool for investigating the processes occurring in the vicinity of their central engine. In particular, studies of very rapid intensity variations, or intra-night optical variability (INOV), where the variations have amplitudes of a few hundredths of a magnitude on hour-like or shorter time scales, enable us to probe their innermost nuclear cores, on the scales of microarcseconds. Blazars define a class of AGNs  made up of optically violent variable quasars, high polarization quasars and BL Lac objects. They show high polarization and violent variability at optical wavelengths, and in the radio band contain compact flat spectrum sources which often exhibit apparent superluminal motion and a high polarization level. These characteristics are generally attributed to synchrotron emission from a relativistic jet with the jet axis oriented at small angles to the observer's line of sight.  Theoretical explanations for the origin of INOV  can be broadly divided into intrinsic and extrinsic categories. One such extrinsic mechanism is refractive interstellar scintillation (e.g., Kedziora-Chudczer et al.\\ 1997), though this is relevant only in the radio band. Another proposed extrinsic mechanism is superluminal microlensing (Gopal-Krishna \\& Subramanian 1991), which can explain the rapidity of variations, though it is unlikely to apply to a large fraction of blazars. The dominant basic model for intrinsic variability invokes shocks propagating through the jet (e.g., Blandford \\& K\\\"{o}nigl 1979; Marscher \\& Gear 1985).  Models designed to explain INOV based on jets involve turbulence behind the jet (Marscher, Gear \\& Travis 1992), light echo effects (Qian et al. 1991), helical filaments (Rosen 1990) or changing directions of the shocks in the jet (Gopal-Krishna \\& Wiita 1992; Nesic et al.\\ 2005). The effectiveness of all these models are enhanced by relativistic effects in the jet, especially when the viewing angle is small. Another family of intrinsic explanations invokes numerous flares or hot spots on the surface of (or in the corona above) the accretion disk (e.g., Mangalam \\& Wiita 1993). Similar models have been proposed to explain the X-ray variations of AGN (e.g., Abramowicz et al.\\ 1992). While such disk perturbations are unlikely to dominate blazar variability they may propagate into or otherwise affect jet dominated emission (e.g., Wiita 2005).   \\subsection{S5 0716$+$714} S5 0716$+$714, classified as a BL Lac object, was discovered in the Bonn-NRAO Radio Survey (K\\\"{u}hr et al.\\ 1981) of flat spectrum radio sources with a 5 GHz flux greater than 1 Jy. It is believed to be at a redshift z $>$ 0.3 (Wagner et al. 1996) due to the lack of the detection of its host galaxy; however, a lower value of z $\\sim$ 0.1 has also been suggested recently (Kadler et al. 2004). This source, an intra-day variable at radio and optical wavelengths as well as a $\\gamma$-ray emitter (Wagner et al.\\ 1996), is a favourite target for variability studies. In a 4 week long monitoring campaign in February 1990, this blazar showed an abrupt transition in its variability pattern from a higher flux level with roughly 1-day time scale during the first week, to a lower flux level with approximately 7-day time scale for the remainder of the campaign, both at optical and centimeter wavelengths (Quirrenbach et al.\\ 1991). Another interesting result is that the radio spectral index of this blazar was found to correlate with intra-night optical variations such that a flattening of the radio spectrum near the optical maxima was found for several quasi-cycles (Qian et al.\\ 1996; Wagner et al.\\ 1996). The intriguing evidence for correlated variability in the radio and optical ranges (Quirrenbach et al. 1991; Wagner et al. 1996) certainly points to an intrinsic origin of the variations in this blazar. These authors further inferred that the {\\it R} band light curves evinced a significant signature of flickering on time scales as short as 15 minutes,  much faster than the apparent quasi-periodic time scales.  Raiteri et al. (2003) present results on multi-band optical observations ranging from 1994$-$2001 and also present radio observations spanning more than 20 years. They report that the long term optical brightness variations of this source appears to have a characteristic timescale of $\\sim$3.3 yrs.  Ghisellini et al.\\ (1997) and Sagar et al.\\ (1999) too have reported results of monitoring of this blazar in BVRI colours during February to April 1994. Their campaigns recorded a few events of INOV with an amplitude of about 5\\% within a few hours, superposed on slower variations.  In order to further investigate the INOV behaviour of this blazar, we carried out a two weeks long optical monitoring campaign during April 1996. These observations covered a total baseline of 16 days and provided a fairly dense temporal coverage on most of the nights. These data, in conjunction with our additional observations during 2000$-$01, are used here to investigate both short-term and long-term variability of this prominent blazar.   \\subsection{BL Lacertae} BL Lac (2200$+$420), the prototype of this class of AGNs, was identified by Schmitt (1968) with the radio source VRO 42.22.01, and its spectrum was found to be featureless by Viswanathan (1969). However, a broad $H{\\alpha}$ line with EW $>$ 5 \\AA ~was later detected by Vermeulen et al.\\ (1995). It is found to be embedded in an elliptical host galaxy (Wurtz, Stocke \\& Yee 1996) with $z = 0.069$. Though large amplitude variations on time scales ranging from days to decades (Webb et al.\\ 1988) and rapid variations of 0.1 mag/hr (Racine 1970) were already known, the advent of CCDs used as N-star photometers conclusively demonstrated INOV in BL Lac (Miller et al.\\ 1989). These authors report rapid changes in the {\\it V} band flux on time scales as short as 1.5 hrs. It brightens almost every year and was reported to be in the active state in 2001 by Mattei et al.\\ (2001). We carried out intra-night BVR monitoring of this blazar for 5 nights in October 2001.  \\subsection{Spectral variability in blazars} Although optical variability on hour-like time scales is now a well established phenomenon for blazars, its relationship to the long-term variability remains unclear. Possible clues could come from monitoring the optical spectrum for correlations with brightness. S5 0716+714 has been intensively studied for variability across the electromagnetic spectrum on various time scales (see Wu et al. 2005 and references therein; Nesci et al. 2005). On inter-night time scales a bluer when brighter correlation was found when the object was in an active or flaring state, but this trend was absent during its relatively quiescent state (Wu et al. 2005). Also, during one night's monitoring, Wu et al. (2005) noticed no spectral hardening with brightness.  During a 11-day long intra-night monitoring campaign of BL Lac during a major outburst in July 1997, clear evidence for the object to become bluer when brighter was noted by Clements \\& Carini (2001). However, they interpreted this as an artefact of an increased contribution from the host galaxy when the blazar was fainter. In contrast, based on their BRI photometry of BL Lacertae on 5 nights during 1999 and 2001, Papadakis et al.\\ (2003) argued that spectral hardening with increasing brightness is evident even after the contribution from the host galaxy is subtracted, hinting that the effect is intrinsic to the blazar. The same significant trend was independently reported by Villata et al.\\ (2004) from the intra-night monitoring of BL Lacertae in UBVRI bands. They also noted that this colour--brightness correlation is much weaker for the long-term brightness variations, which were nearly achromatic on a few day-like time scale. The results of long-term BVRI monitoring between 1995 and 1999 of 8 BL Lac objects, including BL Lacertae and S5 0716$+$714, by D'Amicis et al.\\ (2002) were further analysed by Vagnetti et al.\\ (2003) who found a tendency for bluer colour at higher luminosity for all of them; however, for their dataset, BL Lac showed a fairly strong correlation while 0716$+$714 showed a weaker one.  \\begin{figure*} \\psfig{file=fig1.eps,width=19cm,height=16cm}  % \\caption{DLCs for S5 0716+714 and the two comparison stars (Sect. 3). Filled and open circles denote the differential magnitude in {\\it R} and V bands, respectively. The {\\it V} band DLCs have been shifted by $+$0.4 mag (between blazar and comparison stars) and $-$0.25 mag (between comparison stars themselves), respectively. The time scale for all DLCs is indicated by the horizontal bar, with actual times given in Table 1.} \\end{figure*}  If the trend of spectral hardening with increasing brightness is confirmed to be (nearly) universal, at least for the variations on hour-like time scales, one interesting explanation posits short-term fluctuations of only the electron injection spectral index (e.g., Mastichiadis \\& Kirk 2002; B\\\"{o}ttcher \\& Reimer 2004). However, some blazars are found to show anomalous spectral behaviour (see Ram\\'{i}rez et al.\\ 2004 and references therein). For example, PKS 0736$+$017 showed a tendency for its spectrum to become redder when brighter (Ram\\'{i}rez et al.\\ 2004), both on inter-night and intra-night time scales. The multiband optical observations of the prominent blazars S5 0716+714 and BL Lacertae reported here were carried out to shed more light on the possible dependence of optical spectrum on the brightness state of blazars. In the following section we describe the observations and data reduction. Sect.\\ 3 presents the analysis of the lightcurves, and the results are summarized in Sect.\\ 4.   \\hspace*{-4.0cm} ",
        "conclusions": "Based on the  published long-term light curves of both these blazars (Raiteri et al.  2003; Villata et al.\\ 2004; Nesci et al.\\ 2005), the present INOV observations are seen to coincide with their relatively faint optical states. Nonetheless, two clear events of INOV, one of brightening, and the other of fading, were observed in S5 0716$+$714. The events are well fitted with an exponential intensity profile with essentially the same rate of variation in the {\\it V} and {\\it R} bands. We note a hint of a possible  temporal lag of $\\sim$10 minutes between these two bands with the shorter wavelength leading the longer, as expected in the standard shock-in-jet model; such a lag is not expected in most disk-based variability models (e.g., Wiita 2005). In the present eight nights of our monitoring of S5 0716$+$714 spanning about a fortnight, two noticeable events of inter-night variability (one around JD 2450180 and the other around JD 2450188) of amplitude $\\ge$ 0.2 mag were found. Our long-term monitoring of this blazar, on the other hand, coincided with the maximum brightness phase seen in the lightcurve covering 1994$-$2001, as reported by Raiteri et al. (2003). By combing their data with the present measurements a large 1.5 magnitude flare is detected around JD 2451875 lasting for about 75 days. This is followed by a weaker variation on a similar timescale. No clear evidence of colour variation with brightness was found in either our inter-night or our intra-night monitoring of S5 0716+714.  The other blazar, BL Lacertae, was found to vary strongly in all three bands (B,V,R), on one of the five nights we monitored it and in two of the three bands on two other nights. Also, on the night when strong INOV was observed in all three bands, a temporal lag between the {\\it V} and {\\it R} band lightcurves may be present, with the {\\it V} band variations apparently leading the {\\it R} band variations by a few minutes. Nesci et al. (1998) also noted that no time lag between different bands could be unambiguously detected during intra-night monitoring observations of BL Lac in 1997.  A clear trend was seen in our data for the source to become bluer when brighter on intra-night time scales. Such a clear trend was not evident on inter-night time scales, although the data suggest one may be present. The presence (or absence) of this consistency between the colour-magnitude behaviour  on intra-night and inter-night time scales can provide interesting clues about the origin of blazar variability from hour-like to much longer time sales. In view of this, sustained efforts for multi-colour monitoring of blazars on different time scales will be fruitful.  \\begin{figure} \\centering \\includegraphics{fig8.eps} \\caption{Inter-night CMD for BL Lacertae based on our observations in 2001. The solid line is the linear least-squares fit to the data points. The dotted line is the linear least-squares fit to the data excluding the first data point at V $\\sim$14.9 mag.} \\label{Colour} \\end{figure}   \\begin{figure} \\centering \\includegraphics{fig9.eps} \\caption{Intra-night CMD for BL Lacertae on 22 October 2001. The solid line is the linear least-squares fit to the data points} \\end{figure}   \\begin{figure} \\hspace*{-0.9cm}\\psfig{file=fig10.eps} \\caption{Cross correlation function for the {\\it V} and {\\it R} band lightcurves of BL Lacertae on 22 October 2001. Full circles with error bars refer to the DCF, whereas the solid line is calculated from the ICCF (see Sect.\\ 3.1.1).} \\end{figure}"
    },
    "0512/astro-ph0512008_arXiv.txt": {
        "abstract": "\\vspace{-0.6cm} The Great Observatories Origins Deep Survey (GOODS) combines deep HST and Spitzer imaging with the deepest Chandra/XMM observations to probe obscured AGN at higher redshifts than previous multiwavelength surveys. We present a self-consistent implementation of the AGN unification paradigm, which postulates obscured AGN wherever there are unobscured AGN, to successfully explain the infrared, optical, and X-ray number counts of X-ray sources detected in the GOODS fields. Assuming either a constant ratio of obscured to unobscured AGN of 3:1 (the local value), or a ratio that decreases with luminosity, and including Compton-thick sources, we can explain the spectral shape and normalization of the extragalactic X-ray ``background'' as a superposition of unresolved AGN, predominantly at $z$$\\sim$0.5-1.5 and L$_x$$\\sim$10$^{43}$-10$^{44}$ ergs/s. The possible dependence of the obscured to unobscured ratio with redshift is not well constrained; present data allow it to decrease or increase substantially beyond $z$$\\sim$1. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/hep-ph0512367_arXiv.txt": {
        "abstract": "We discuss the possibility of generation of baryon inhomogeneities in a quark-gluon plasma phase due to moving $Z(3)$ interfaces. By modeling the dependence of effective mass of the quarks on the Polyakov loop order parameter, we study the reflection of quarks from collapsing $Z(3)$ interfaces and estimate resulting baryon inhomogeneities in the context of the early universe. We argue that in the context of certain low energy scale inflationary models, it is possible that large $Z(3)$ walls arise at the end of the reheating stage. Collapse of such walls could lead to baryon inhomogeneities which may be separated by large distances near the QCD scale. Importantly, the generation of these inhomogeneities is insensitive to the order, or even the existence, of the quark-hadron phase transition. We also briefly discuss the possibility of formation of quark nuggets in this model, as well as baryon inhomogeneity generation in relativistic heavy-ion collisions. ",
        "introduction": "Generation of baryon inhomogeneities in the early universe can have important implications for nucleosynthesis, and for the possibility of creating compact baryon rich objects \\cite{wtn}. Though, current observations do not support any strong deviation from the standard big-bang nucleosynthesis  calculations. Calculations of inhomogeneous big bang nucleosynthesis resulting from an inhomogeneous distribution of baryons in the universe, (such as those in ref. \\cite{ibbn1,ibbn2}), therefore, can be used to constrain the baryon inhomogeneities present in the early universe.  There have been numerous investigations of the nature of baryon inhomogeneities generated during a first order quark-hadron phase transition \\cite{wtn,fuller}. In these investigations, baryon inhomogeneities arise due to moving bubble walls at the transition, with baryons getting concentrated in the remaining localized quark-gluon plasma (QGP) regions.  Main problems in implementing the scenario of ref. \\cite{wtn} have been regarding the nature of the quark-hadron phase transition as well as the relevant length scales. Lattice calculations \\cite{ltc} tell us that for realistic values of quark masses, quark-hadron transition is at best a weak first order transition, and most likely it is a cross-over. The scenario of ref. \\cite{wtn} does not work in this case. Even if one allows for a possibility of strong first order transition, relevant length and time scales are such that the resulting baryon inhomogeneities are separated by very small distances. Typical separation between such baryonic lumps is of the order of separation between the nucleation sites of the hadronic bubbles, which is at most of the order of few cm at the end of the quark-hadron transition for homogeneous nucleation\\cite{fuller,impur}. In order that these baryonic lumps survive various dissipative processes, this separation needs to be at least of order of a meter at the transition stage \\cite{bfluct}. There have been discussions of larger separations between baryon inhomogeneities invoking impurity induced inhomogeneous bubble nucleation \\cite{impur}, presence of density fluctuations \\cite{fluc,layek} etc. However, all these scenarios still depend crucially on the assumption of a first order phase transition, and will not work if the quark-hadron transition was a cross-over.  In this paper we propose a different scenario where baryon inhomogeneities are produced not due to moving quark-hadron phase boundaries, but due to moving $Z(N)$ interfaces. $Z(N)$ interfaces arise when one uses the expectation value of the Polyakov loop, $l(x)$, as the order parameter for the confinement-deconfinement phase transition of an SU(N) gauge theory \\cite{plkv}.  This order parameter transforms non-trivially under the center $Z(N)$ of the  SU(N) group  and is non-zero above the critical temperature $T_c$. This breaks the global $Z(N)$ symmetry spontaneously above $T_c$, while the symmetry is restored below $T_c$ in the confining phase where this order parameter vanishes. For QCD with SU(3) color group, spontaneous breaking of the discrete $Z(3)$ symmetry in the QGP phase leads to the existence of domain walls (interfaces) across which $l(x)$ interpolates between different $Z(3)$ vacua.  The properties and physical consequences of these $Z(3)$ interfaces have been discussed in the literature \\cite{zn}. Though, we mention that it has also been suggested  that these interfaces should not be taken as physical objects in the Minkowski space \\cite{smlg}. Similarly, it has also been subject of discussion whether it makes sense to talk about this $Z(3)$ symmetry in the presence of quarks \\cite{qurk1}. The presence of quarks can be interpreted as leading to explicit breaking of $Z(3)$ symmetry, lifting the degeneracy of different $Z(3)$ vacua \\cite{qurk2,psrsk,psrsk2,veff}. In this approach, with quarks, $Z(3)$ interfaces become unstable and move away from the region with the unique true vacuum. Thus, in the context of cosmology, if these walls were produced at some early stage (say after GUT scale inflation), it is likely that they will quickly disappear due to this pressure difference between different $Z(3)$ vacua. However, we will argue (in section III) that in the context of certain low energy scale inflationary models it is possible that large $Z(3)$ domain walls may arise in the QGP phase near the quark-hadron transition stage and may lead to observational effects.  The basic idea of our model is that as $l(x)$ is the order parameter for the quark-hadron transition, physical properties such as effective mass of the quarks should be determined in terms of $l(x)$.  This also looks natural from the expected correlation between the chiral condensate and the Polyakov loop.  Thus, if there is spatial variation in the value of $l(x)$ in the QGP phase then effective mass of the quark traversing that region should also vary. For regions where $l(x) = 0$, quarks should acquire constituent mass as appropriate for the confining phase. As we will see below, $l(x)$ varies across a $Z(3)$ interface, acquiring small magnitude in the center of the wall. A quark passing through this interface, therefore, experiences a nonzero potential barrier leading to non-zero reflection coefficient for the quark. Due to this, as a closed domain wall collapses, quarks inside will stream through it.  With a non-zero reflection coefficient, net baryon number density inside will grow, somewhat  in the manner as in the conventional treatments of collapsing quark-hadron phase boundaries. This will lead to formation of baryonic lumps.  Important thing to realize is that all this happens in the QGP phase itself, with any possible quark-hadron transition being completely irrelevant to this discussion. The only relevance of the quark-hadron transition is that in the hadronic phase $l(x) = 0$ so all $Z(3)$ domain walls disappear. The final structure of the baryon inhomogeneities will therefore be decided by those $Z(3)$ interfaces which are last to collapse. As mentioned above, we will argue in section III that it is possible that the size and separation of different collapsing domain walls may be of the order of a fraction of the horizon size just above the quark-hadron transition stage, i.e. of order of a km. If such large domain walls could form then the number of baryons trapped inside can be very large. Also, due to larger mass of the strange quark, reflection coefficient for them is larger than that for the u and d quarks. This leads naturally to strangeness rich quark nugget formation which, as we will show, can have baryon number as large as about 10$^{44}$ within a size of order 1 meter.  In a previous paper we have shown that at the intersection of the three different $Z(3)$ interfaces $l(x)$ vanish due to topological considerations, leading to a topological string whose core is in the confining phase \\cite{znstr}. Structure of this string is similar to the standard axionic string which forms at the junction of axionic domain walls \\cite{axion}. With quarks contributing to explicit $Z(3)$ symmetry breaking, this will lead to decay of $Z(3)$ interfaces along with decay of the associated strings. As $l(x) = 0$ in the core of these strings, collapsing string loops will have larger reflection coefficients for quarks and will also contribute to formation of baryon inhomogeneities. However, unless this string network is very dense, large scale baryon inhomogeneities will mostly result from collapsing $Z(3)$ interfaces.  The mechanism discussed in this paper will also lead to generation of baryon fluctuations in the QGP formed in relativistic heavy-ion collision experiments, with the walls forming during the initial thermalization stage. The effects of explicit symmetry breaking due to quarks on the evolution of wall etc., as mentioned above, will not be much relevant there because of very short time scale available for the evolution of QGP. We plan to study this using detailed computer simulations in a future work.  The paper is organized in the following manner. In section II we discuss structure of $Z(N)$ walls and give numerical results for the profile of $Z(3)$ walls for the case of QCD. Section III discusses how $Z(3)$ walls can form in the early universe. In section IV baryon inhomogeneity generation due to quark reflection from collapsing $Z(3)$ walls is estimated. Numerical results and discussion are given in section V. ",
        "conclusions": "In Fig.2 we have given plots of $n_i$ vs. time (in microseconds) and of $\\rho$  vs. R (in meters) for T = 200 MeV and for the choice of $m_0$ = 300 MeV in Eq.(5).  (Again, initial values of $n_i$, $n_o$ are normalized to the average baryon density of the universe $n_{av}$. To get absolute values of these densities, and of $\\rho$, one should multiply by $n_{av}$.) We have taken the number of domain walls in a horizon volume $N_d$ to be 10. We find that the size of the region inside which the baryon overdensity $\\rho > 1000$ is about 10 m for u,d quarks while the size is about 60 m for the strange quark case. Baryon density sharply rises for small $R$.  We see that for $R < 1 m$, $\\rho$ rises to a value of about 20,000 for u,d quarks and to a value of about 6 $\\times 10^5$ for the strange quark. These overdensity magnitudes and sizes are large enough that they can survive until the time of nucleosynthesis and affect nuclear abundances. Typical separation between the inhomogeneities is the inter-domain wall separation near the QCD scale (below which walls disappear), and hence can be very large in our model, of order of a km. (Of course, with the assumption that large size walls arise at the end of reheating stage in a low scale inflationary model, as discussed in section III.) This corresponds to about 100- 200 km length scale at the nucleosynthesis epoch, which is  precisely the range of length scales which can have optimum effects on nucleosynthesis calculations  in ref.  \\cite{ibbn2}.  \\begin{figure} \\epsfig{file=fig2a.eps,height=65mm} \\epsfig{file=fig2b.eps,height=53mm} \\caption{Plots of $n_i$ vs. time $t$ (in microseconds), and $\\rho$  vs. R (in meters). The origin for $t$ is chosen at the beginning of the wall collapse.  Solid curves are for s quark and dashed curves for u,d quarks.} \\label{fig:fig2} \\end{figure}  When we consider only one domain wall in the horizon ($N_d$ = 1) then overdensities are larger. For example for the above cases, we find that within $R < 1$ m, $\\rho$ is larger by a factor of 2 to 4. Overdensities become much smaller for the u,d quark case if we take $m_0$ in Eq.(5) to be equal to $m_q$ (instead of 300 MeV), as the potential barrier becomes much smaller than the typical quark energy leading to very small reflection coefficient. For example, for other parameters of Fig.2, $\\rho$ is about 20 for $R < 1$ m for u,d quark. However, for the strange quark even with $m_0 = m_q$ the potential barrier is high enough with significant reflection of quarks and leads to $\\rho = 120000$ for $R < 1$ m. For comparison we have also calculated overdensities occurring at T = 300 MeV. These are much smaller, first due to smaller domain wall width, and secondly due to larger value of $l$ in the domain wall (see, Fig.1), leading to smaller potential barrier (height as well as width). For example, with $m_0$ = 300 MeV, within $R < 1$ m we get $\\rho = 5000$ for $s$ quark, and $\\rho = 400$ for u,d quarks.  With large overdensities occurring as in Fig.2, there may be possibility of quark nugget formation \\cite{ngt}. Indeed we find that for certain cases, e.g. with the parameters of Fig.2, total number of baryons can be very large, $\\sim 10^{44}$ within $R = 1$ m. These regions will be dominated by strange quarks as is clear from Fig.2. These seem to be favorable conditions for the formation of stable quark nuggets. If these survive cooling down through $T_c$, and survive until present then they may constitute dark matter, without affecting microwave background anisotropy or nucleosynthesis constraints.  We summarize main features of our model. We have discussed formation and evolution of $Z(3)$ domain walls in the early universe. We have argued that, in the context of low scale inflationary models with reheat temperature of order of few GeV, it is possible that large $Z(3)$ walls can arise near the QCD scale. (We also briefly mentioned other possibilities where large $Z(3)$ walls can arise in inflationary models based on thermal inflation, or warm inflation etc.) We study baryon inhomogeneities resulting from these walls.  In our model, baryon inhomogeneities are produced not due to moving quark-hadron phase boundaries as in the conventional treatments, but due to moving $Z(3)$ domain walls. The variation in the value of the Polyakov loop order parameter across the wall leads to non-zero reflection coefficient for the quarks. As a closed domain wall collapses, a fraction of quarks inside it remains trapped leading to production of baryon inhomogeneities. Important thing is that all this happens in the QGP phase itself, with any possible quark-hadron transition being completely irrelevant.  We have assumed that near $T_c$, the pressure difference between the metastable $Z(3)$ vacua and the true vacuum may be small so that surface tension may play a dominant role in the early evolution of domain walls, which form as the temperature of the universe crosses $T_c$ during reheating stage at the end of inflation.  The separation of the resulting inhomogeneities is then the separation between different collapsing  domain walls, which may be of the order of a fraction of the horizon size near the quark-hadron transition stage. Resulting overdensities then have large enough magnitudes and sizes that they can survive until the stage of nucleosynthesis and affect the abundances of elements. We also find that if such large walls can form then strangeness rich quark nuggets of large  baryon number ($~ 10^{44}$) can form in our model.  If the effects of pressure difference do not remain subdominant near $T_c$ in the coarsening dynamics due to surface tension of walls, then resulting walls will not be as large. Still resulting baryon inhomogeneities may have large enough scales to survive until nucleosynthesis and affect abundances of elements. In view of tight constraints on models of inhomogeneous nucleosynthesis, our results can be used to constrain various models of low scale inflation (or other inflationary models, as discussed above).  The mechanism discussed in this paper will also lead to generation of baryon fluctuations in the QGP formed in relativistic heavy-ion collision experiments, with the walls forming during the initial thermalization stage. The effects of explicit symmetry breaking due to quarks, as discussed above, will not be much relevant there because of very short time scale available for the evolution of QGP. However,  one cannot use simplifying assumptions about coarsening of $Z(3)$ walls for the heavy-ion case, as one can do for the case of the universe. Similarly, because of rapid cooling due to expansion, one will have to use time dependent potential barrier for estimating quark reflection from $Z(3)$ walls.  We plan to study this using detailed computer simulations in a future work."
    },
    "0512/astro-ph0512536_arXiv.txt": {
        "abstract": "In this work, the relation between muon production in extensive air showers and features of hadronic multiparticle production at low energies is studied. Using CORSIKA, we determine typical energies and phase space regions of secondary particles which are important for muon production in extensive air showers and confront the results with existing fixed target measurements. Furthermore possibilities to measure relevant quantities of hadron production in existing and planned accelerator experiments are discussed. ",
        "introduction": "The energy spectrum and composition of the cosmic rays with energies above $10^{15}$\\,eV are typically derived from measurements of the number of electrons and muons produced in extensive air showers (EAS) at ground. However, the results of such a shower analysis are strongly dependent on the hadronic interaction models used for simulating reference showers \\cite{kascade_holger}.  Therefore it is important to study in detail the role of hadronic interactions and in particular the energy and secondary particle phase space regions that are most important for the observed characteristics of EAS.  The electromagnetic component of a shower is well determined by the depth of maximum and the energy of the shower. Due to the electromagnetic cascade, having a short radiation length of \\unit[$\\sim 36$]{g/cm$^2$}, any information on the initial distribution of photons produced in $\\pi^0$ decays is lost. Therefore the electromagnetic shower component depends on the primary particle type only through the depth of shower maximum. In contrast, the muon component is very sensitive to the characteristics of hadronic interactions.  Once the hadronic shower particles have reached an energy at which charged pions and kaons decay, they produce muons which decouple from the shower cascade. The muons propagate to the detector with small energy loss and deflection and hence carry information on hadronic interactions in EAS. Due to the competition between interaction and decay, most of the muons are decay products of mesons that are produced in low-energy interactions. Therefore it is not surprising that muons in EAS are particularly sensitive to hadronic multiparticle production at low energy \\cite{EngelISMD1999, Drescher2003}. Recent model studies show that even at ultra-high shower energies the predictions on the lateral distribution of shower particles depend strongly on the applied low-energy interaction model \\cite{Drescher2004, Heck2003}. ",
        "conclusions": "Due to the interplay between decay and interaction of pions and kaons, low energy hadronic interactions are very important for muon production in extensive air showers. With increasing lateral distance the mean energy of these interactions, which are mainly initiated by pions and nucleons, decreases.  The most important interaction energies and phase space regions are accessible for fixed target experiments. However, so far only measurements with protons as projectiles and with a very limited secondary particle phase space exist. The situation could be improved considerably, if data from fixed target experiments with large acceptance detectors such as HARP, NA49 and MIPP would be analysed for minimum bias collisions especially for $\\pi$+C reactions.  Fixed target measurements can contribute to improving low energy interaction models. Due to the sensitivity of EAS to low energy hadron production, such improvements will increase the reliability of air shower simulations.  Finally it should be noted that multi-detector installations such as KASCADE can be used to check the consistency of the simulation of hadronic interactions by comparing different observables of EAS and their correlation with model predictions.  \\noindent  {\\bf Acknowledgements:} The authors thank Dieter Heck for many fruitful discussions and help with modifying CORSIKA to include the muon ancestor information."
    },
    "0512/astro-ph0512193_arXiv.txt": {
        "abstract": "{Aims. Filamentary structures of early type stars are found to be a common feature of the Magellanic Clouds formed at an age of about $0.9-2\\times10^8\\,yr$. As we go to younger ages these large structures appear fragmented and sooner or later form young clusters and associations. In the optical domain we have detected 56 such large structures  of young objects, known as stellar complexes in the LMC for which we give coordinates and dimensions. We also investigate star formation activity and evolution of these stellar complexes and define the term ``starburst region''. Methods. IR  properties of these regions have been investigated using IRAS data. A colour-magnitude diagram (CMD) and a two-colour diagram from IRAS data of these regions ware compared with observations of starburst galaxies and cross-matching with HII regions and SNRs was made . Radio emission maps at 8.6-GHz and the CO (1$\\rightarrow$0) line were also cross correlated with the map of the stellar complexes. Results. It has been found that nearly 1/3 of the stellar complexes are extremely active resembling the IR behaviour of starburst galaxies and HII regions. These stellar complexes illustrating such properties are called here ``starburst regions''. They host an increased number of HII regions and SNRs. The main starburst tracers are their IR luminosity ($F_{60}$ well above 5.4 Jy) and the 8.6-GHz radio emission. Finally the evolution of all stellar complexes is discussed based on the CO emission. ",
        "introduction": "Due to its proximity to our Galaxy and  the SMC, the LMC offers two important advantages: i) the possibility to study its stellar content down to low stellar masses and ii) the possibility to follow the consequences of the close interaction, about $2-4 \\times 10^8\\,yr$ ago (Gardiner \\& Noguchi \\cite{gardiner}; Kunkel, Demers \\& Irwin \\cite{kunkel}), with the SMC. Stellar complexes are defined as large  structures (clump like) in galaxies dominated by recently formed stellar component mixed with very young clusters, stellar associations and gas (Martin et al. \\cite{martin}; van den Bergh \\cite{van}; Elmegreen \\& Elmegreen \\cite{elmegreen2}; Feitzinger \\& Braunsfurth \\cite{feitz}; Elmegreen \\& Elmegreen \\cite{elmegreen3}; Ivanov \\cite{ivanov}; Larson et al. \\cite{larson1}; Efremov\\cite{efremov1}; Larson et al. \\cite{larson2}; Efremov \\& Chernin \\cite{efremov3}; Elmegreen et al. \\cite{elmegreen4}; Kontizas et al. \\cite{kontizas1}; Efremov \\cite{efremov2}; Battinelli, Efremov \\& Magnier \\cite{battinelli}; Elmegreen \\& Efremov \\cite{elmegreen1}; Kontizas et al. \\cite{kontizas2}). So far they have been detected as regions with young stellar component and enhanced stellar number density, compared to the surrounding region. In the LMC  stellar complexes have been detected using UK Schmidt plates (Maragoudaki et al. \\cite{maragoudaki1}). In the optical domain we have detected 56 large  structures  of young objects. These stellar complexes are associated with the already known nine Shapley constellations (Shapley \\cite{shapley}), which are referred to as large stellar structures of early type  supergiants.  Helou (\\cite{helou}) found that the IRAS $log(F_{60}/F_{100})$ versus $log(F_{12}/F_{25})$ colour-colour diagram of ``normal galaxies'' shows a distribution that extends continuously from the cool ($\\sim$20K) relatively constant ``cirrus'' emission from the neutral medium to the warmer ($\\sim$40-50K) emission from the active, starburst end.  IRAS far-IR fluxes for 6 HII regions in the LMC have been studied by DeGioia-Eastwood (\\cite{gioia}). These regions are the sites of massive star formation, where the radiative heating source is young stars rather than the general interstellar radiation field. Such regions are expected to lie in stellar complexes and/or to be the stage just before the stellar complex is formed.  The highly structured diffuse X-ray emission of the LMC has been imaged in detail by ROSAT. The brightest regions are found east of LMC X-1  and in the 30 Doradus region. The latter is strikingly similar to the optical picture. There is a strong correlation between diffuse features in the X-ray image and ESO colour images of the LMC in visible light. Bright knots in the X-ray map correspond to HII emission in the optical (Westerlund \\cite{west}). The distribution of the SNRs in the LMC shows (i) a clumping of objects in the 30 Doradus region, (ii) several remnants within the Bar of the LMC and (iii) that the remainder of the remnants are found in super-associations. Most SNRs in the Bar are in regions where many young clusters are located, clusters as young as $\\le\\,10\\,Myr$ (Westerlund \\cite{west}).  The LMC has been well studied in the radio continuum, which is connected to the formation of new and massive stars. Fukui et al. (\\cite{fukui1}, \\cite{fukui2}) found that the molecular clouds in the LMC have a good correlation with the youngest ($\\la$10~Myr) stellar clusters, while there is little spatial correlation of these clouds with SNRs or with the older stellar clusters.  In this paper we describe how the stellar complexes were detected, using optical data and we give their location and dimensions. The total flux of each structure is presented. Their derived IR  properties are compared to IR properties of galaxies in order to identify the distinction between ordinary stellar complexes and active star formation regions. Possible connection of the complexes determined here with known X-ray detected SNRs and candidates as well as HII regions is discussed. Finally, their correlation between high resolution radio data at 8.6-GHz and CO emission is examined, in order to determine the activity of the complexes and their evolution. ",
        "conclusions": "From optical observations, 56 stellar complexes are revealed in an area of $6\\times 7 \\deg^{2}$ of the LMC. The IR properties are studied using the IRAS fluxes. The ratio of the flux densities $F_{60}/F_{100}$, the flux density $F_{60}$ as discussed above, the 8.6-GHz and CO radio emission, will be used in order to classify these complexes.  As mentioned in section 3, Lehnert \\& Heckman (\\cite{lehnert}) showed that starburst galaxies are IR ``warm'' and ``bright''. In the same manner, we adopt the equivalent term {\\bf ``starburst regions''} for complexes that are found to fulfil these criteria, namely $F_{60}/F_{100} \\ge 0.4$ and $F_{60} \\ge 5.4$ Jy. In addition, complexes that do not fulfil the previously stated criteria will be called {\\bf ``active complexes''}.  A more detailed study of high resolution 8.6-GHz map (Fig.~\\ref{co8.6}) revealed that all complexes with F$_{60}$ well above the 5.4 Jy limit are very well correlated with the 8.6-GHz radio emission. In contrast, complexes with F$_{60}$ well below the 5.4 Jy limit lack 8.6-GHz emission. However, 7 out of the 56 complexes (A5, A6, A7, C5, C7, C22 and A31) have F$_{60}$ close to the classification limit and deviate from the above behaviour. Complex C22, for example, is well correlated with 8.6-GHz emission; its {\\it F$_{60}$}(5.07$\\pm0.20$ Jy) though is below the accepted limit. In contrast A31 has {\\it F$_{60}$=5.98$\\pm0.20$ Jy}, which would classify it as a starburst region; however, there is no 8.6-GHz emission associated with it. Since the $F_{60}$ criterion is not trustworthy for these regions, we additionally use the 8.6-GHz emission to classify the complexes as ``starburst candidates'' or ``active complexes candidates'' depending on its existence. Hence, using the 8.6-GHz radiation as an additional tracer in tandem with the $F_{60}$, we can classify the stellar complexes into 13 starbursts, 5 starburst candidates, 2 active complexes candidates and 36 active complexes as seen in Table ~\\ref{starbursts}. The characterisation of each complex is given in Table ~\\ref{Properties}, column 10.  \\begin{table} \\begin{center} \\caption{Complexes characterisation.} \\begin{tabular}{lccc} \\hline \\hline Complex Type             & $F_{60}$        & 8.6      &  No. of  \\\\ & (Jy)            & (GHz)    & Complexes \\\\ \\hline starburst                & $F_{60}>$5.4    & yes      &   13              \\\\ starburst candidate      & $F_{60}\\la$5.4  & yes      &   5               \\\\ active complex candidate & $F_{60}\\ga$5.4  & no       &   2               \\\\ active complex           & $F_{60}<$5.4    & no       &   36              \\\\ \\hline \\end{tabular} \\label{starbursts} \\end{center} \\end{table}  Moreover, a discrimination concerning the evolution of all the complexes was attempted based on the CO data. More than half of the starburst and starburst candidate regions show enhanced CO emission indicating ongoing and future evolution, while the rest of them are thought to be more or less evolved. Regarding the active stellar complexes (and the two candidates), we found only 5 of them with significant amount of CO emission, giving additional indication of current star formation and revealing potential future starbursts. The lack of CO in the rest of them could indicate that these complexes have never been starbursts or that their starburst activity previously exhausted the molecular gas.   \\begin{figure} \\centering \\includegraphics[angle=0,width=8cm]{Figure4.ps} \\caption{The HII regions (triangles) and SNRs (squares), plotted over the stellar complexes as defined here.} \\label{SBR} \\end{figure}  In Fig.~\\ref{SBR} we plotted the known SNRs and HII regions based on the literature (Williams et al. \\cite{williams}; Haberl \\& Pietsch \\cite{haberl}; Filipovic et al. \\cite{filip} and Sasaki, Haberl \\& Pietsch \\cite{sasaki}) over the stellar complexes with triangles and squares respectively. The overplotted HII regions show an increased concentration in and around the very active star forming regions, as expected. Most of the defined starburst areas (and the candidates) are associated with SNRs. The detected complexes are also cross-matched with the catalogue of stellar associations (Gouliermis et al. \\cite{gouliermis}) and that of nebulae in the LMC (Davies, Elliot \\& Meaburn \\cite{davies}). The majority of the complexes are loci of stellar associations, whereas all of them are associated with a large number of nebulosities (Table ~\\ref{Properties}).   \\begin{table*}[h] \\begin{center} \\tiny{ \\caption{List of the LMC stellar complexes. Col. 1 indicates the adopted identification name. A denotes aggregate, C complex and SC super-complex. For a structure found inside a larger one, the name of the later is given in parenthesis. Col. 2, 3 and 4 give the RA, DEC of the ``centre'' of each stellar complex and its dimension. Col. 5,6, 7 show the number of SNRs and candidates, stellar associations and number of Nebulae and Henize objects found in each complex. Col. 8 and 9 indicate presence of 8.6-GHz and CO line emission respectively. Finally in Col. 10 the complexes are characterised. I indicates starbursts, II starburst candidates, III active complexes candidates and IV active complexes.} \\begin{tabular}{lccccccccc} Stellar grouping & RA (1950) & DEC (1950) & Dimension & No.SNRs &Association ID & No.Nebulae & 8.6 & CO & Characterisation\\\\ & h   m   s  & deg  m  s  & (pc)&  & & & GHz &line & \\\\ \\hline \\hline SC1       & 05:44:39 & -67:15:18 & 1421.5 & 1            & 173,257     & 46 & -       & -       & IV \\\\ C1 (SC1)  & 05:43:30 & -67:17:29 & 355.3  & 0            & 173,257     & 57 & -       & -       & IV \\\\ A1        & 05:52:53 & -67:29:46 & 205    & 0            & 173,257     & 22 & -       & -       & IV \\\\ A2        & 05:43:44 & -67:19:37 & 191    & 1            & 173,257     & 31 & -       & -       & IV \\\\ SC2       & 05:31:22 & -66:54:42 & 1526.5 & 7            &             & 24 & -       & -       & IV \\\\ C2 (SC2)  & 05:32:45 & -67:05:11 & 749    & 1            & 173         & 41 & -       & -       & IV \\\\ C3 (C2)   & 05:32:47 & -66:58:22 & 572    & 1            &             & 32 & -       & -       & IV \\\\ A3 (SC2)  & 05:29:45 & -66:37:17 & 204    & 0            &             & 31 & -       & -       & IV \\\\ A4 (C2)   & 05:30:53 & -66:39:15 & 245    & 0            &             & 26 & -       & -       & IV \\\\ A5 (SC2)  & 05:32:55 & -67:32:29 & 164    & 1            & 173,257     & 41 & -       & -       & III\\\\ A6        & 05:25:58 & -66:12:18 & 299    & 1            &             & 20 & $\\surd$ & $\\surd$ & II \\\\ A7        & 05:27:01 & -67:27:42 & 299    & 1            & 173,257     & 10 & $\\surd$ & -       & II \\\\ C4        & 05:13:28 & -67:20:57 & 314    & 0            & 173,257     & 16 & -       & -       & IV \\\\ A8 (C4)   & 05:13:10 & -67:22:34 & 246    & 0            & 173,257     & 16 & -       & -       & IV \\\\ A9        & 05:10:55 & -67:11:35 & 192    & 0            & 173,257     & 26 & -       & $\\surd$ & IV \\\\ A10       & 04:57:01 & -66:28:27 & 245    & 1            &             & 7  & $\\surd$ & $\\surd$ & I  \\\\ C5        & 05:20:58 & -68:10:16 & 765    & 2            &             & 17 & $\\surd$ & $\\surd$ & II \\\\ C6 (C5)   & 05:22:13 & -67:58:26 & 342.6  & 1            & 173,257     & 36 & $\\surd$ & $\\surd$ & I  \\\\ A11 (C6)  & 05:22:13 & -67:56:44 & 205    & 1            & 173,257     & 17 & $\\surd$ & $\\surd$ & I  \\\\ C7        & 05:07:30 & -68:44:27 & 545    & 1            &             & 39 & $\\surd$ & $\\surd$ & II \\\\ A12 (C7)  & 05:09:04 & -68:49:37 & 287    & 1            &             & 17 & $\\surd$ & $\\surd$ & I  \\\\ A13 (C7)  & 05:05:38 & -68:39:12 & 245    & 0            &             & 27 & -       & -       & IV \\\\ A14       & 05:04:00 & -68:57:45 & 230    & 0            &             & 11 & -       & -       & IV \\\\ C8        & 04:55:33 & -69:32:52 & 628    & 1            &             & 12 & -       & -       & IV \\\\ A15 (C8)  & 04:56:55 & -69:30:18 & 273    & 1            &             & 10 & -       & -       & IV \\\\ C9        & 05:09:26 & -70:08:58 & 328    & 0            & 150,153,151 & 19 & -       & -       & IV \\\\ A16 (C9)  & 05:08:09 & -70:05:05 & 246    & 0            & 150,151     & 10 & -       & -       & IV \\\\ SC3       & 05:04:44 & -70:29:23 & 1093.2 & 2            & 150,153,151 & 23 & -       & -       & IV \\\\ C10 (SC3) & 05:05:01 & -70:32:56 & 478.3  & 1            & 150,153,151 & 13 & -       & -       & IV \\\\ C11(C10)  & 05:06:16 & -70:40:18 & 328    & 0            & 150,153,151 & 17 & -       & -       & IV \\\\ C12       & 04:58:02 & -70:53:27 & 574    & 0            &             & 2  & -       & -       & IV \\\\ A17       & 04:56:28 & -71:28:37 & 164    & 0            &             & 0  & -       & -       & IV \\\\ C13       & 05:30:55 & -71:50:37 & 574    & 0            & 76          & 6  & -       & -       & IV \\\\ C14       & 05:39:55 & -71:11:22 & 328    & 0            & 76          & 10 & -       & $\\surd$ & IV \\\\ C15       & 05:35:42 & -71:12:22 & 437    & 0            & 76          & 17 & -       & $\\surd$ & IV \\\\ A18       & 05:31:26 & -71:06:11 & 218.64 & 2            & 76          & 21 & $\\surd$ & $\\surd$ & I  \\\\ C16       & 05:47:50 & -70:41:07 & 615    & 1 very close & 150,151     & 15 & -       & $\\surd$ & IV \\\\ A19 (C16) & 05:47:53 & -70:44:31 & 210    & 0            & 150,153,151 & 14 & -       & $\\surd$ & IV \\\\ A20       & 05:49:53 & -70:05:51 & 164    & 0            & 150,153,151 & 10 & -       & -       & IV \\\\ C17       & 05:39:42 & -69:21:58 & 724    & 10           & 172         & 23 & $\\surd$ & $\\surd$ & I  \\\\ C18 (C17) & 05:38:17 & -69:26:16 & 463    & 3            & 172         & 50 & $\\surd$ & -       & I  \\\\ A21 (C17) & 05:36:42 & -69:10:51 & 272.3  & 5            & 172         & 29 & $\\surd$ & -       & I  \\\\ A22 (C18) & 05:39:34 & -69:28:00 & 218    & 1            & 172         & 47 & $\\surd$ & -       & I  \\\\ A23 (C18) & 05:36:59 & -69:24:30 & 245    & 1            & 172         & 32 & $\\surd$ & -       & I  \\\\ A24 (C17) & 05:40:52 & -69:38:12 & 298    & 2            & 172         & 44 & $\\surd$ & -       & I  \\\\ A25       & 05:35:48 & -68:55:28 & 190.61 & 0            &             & 29 & $\\surd$ & $\\surd$ & I  \\\\ A26       & 05:35:18 & -69:43:05 & 231.45 & 0            & 172         & 46 & $\\surd$ & -       & I  \\\\ A27       & 05:36:33 & -70:10:26 & 272.3  & 0            & 150,153,151 & 32 & -       & -       & IV \\\\ C19       & 05:30:00 & -69:11:52 & 792.57 & 2            & 172         & 10 & -       & -       & IV \\\\ C20 (C19) & 05:29:42 & -69:08:07 & 517.37 & 2            & 172         & 46 & -       & -       & IV \\\\ A28 (C20) & 05:31:34 & -69:15:19 & 163    & 0            & 172         & 63 & -       & -       & IV \\\\ A29 (C20) & 05:28:27 & -69:04:25 & 285    & 2            & 172         & 71 & -       & -       & IV \\\\ C21       & 05:26:03 & -69:51:28 & 463    & 1 very close & 172         & 36 & -       & -       & IV \\\\ A30 (C21) & 05:27:03 & -69:50:03 & 163    & 0            & 172         & 28 & -       & -       & IV \\\\ C22       & 05:19:27 & -69:34:39 & 763    & 2            & 177         & 38 & $\\surd$ & $\\surd$ & II \\\\ A31 (C22) & 05:17:09 & -69:31:39 & 300    & 0            & 772         & 20 & -       & -       & III\\\\ \\hline \\end{tabular} \\label{Properties} } \\end{center} \\end{table*}  \\begin{table*}[h] \\begin{center} \\tiny{ \\caption{List of the derived IRAS fluxes (MJy/sr per pixel), luminocities (in units of solar luminocities) of the identified stellar groupings at 12, 25, 60 and 100 $\\mu m$, the IRAS ratio $F_{60}/F_{100}$ and the flux at 60\\, $\\mu m$ in Jy per pixel.} \\begin{tabular}{lcccccccccc} Stellar grouping & f12 & L12 & f25 & L25 & f60 & L60 & f100 & L100 & $F_{60}/F_{100}$ & $ F_{60} $   \\\\ \\hline \\hline SC1          & 0.18 & 2.8E+06 &  0.50 & 3.7E+06 &   3.63 & 1.1E+07 &   8.87 & 1.7E+07 & 0.41 &  0.69 $ \\pm $ 0.25 \\\\ C1 (SC1)     & 0.31 & 0.3E+06 &  1.02 & 0.5E+06 &   5.38 & 1.1E+06 &  13.38 & 1.6E+06 & 0.40 &  1.02 $ \\pm $ 0.23 \\\\ A1           & 0.36 & 0.1E+06 &  1.28 & 0.2E+06 &   0.67 & 0.4E+06 &   2.16 & 0.8E+06 & 0.31 &  0.13 $ \\pm $ 0.49 \\\\ A2           & 0.47 & 0.1E+06 &  1.54 & 0.2E+06 &   6.32 & 0.4E+06 &  19.39 & 0.7E+06 & 0.33 &  1.20 $ \\pm $ 0.23 \\\\ SC2          & 1.14 & 2.0E+07 &  4.46 & 3.8E+07 &  13.96 & 5.0E+07 &  33.50 & 7.2E+07 & 0.42 &  2.65 $ \\pm $ 0.21 \\\\ C2 (SC2)     & 0.16 & 0.7E+06 &  0.56 & 1.1E+06 &   4.03 & 3.5E+06 &   9.34 & 4.9E+06 & 0.43 &  0.77 $ \\pm $ 0.24 \\\\ C3 (C2)      & 0.11 & 0.3E+06 &  0.52 & 0.6E+06 &   2.08 & 1.1E+06 &   5.13 & 1.5E+06 & 0.41 &  0.40 $ \\pm $ 0.29 \\\\ A3 (SC2)     & 0.17 & 0.5E+06 &  1.10 & 0.2E+06 &   0.98 & 0.6E+05 &   3.51 & 0.1E+06 & 0.28 &  0.19 $ \\pm $ 0.40 \\\\ A4 (C2)      & 0.12 & 0.6E+05 &  0.89 & 0.2E+06 &   1.06 & 0.1E+06 &   3.46 & 0.2E+06 & 0.31 &  0.20 $ \\pm $ 0.38 \\\\ A5 (SC2)     & 1.39 & 0.3E+06 &  3.80 & 0.4E+06 &  30.79 & 1.3E+06 &  64.60 & 0.6E+06 & 0.48 &  5.85 $ \\pm $ 0.20 \\\\ A6           & 1.16 & 0.8E+06 &  2.80 & 0.9E+06 &  25.17 & 3.4E+06 &  58.76 & 4.8E+06 & 0.43 &  4.78 $ \\pm $ 0.20 \\\\ A7           & 0.82 & 0.6E+06 &  2.89 & 0.9E+06 &  27.08 & 3.7E+06 &  50.10 & 4.1E+06 & 0.54 &  5.15 $ \\pm $ 0.20 \\\\ C4           & 0.68 & 0.5E+06 &  1.53 & 0.6E+06 &  11.15 & 1.7E+06 &  26.38 & 0.4E+06 & 0.42 &  2.12 $ \\pm $ 0.21 \\\\ A8 (C4)      & 0.80 & 0.4E+06 &  1.90 & 0.4E+06 &  12.30 & 1.1E+06 &  28.14 & 1.5E+06 & 0.44 &  2.34 $ \\pm $ 0.21 \\\\ A9           & 0.52 & 0.1E+06 &  1.66 & 0.2E+06 &   8.07 & 0.5E+06 &  23.91 & 0.8E+06 & 0.34 &  1.53 $ \\pm $ 0.22 \\\\ A10          & 2.26 & 1.0E+06 &  8.83 & 1.9E+06 &  69.77 & 6.5E+06 & 127.91 & 7.1E+06 & 0.55 & 13.26 $ \\pm $ 0.20 \\\\ C5           & 1.61 & 7.2E+06 &  4.71 & 1.0E+07 &  28.23 & 2.5E+07 &  61.20 & 3.3E+07 & 0.46 &  5.36 $ \\pm $ 0.20 \\\\ C6 (C5)      & 0.10 & 0.9E+05 &  6.66 & 2.9E+06 &  54.65 & 9.9E+06 &  97.19 & 1.1E+07 & 0.56 & 10.38 $ \\pm $ 0.20 \\\\ A11 (C6)     & 3.21 & 1.0E+06 & 11.89 & 1.8E+06 &  89.85 & 5.8E+06 & 142.78 & 5.6E+06 & 0.63 & 17.07 $ \\pm $ 0.20 \\\\ C7           & 0.96 & 2.2E+06 &  2.31 & 2.5E+06 &  20.22 & 9.3E+06 &  42.25 & 1.2E+07 & 0.49 &  3.84 $ \\pm $ 0.21 \\\\ A12 (C7)     & 1.89 & 1.2E+06 &  5.31 & 1.6E+06 &  42.20 & 5.4E+06 &  79.59 & 6.1E+06 & 0.53 &  8.02 $ \\pm $ 0.20 \\\\ A13 (C7)     & 0.59 & 0.3E+06 &  1.56 & 0.4E+06 &   9.20 & 0.8E+06 &  22.76 & 1.3E+06 & 0.40 &  1.75 $ \\pm $ 0.22 \\\\ A14          & 0.94 & 0.4E+06 &  1.90 & 0.4E+06 &  13.16 & 1.1E+06 &  35.27 & 1.7E+06 & 0.37 &  2.50 $ \\pm $ 0.21 \\\\ C8           & 0.90 & 2.7E+06 &  2.44 & 3.6E+06 &  16.74 & 1.0E+07 &  33.96 & 1.2E+07 & 0.49 &  3.18 $ \\pm $ 0.21 \\\\ A15 (C8)     & 0.81 & 0.4E+06 &  1.86 & 0.5E+06 &  11.02 & 1.3E+06 &  25.97 & 1.8E+06 & 0.42 &  2.09 $ \\pm $ 0.21 \\\\ C9           & 1.20 & 1.0E+06 &  3.67 & 1.5E+06 &   9.98 & 1.7E+06 &  25.75 & 2.6E+06 & 0.39 &  1.90 $ \\pm $ 0.22 \\\\ A16 (C9)     & 0.72 & 0.3E+06 &  1.77 & 0.4E+06 &   8.95 & 0.8E+06 &  22.08 & 1.2E+06 & 0.41 &  1.70 $ \\pm $ 0.22 \\\\ SC3          & 0.40 & 3.7E+06 &  0.87 & 3.9E+06 &   3.87 & 7.1E+06 &   9.74 & 1.1E+07 & 0.40 &  0.74 $ \\pm $ 0.25 \\\\ C10 (SC3)    & 0.46 & 0.8E+06 &  1.07 & 0.9E+06 &   3.38 & 1.2E+06 &   8.85 & 1.9E+06 & 0.38 &  0.64 $ \\pm $ 0.25 \\\\ C11 (C10)    & 0.60 & 0.5E+06 &  1.44 & 0.6E+06 &   4.82 & 0.8E+06 &  12.47 & 1.2E+06 & 0.39 &  0.92 $ \\pm $ 0.24 \\\\ C12          & 0.37 & 0.9E+06 &  0.87 & 1.1E+06 &   1.82 & 0.9E+06 &   5.75 & 1.8E+06 & 0.32 &  0.35 $ \\pm $ 0.30 \\\\ A17          & 0.69 & 0.1E+06 &  2.14 & 0.2E+06 &   1.89 & 0.8E+05 &   7.51 & 0.2E+06 & 0.25 &  0.36 $ \\pm $ 0.30 \\\\ C13          & 0.45 & 1.1E+06 &  0.79 & 9.6E+05 &   2.51 & 1.3E+06 &   8.61 & 2.6E+06 & 0.29 &  0.48 $ \\pm $ 0.27 \\\\ C14          & 0.99 & 0.8E+06 &  1.81 & 0.7E+06 &  10.17 & 1.7E+06 &  31.07 & 3.1E+06 & 0.33 &  1.93 $ \\pm $ 0.22 \\\\ C15          & 0.62 & 0.9E+06 &  1.13 & 0.8E+06 &   6.19 & 1.8E+06 &  18.46 & 3.3E+06 & 0.34 &  1.18 $ \\pm $ 0.23 \\\\ A18          & 1.86 & 0.7E+06 &  4.59 & 0.8E+06 &  38.36 & 2.8E+06 &  78.27 & 3.5E+06 & 0.49 &  7.29 $ \\pm $ 0.20 \\\\ C16          & 0.67 & 1.9E+06 &  0.87 & 1.2E+06 &   4.16 & 2.4E+06 &  17.37 & 6.1E+06 & 0.24 &  0.79 $ \\pm $ 0.24 \\\\ A19 (C16)    & 1.08 & 0.4E+06 &  1.86 & 0.3E+06 &   4.24 & 0.3E+06 &  21.73 & 0.9E+06 & 0.20 &  0.81 $ \\pm $ 0.24 \\\\ A20          & 0.05 & 0.1E+05 &  0.01 & 0.1E+04 &  13.07 & 0.5E+06 &  27.62 & 0.7E+06 & 0.47 &  2.48 $ \\pm $ 0.21 \\\\ C17          & 7.21 & 2.9E+07 & 36.90 & 7.2E+07 & 248.00 & 2.0E+08 & 327.45 & 1.6E+08 & 0.76 & 47.12 $ \\pm $ 0.20 \\\\ C18 (C17)    & 3.88 & 6.4E+06 & 16.18 & 1.3E+07 & 149.37 & 4.9E+07 & 222.33 & 4.4E+07 & 0.67 & 28.38 $ \\pm $ 0.20 \\\\ A21 (C7)     & 7.62 & 4.3E+06 & 43.93 & 1.2E+07 & 293.57 & 3.4E+07 & 369.90 & 2.5E+07 & 0.79 & 55.78 $ \\pm $ 0.20 \\\\ A22 (C18)    & 4.26 & 1.6E+06 & 15.66 & 2.8E+06 & 153.80 & 1.1E+07 & 238.40 & 1.0E+07 & 0.65 & 29.22 $ \\pm $ 0.20 \\\\ A23 (C18)    & 2.30 & 1.1E+06 &  8.35 & 1.9E+06 &  77.75 & 7.2E+06 & 128.05 & 7.1E+06 & 0.61 & 14.77 $ \\pm $ 0.20 \\\\ A24          & 5.50 & 3.8E+06 & 21.67 & 7.1E+06 & 190.57 & 2.6E+07 & 276.81 & 2.3E+07 & 0.69 & 36.20 $ \\pm $ 0.20 \\\\ A25          & 1.42 & 0.4E+06 &  4.05 & 0.5E+06 &  43.42 & 2.4E+06 &  75.47 & 2.5E+06 & 0.57 &  8.25 $ \\pm $ 0.20 \\\\ A26          & 2.62 & 1.1E+06 &  7.49 & 1.5E+06 &  73.60 & 6.1E+06 & 126.18 & 6.3E+06 & 0.58 & 13.98 $ \\pm $ 0.20 \\\\ A27          & 1.64 & 0.9E+06 &  2.51 & 0.7E+06 &  24.35 & 2.8E+06 &  63.05 & 4.3E+06 & 0.39 &  4.63 $ \\pm $ 0.20 \\\\ C19          & 0.77 & 3.7E+06 &  1.87 & 4.3E+06 &  16.38 & 1.6E+07 &  32.32 & 1.9E+07 & 0.51 &  3.11 $ \\pm $ 0.21 \\\\ C20 (C19)    & 0.79 & 1.6E+06 &  1.90 & 1.9E+07 &  14.00 & 5.8E+06 &  28.65 & 7.1E+06 & 0.49 &  2.66 $ \\pm $ 0.21 \\\\ A28 (C20)    & 1.30 & 0.3E+06 &  3.42 & 0.3E+06 &  19.26 & 0.8E+06 &  39.05 & 1.0E+06 & 0.49 &  3.66 $ \\pm $ 0.21 \\\\ A29 (C20)    & 1.01 & 0.6E+06 &  2.66 & 0.8E+06 &  14.18 & 1.8E+06 &  29.49 & 2.2E+06 & 0.48 &  2.69 $ \\pm $ 0.21 \\\\ C21          & 1.02 & 1.7E+06 &  1.74 & 1.4E+06 &  14.49 & 4.8E+06 &  32.92 & 6.6E+06 & 0.44 &  2.75 $ \\pm $ 0.21 \\\\ A30 (C21)    & 1.01 & 0.2E+06 &  2.40 & 0.2E+06 &   7.42 & 0.3E+06 &  22.65 & 5.6E+06 & 0.33 &  1.41 $ \\pm $ 0.22 \\\\ C22          & 1.28 & 5.7E+06 &  2.41 & 5.2E+06 &  26.67 & 2.4E+07 &  50.59 & 2.7E+07 & 0.53 &  5.07 $ \\pm $ 0.20 \\\\ A31 (C22)    & 1.53 & 1.0E+06 &  3.05 & 1.0E+06 &  31.41 & 4.4E+06 &  56.30 & 4.7E+06 & 0.56 &  5.98 $ \\pm $ 0.20 \\\\ \\hline \\end{tabular} \\label{IrasFluxes} } \\end{center} \\end{table*}"
    },
    "0512/astro-ph0512120_arXiv.txt": {
        "abstract": "The ZEPLIN collaboration has recently published its first result presenting a maximum sensitivity of $1.1 \\times 10^{-6}$ picobarn for a WIMP mass of $\\approx$ 60 GeV. The analysis is based on a discrimination method using the different time distribution of scintillation light generated in electron recoil and nuclear recoil interactions. We show that the methodology followed both for the calibration of the ZEPLIN-I detector response and for the estimation of the discrimination power is not reliable enough to claim any background discrimination at the present stage. The ZEPLIN-I sensitivity appears then to be in the order of 10$^{-3}$ picobarn, three orders of magnitude above the claimed 1.1 10$^{-6}$ picobarn. ",
        "introduction": "Most of the progress in Dark Matter Direct Detection sensitivity is associated with the advent of a new generation of discriminating detectors, able to reject the important radioactive background associated mostly with electron recoils, and making it possible to identify a small population of nuclear recoil interactions. Three cryogenic experiments, CDMS~\\cite{cdms}, EDELWEISS~\\cite{edelweiss} and CRESST~\\cite{cresst}, have applied this discrimination scheme using the simultaneous measurement of charge and phonon signals, and of light and phonon signals. These three experiments have published sensitivities to spin-independent WIMP interactions ranging from $\\approx$~1.6 10$^{-7}$ picobarn to $\\approx$~2.0 10$^{-6}$  picobarn at maximum sensitivity~\\cite{cdms,edelweiss,cresst}. The ZEPLIN-I experiment has followed a parallel discrimination strategy by measuring two quantities, the visible energy and the scintillation time constant. Despite a much higher background rate, an energy resolution of 100 \\% or larger for events of interest and a statistical discrimination, this experiment claims a sensitivity of 1.1 10$^{-6}$ picobarn, requiring a background subtraction at the 99.9 \\% level~\\cite{zeplin}. ",
        "conclusions": "No demonstration has been given in the ZEPLIN-I analysis that nuclear recoils and electron recoils differ at the low energies (a few keV of visible energy) where most WIMP interactions are expected. Previous calibrations by the same group indicate just the opposite: nuclear and electron recoil scintillation time constants appear to converge when the detected energy is decreased.  In addition, the calibration procedure described in the latest ZEPLIN-I publication shows a major methodological problem: it is when the neutron source is removed that the highest neutron/gamma signal ratio is observed. A tentative but unproven population of \"ambient neutrons\", in the absence of a neutron source, is then used as a calibration, but a noise cut overlaps the \"ambient neutron\" selection. Therefore, it is not possible to exclude at present that this small \"ambient neutron\" population is in fact a photomultiplier noise tail, induced for instance by cosmic rays.  Waiting for more precise and consistent measurements, in an energy region where the present resolution is everywhere $\\geq 100\\%$, all events recorded at low energies should be considered as potential WIMP candidates. The maximum ZEPLIN-I sensitivity then becomes conservatively 10$^{-3}$ picobarn, three orders of magnitude above the claimed 1.1 10$^{-6}$ picobarn."
    },
    "0512/astro-ph0512470.txt": {
        "abstract": "{A molecular survey recently performed ahead of HH~2 supports the idea that the observed molecular enhancement is due to UV radiation from the HH object.} {The aim of the present work is to determine whether all HH objects with enhanced HCO$^+$ emission ahead of them also exhibit the same enhanced chemistry as HH~2. We thus observed several molecular lines at several positions ahead of five Herbig-Haro objects where enhanced HCO$^+$ emission was previously observed.} {We mapped the five Herbig-Haro objects using the IRAM-30 m. For each position we searched for more than one molecular species, and where possible for more than one transition per species. We then estimated the averaged beam column densities for all species observed and also performed LVG analyses to constrain the physical properties of the gas.} {The chemically richest quiescent gas is found ahead of the HH~7-11 complex, in particular at the HH~7-11 A position. In some regions we also detected a high velocity gas component.  We find that the gas densities are always higher than those typical of a molecular cloud while the derived temperatures are always quite low, ranging from 10 to 25 K. The emission of most species seems to be enhanced with respect to that of a typical dense clump, probably due to the exposure to a high UV radiation from the HH objects. Chemical differentiation among the positions is also observed. We attempt a very simple chemical analysis to explain such differentiation.} {} ",
        "introduction": "Herbig-Haro (HH) objects are regions where protostellar jets interact hydro-dynamically with the surrounding molecular cloud, resulting in strong atomic line emission characteristic of shocks.  There is now ample and increasing evidence that HH objects also have a less direct but very diverse chemical influence on nearby molecular clumps which have not yet been reached completely by the flow itself. In fact, the radiation generated in the HH objects, as well as the shocks produced by the HH objects alter the chemical composition of the molecular gas surrounding them in a very complex way.  An example of the complex interaction of the HH object with its surroundings, and the most studied one, is the HH~2 region (e.g. \\cite{girart02}; \\cite{dent03}; \\cite{lef05}; \\cite{girart05}). Some of the gas ahead of the HH object seems to be altered solely by the UV radiation: in the last decade or so quiescent, cold, molecular condensations have been found ahead of several HH objects. These cool condensations were found in enhanced emission of HCO$^+$ and NH$_3$, which are tracers of high density gas (e.g. \\cite{rudolph88}; \\cite{torrelles92}; \\cite{girart94}). These emissions probably arose because the UV radiation generated in the HH shock evaporated the icy mantles on dust grains in small density enhancements in the cloud near to the head of the jet (\\cite{girart94}).  Theoretical and further observational studies (\\cite{vw99}, hereafter VW99; \\cite{viti03} and \\cite{girart05}) confirmed this picture and showed that, because of the quiescent nature of such clumps (clear from the linewidths of the emission as well as the derived kinetic temperature of the gas), the high abundances of HCO$^+$ and other species could not be due to shock chemistry.   In this paper, we are particularly interested in these anomalous, quiescent enhanced molecular emissions from localized regions in the vicinity of the HH objects.  The VW99 models succeeded in explaining the known enhancements of HCO$^+$ and NH$_3$ relative to CS, predicted the amount by which these species should be increased, but also indicated other species which should be enhanced (by factors of 10--200 with respect to dark molecular clouds e.g L1554), in particular CH$_3$OH, H$_2$S, SO, SO$_2$, H$_2$CO, C$_3$H$_4$, H$_2$CS. Such predictions were confirmed by recent observations and further modelling of HH~2 (\\cite{girart02}, \\cite{viti03}; \\cite{girart05}). It is important to note however that while single dish observations (\\cite{girart02}, \\cite{viti03}) show that the chemistry of the illuminated $quiescent$ clumps ahead of the HH~2 object is essentially consistent with the UV field from the HH~2 front releasing the grain mantles and hence enhancing the gas chemistry, BIMA maps of such region in several molecules expected to be enhanced (\\cite{girart05}) have shown how the HH~2 object affects the chemistry and dynamic of the surrounding environment in different degrees. In summary, although Girart et al. (2005) concluded that a very complex morphological kinematical and chemical structure of the molecular gas is present ahead of the HH~2 object, % %four chemically differentiated regions ahead of the HH~2 object were %observed. While the chemistry in three of the four regions could only %be explained by a combination of UV, shock and PDR chemistry, the %quiescent clump (called SO$_2$ clump in \\cite{girart05}) was well %explained by the VW99 models where UV photoevaporation and %photochemistry are the main paths to the chemical enhancements. %this %study has they also confirmed that there is indeed an apparently quiescent molecular component that is more exposed to the UV radiation from the HH~2 front; the analysis of observational molecular lines from this region can be explained by UV radiation from the HH shock lifting the ice mantles from grains of a multi-density component clump and the consequent chemistry that arises is as complex and varied as indicated in the VW99 models.  The question we attempt to answer in this paper is the following: is the rich photochemistry characteristic of UV illumination peculiar to the HH~2 or is it common to all HH objects?  The simplest and most direct way to answer such question is to perform a survey of molecular species predicted to be enhanced by the VW99, ahead of several HH objects where enhanced quiescent HCO$^+$ and/or NH$_3$ have already been detected. We have now performed such a survey with the IRAM 30-meter and we report the observational results in this paper. The main aims of this paper are: 1) determine whether other HH objects exhibit the same peculiar enhanced chemistry of HH~2 as predicted by VW99 and 2) assuming the HH object to be the source of the UV radiation, look for correlations between the enhanced molecular emission and the distance of such emission to the HH object.  We report our observations in \\S~2; in \\S~3 we present our results and the derivation of beam averaged column densities. In \\S~4 we analyse our data by means of radiative transfer modelling. In \\S~5 we briefly discuss the chemical characteristics of the gas and in \\S~6 we summarize our findings. ",
        "conclusions": "We have performed a survey of molecular species ahead of several HH objects, where enhanced \\hco\\ emission was previously observed. Recent modelling have supported the idea that such enhancement is due to UV radiation from the HH object lifting the ice mantles from the grains in the clumpy medium ahead of the jet, which then would drive a photochemistry. If this model is correct then other molecular species (such as \\met, SO etc) should also be enhanced. A recent molecular survey of the gas ahead of HH~2 seems to support such model (\\cite{girart05}). The aims of this present survey were to investigate: 1) whether the rich chemistry observed ahead of HH~2, and predicted by the VW99 model, was indeed characteristic of the gas ahead of $all$ Herbig-Haro objects; 2) the chemical differentiation that may arise among clumps ahead of different objects, which may differ in radiation field strengths. \\par We have selected 4 HH objects regions: in some of these regions enhanced \\hco\\ emission had been detected from several clumps. Hence, in total we observed ten positions.  For each position we searched for more than one species and, if possible, for more than one transition per species. We find that, chemically, the richest gas was that ahead of the HH 7--11 complex, in particular at the HH~7--11 A position. In some regions we also detected a high velocity gas component. Beside estimating the averaged beam column densities for all the species observed, we have performed, where possible, LVG analyses to constrain the physical properties of the gas. The main conclusions from our LVG analysis are: \\begin{enumerate} \\item The gas densities of all the clumps observed are always higher than those typical of molecular clouds by at least a factor of 5, but probably higher. \\item The derived gas temperatures are alway quite low ($<$ 15 K) apart for the clumps ahead of HH~7--11 where the maximum temperature does not in any case exceed 25 K. \\item \\hco, SO and CS column densities derived from our LVG and RADEX analyses are higher than the beam averaged ones. This indicates that the size of the emitting region is smaller than the beam size. \\item The emitting size derived from the LVG analysis of CS are larger than those estimated from the \\hco\\ analysis. It may be that these two species arise from different gas components, consistent with previous studies (\\cite{girart05}). In addition, the \\hco\\ abundance is enhanced with respect to the starless core in L1544, whereas the CS has similar abundances. All this suggests that at least a fraction of the CS emission may arise from a cloud component or clumps not illuminated by he HH objects. \\end{enumerate}  By comparing column densities among the several clumps we also attempt a very simple chemical analysis in light of the photochemical model. We find that the abundances of the observed species ahead of the surveyed HH objects are higher than those found in a typical starless core and that by comparing the column densities among the different clumps we find that the clumps ahead of HH~7--11, HH~34 and HH~1 are of the same chemical nature, supporting the idea that the gas ahead of all HH objects exhibit a rich chemistry."
    },
    "0512/astro-ph0512585_arXiv.txt": {
        "abstract": "We investigate the birth and evolution of Galactic isolated radio pulsars. We begin by estimating their birth space velocity distribution from proper motion measurements of \\cite{2002ApJ...571..906B, 2003AJ....126.3090B}. We find no evidence for multimodality of the distribution and favor one in which the absolute one-dimensional velocity components are exponentially distributed and with a three-dimensional mean velocity of \\mbox{$380^{+40}_{-60}$ km s$^{-1}$}. We then proceed with a Monte Carlo-based population synthesis, modelling the birth properties of the pulsars, their time evolution, and their detection in the Parkes and Swinburne Multibeam surveys. We present a population model that appears generally consistent with the observations. Our results suggest that pulsars are born in the spiral arms, with a Galactocentric radial distribution that is well described by the functional form proposed by \\cite{2004A&A...422..545Y}, in which the pulsar surface density peaks at radius \\mbox{$\\sim3$ kpc}. The birth spin period distribution extends to several hundred milliseconds, with no evidence of multimodality. Models which assume the radio luminosities of pulsars to be independent of the spin periods and period derivatives are inadequate, as they lead to the detection of too many old simulated pulsars in our simulations. Dithered radio luminosities proportional to the square root of the spin-down luminosity accommodate the observations well and provide a natural mechanism for the pulsars to dim uniformly as they approach the death line, avoiding an observed pile-up on the latter. There is no evidence for significant torque decay (due to magnetic field decay or otherwise) over the lifetime of the pulsars as radio sources (\\mbox{$\\sim100$ Myr}). Finally, we estimate the pulsar birthrate and total number of pulsars in the Galaxy. ",
        "introduction": "\\label{intro} The birth and evolution of pulsars are of considerable interest. The spatial distribution of pulsars at birth may be used to associate them with their progenitors. Their initial spin periods may also be related to those of the progenitors' cores, which are not well predicted by theory due to differential rotation \\citep{1978ApJ...220..279E}. The birth properties of neutron stars are also intimately related to the physics of core-collapse supernovae in which most are thought to be formed. For example, their birth space velocities, spin periods, and magnetic fields put severe constraints on the mechanisms that may account for their high observed velocities \\citep{lai03}. The comparison of the Galactic pulsar birth and supernova rates may provide further insight into supernovae by quantifying the fraction which leaves behind a neutron star. It may also be possible to clarify whether all neutron stars unaffected by peculiar conditions (such as accretion from a binary companion or an anomalously high magnetic field) are in fact radio pulsars, although not always beamed toward us. Knowledge of the spin evolution of pulsars is particularly valuable in elucidating whether magnetic field decay occurs in isolated neutron stars. A correlation between the spin down of pulsars and the evolution of their radio luminosity may finally shed light on the long-standing problem of the pulsar radio emission mechanism.  One way to probe the birth and evolution of pulsars is to study the population as a whole. While the earlier efforts in this field relied on simplified analytical or semi-analytical treatments \\citep[e.g.,][]{1970ApJ...160..979G, 1971IAUS...46..165L, 1977ApJ...215..885T, 1977MNRAS.179..635D, 1985MNRAS.213..613L}, the trend over the last two decades has been to attempt detailed computational modelling of the pulsar population and the selection effects affecting the observed sample, often including the evolution of the synthetic pulsars from birth to detection \\citep[e.g.,][]{1987A&A...178..143S, 1989ApJ...345..931E, 1992A&A...254..198B, 1997A&A...322..477H, 1997MNRAS.289..592L, 2002ApJ...568..289A, 2002ApJ...565..482G, 2004ApJ...604..775G}. The broad goal of these ``Monte Carlo\" (MC) simulation studies is to statistically reproduce the pulsar sample observed in actual surveys and test whether a given population model is consistent with the data.  Unfortunately, the conclusions of pulsar population investigations have often been conflicting.  One long standing issue is that of magnetic field decay. For example, \\cite{1970ApJ...160..979G}, \\cite{1985MNRAS.213..613L}, \\cite{1990ApJ...352..222N}, \\cite{2002ApJ...565..482G}, and \\cite{2004ApJ...604..775G} have suggested that field decay occurs on short time scales \\mbox{$\\sim2.5-5$ Myr}. Similar statistical studies, e.g. those of \\cite{1987A&A...178..143S}, \\cite{1992A&A...254..198B}, and \\cite{1997MNRAS.289..592L}, reached the opposite conclusion that the magnetic field of pulsars does not decay significantly during their lifetime as radio sources, implying decay time constants \\mbox{$\\gtrsim100$ Myr}. Meanwhile, theoretical arguments have mostly supported that field decay is unimportant for typical neutron stars \\citep[][and references therein]{1969Natur.224..673B, 1992ApJ...395..250G}. Since the observed kinetic age versus characteristic age diagram (see, e.g., Lyne, Anderson, \\& Salter 1982\\nocite{1982MNRAS.201..503L} and Harrison, Lyne, \\& Anderson 1993\\nocite{1993MNRAS.261..113H}, who used it to argue in favor of field decay) has been shown to be an artefact of selection effects \\citep{1997MNRAS.289..592L}, direct empirical evidence for field decay is also lacking. However, evolution of the inclination angle between the magnetic and spin axes of a pulsar can also produce torque variations (see section \\ref{evolution rotational}) that are difficult to disentangle from evolution of the magnetic field itself on the basis of spin kinematics. We will thus henceforth speak more generally of ``torque decay\"\\footnote{The magnetic torque on a rotating neutron star decreases with time even if the magnetic field and its orientation are constant, owing to its decreasing rotational frequency. Here, we are referring to any potential decay in addition to this expected decrease.}.  Another matter of debate, introduced by \\cite{1981JApA....2..315V}, concerns the ``injection\" of a subpopulation of pulsars with birth spin period \\mbox{$\\sim$0.5 s}, in addition to a population of Crab-like pulsars born with short periods \\mbox{$\\lesssim$ 100 ms}. Evidence for injection was found by \\cite{1981JApA....2..315V} in an extension of the ``pulsar current\" analysis proposed by \\cite{1981MNRAS.194..137P}, in which one considers the flow of pulsars in the period-period derivative plane. Although the pulsar current analysis is model-free in the sense that it does not require explicit assumptions regarding the luminosity and spin-down laws of the pulsars, \\cite{1985MNRAS.213..613L} argued that the conclusion of \\cite{1981JApA....2..315V} was subject to considerable statistical uncertainty and suffered from an incomplete treatment of the selection effects plaguing pulsar surveys. Nevertheless, injection received further support from a number of independent analyses \\citep{1986ApJ...304..140C, 1987ApJ...319..162N, 1990ApJ...352..222N}, while similar studies maintained that the data did not require it \\citep{1987A&A...178..143S, 1992A&A...254..198B}. Refining the pulsar current analysis by restricting attention to pulsars above a luminosity cut-off of \\mbox{10 mJy kpc$^{2}$} at 400 MHz, for which significant statistics were available, \\cite{1993MNRAS.263..403L} found no evidence for injection, although they could not exclude that it may affect the fainter end of the pulsar population. Injection appears to have since lost popularity. Most recently, \\cite{2004ApJ...617L.139V} presented a pulsar current analysis of the Parkes Multibeam pulsar survey \\citep{2001MNRAS.328...17M} data which, although it suggests that many, perhaps \\mbox{40\\%}, of pulsars are born with periods in the range \\mbox{0.1-0.5 s}, did not show evidence of distinct subpopulations.  Lately, several authors have attempted to characterize the distribution of the birth space velocities of pulsars. Different analyses initially led to quantitative disagreements regarding the mean velocity \\citep{1994Natur.369..127L, 1997MNRAS.291..569H, 1997MNRAS.289..592L}. Somewhat surprisingly, other workers have found evidence in favor of a bimodal velocity distribution \\citep{1998ApJ...505..315C, 2002ApJ...568..289A, 2003AJ....126.3090B}, so that the issue of subpopulations of ordinary radio pulsars with different birth properties is not completely resolved. The origin of the two distribution components has not been conclusively identified. In fact, the two components themselves remain to be directly exhibited in the proper-motion data and the different authors disagree on their modes and relative weight. Analysing a large sample of pulsar proper motion measurements, \\cite{2005MNRAS.tmp..475H} directly challenged the multimodality of the distribution.  Multiple reasons motivate us to reconsider the birth and evolution of isolated radio pulsars in this paper.  First, pulsar astronomy has in the recent years seen a number of major advances. The Parkes (PM) and Swinburne (SM) Multibeam pulsar surveys \\citep{2001MNRAS.328...17M, 2001MNRAS.326..358E} at \\mbox{1.4 GHz} have approximately doubled the number of known pulsars. Over 1500 objects are now cataloged in the Australia Telescope National Facility (ATNF) Pulsar Database\\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat/} \\citep{2005AJ....129.1993M}. \\cite{NE2001} have introduced a new model of the Galactic free electron density (NE2001) to supersede the previously standard model of \\cite{1993ApJ...411..674T} (TC93), providing an updated dispersion measure distance scale. New astrometric measurements using interferometry provide a sample relatively free of brightness and distance biases from which to estimate the space velocity distribution of pulsars \\citep{2002ApJ...571..906B, 2003AJ....126.3090B}.  Second, in spite of poorly understood underlying physics, the recent simulations have become increasingly sophisticated. Departing from the common practice of modelling the pulsar pseudo-luminosity (where pseudo-luminosity $\\times$ distance$^{2} \\equiv$ flux density), \\cite{2002ApJ...568..289A} have for instance introduced a standard-candle geometrical model of physical luminosity. While this luminosity model is certainly a valuable step toward more realistic simulations, it is somewhat speculative in several respects (e.g., Gaussian shape and angular size of the radio beams, and the relative radiated power contribution of each). An important feature of the analysis of \\cite{2002ApJ...568..289A} is the use of indirect information (i.e. other than proper motions) contained in the observed pulsar sample in the inferrence of their birth velocity distribution of pulsars, relying on a detailed modelling of the selection effects of major radio surveys at 400 MHz. Since the accuracy of this modelling is limited by the validity of the assumptions made regarding the intrinsic properties of the pulsars, it is a fair question to ask whether their conclusion that the birth space velocity distribution of isolated pulsars has two well defined components, corresponding to low (\\mbox{$\\sim90$ km s$^{-1}$}) and high velocities (\\mbox{$\\sim500$ km s$^{-1}$}) is really required by the data. A previous analysis by \\cite{1998ApJ...505..315C} omitting a detailed treatment of the selection effects suggest that evidence for two components is indeed present in the data, but this claim has been disputed by \\cite{2005MNRAS.tmp..475H}. Adopting a slight modification of the \\cite{2002ApJ...568..289A} luminosity model to include spectral dependence, \\cite{2004ApJ...604..775G} claimed evidence for magnetic field decay on a time scale \\mbox{$\\sim2.8$ Myr}. Because pulsars are ultimately detected through their electromagnetic fluxes, it appears that the conclusions of pulsar population simulations are highly dependent on the assumed luminosity model. Since there is no strong independent support for any particular luminosity model, we may also consider whether the complexity of some of the simulations is absolutely needed.  Throughout, we follow two guiding principles. First, because of the important uncertainties in modelling pulsars from birth to detection, we rely on independent results obtained by more direct means whenever possible. Second, inspired by Ockham's razor, we strive for simplicity over sophistication. We further attempt to be as self-consistent as possible. While \\cite{2004ApJ...604..775G} implemented the NE2001 model, for instance, they based their luminosity model and their birth velocity distribution on those of \\cite{2002ApJ...568..289A}, who used TC93.  We begin by estimating the pulsar birth space velocity distribution directly from proper motion measurements in section \\ref{kick vel}. In section \\ref{pop synth}, we fix this distribution and proceed with Monte Carlo simulations of pulsar birth, evolution, and detection. We discuss our results and their implications in section \\ref{discussion} and conclude in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} Motivated by recent advances in pulsar astronomy (a large, homogeneous sample of detections by the Parkes and Swinburne Multibeam pulsar surveys; an updated model of the Galactic free electron density, NE2001; and new astrometric measurements), we have revisited the problem of the birth and evolution of isolated radio pulsars.  We started by estimating the pulsar birth velocity distribution directly from proper motion measurements of \\cite{2002ApJ...571..906B, 2003AJ....126.3090B} in section \\ref{kick vel}. A single-component Maxwellian distribution appears to be inadequate due to the detection of a few very high-velocity objects. However, we do not find evidence for multimodality of the velocity distribution, as alternative single-parameter models with heavier tails accommodate the observations equally well as a two-component Maxwellian. The exact shape of the velocity distribution is not well constrained. We adopted a model in which the absolute one-dimensional birth velocity components are exponentially distributed and with three-dimensional mean \\mbox{$\\langle v_{3D} \\rangle = 380^{+40}_{-60}$ km s$^{-1}$}.  In section \\ref{pop synth}, we used this velocity distribution as input to a more general Monte Carlo-based population synthesis. We described parametric prescriptions for the birth properties (location, velocity, spin period, magnetic field, and radio luminosity) of the pulsars and their time evolution (spatial, rotational, and radio emission shut-off). We then generated synthetic pulsar populations on which we performed simulations of the PM and SM surveys. By comparing the observed samples in the simulations with the real detections, we determined ``optimal\" model parameters.  The Galactocentric radial distribution of pulsar formation appears consistent with the functional form proposed by \\cite{2004A&A...422..545Y}, which incorporates a deficit in surface density near the Galactic Center. Although the \\cite{2004A&A...422..545Y} distribution was derived for the present-day distribution of evolved pulsars, our simulations suggest that younger pulsars are preferentially detected, a bias not treated in the \\cite{2004A&A...422..545Y}, and that it may in fact be a better approximation to the birth distribution. This distribution is qualitatively consistent (peaking midway between the Galactic Center and the Sun) with that of massive main-sequence stars, the purported pulsar progenitors, although there is a difference of \\mbox{$\\sim1.5$ kpc} between the radii of peak density, possibly due to remaining imperfections in the free electron density model. Spiral arm structure is required to reproduce the spatial distribution of pulsars. We have modelled this structure using fixed spirals, but there are apparent deficiencies. Proper modelling of the angular motion of the spirals with respect to the Sun may be needed for more accurate results.  Pulsar radio luminosities independent of the period and period derivative can safely be ruled out. They lead to the detection of too many old synthetic pulsars, as indicated by the exceedingly large observed scale height and a clear pile-up of detections on the death line. A model in which the radio (pseudo-)luminosity is, before dithering, proportional to $P^{-1.5} P^{0.5}$ (the square root of the spin-down luminosity) favors the detection of younger objects and appears consistent with the observations. In this model, the undithered luminosity is proportional to the voltage drop available for particle acceleration in magnetosphere models, which may be related to the physics of the radio emission mechanism. Also, the modelled fluxes are independent of the viewing geometry, provided that the pulsar beam crosses the line of sight, which supports the hypothesis of sharp-edged beams with random component locations.  We do not find evidence for significant torque decay (due to magnetic field decay or otherwise) over the lifetime of the pulsars as radio sources (\\mbox{$\\sim100$ Myr}). The conflicting conclusion of \\cite{2002ApJ...565..482G, 2004ApJ...604..775G}, who found evidence for magnetic field decay on time scales \\mbox{$\\lesssim5$ Myr}, is most likely due to different assumed (effective) dependences of the radio luminosity on the period and period derivative of the pulsar, which resulted in a pile-up of synthetic pulsars on the death line without simulated field decay. Our preferred luminosity model neatly avoids such a pile-up, as the undithered luminosity contours are parallel to the death line, so that the pulsars uniformly dim as they approach it. While it is unclear whether we have identified the correct radio pulsar luminosity law, this argues that conclusions regarding torque decay in isolated neutron stars based on simulation studies are strongly dependent on assumptions regarding the luminosities of the pulsars. We have demonstrated that it is possible to avoid the conclusion of magnetic field decay using a very simple and natural luminosity model.  Our method does not precisely constrain the distribution of birth spin periods of the pulsars, because the period of a pulsar is asymptotically independent of its initial value. However, in order to avoid an excess of simulated detections with periods \\mbox{$\\lesssim100$ ms}, many pulsars must be formed with greater periods. We found that a Gaussian distribution with mean \\mbox{300 ms} and standard deviation \\mbox{150 ms} accommodates the observations well. This suggests that characteristic ages, which assume that the initial period is negligible compared to the observed one, are generally poor age indicators, at least for young pulsars.  We estimate the Galactic pulsar birthrate to be $\\sim2.8$ pulsars per century. After accounting for a fraction of core-collapse supernovae forming black holes, this value is consistent with all neutron stars, except for a small fraction in X-ray binaries and magnetars, being born as radio pulsars, although not necessarily beamed toward us. Using the beaming model of \\cite{1998MNRAS.298..625T}, the Galaxy is estimated to contain $\\sim120,000$ potentially observable ordinary pulsars.  An important point to emphasize is that, as is also the case with previous work, we have merely provided evidence for consistency of the observational data with a particular scenario. The value of our study lies in the fact that our proposed model has, arguably, minimal complexity and does not involve any controversial component. We have also implemented the recent advances in pulsar astronomy in a generally self-consistent manner. That very similar studies differing principally by assumptions about poorly constrainted aspects -- such as the radio luminosity of pulsars - led to conflicting conclusions -- notably regarding torque decay - is strongly suggestive that pulsar population simulations require further independent input before they can be used to draw definitive conclusions. Perhaps the single most important breakthrough would be an independently verified luminosity model, as it is what ultimately determines the detectability of pulsars. Until then, Ockham's razor should be applied when considering the results of such studies. Meanwhile, the pulsar data set is poised to continue to be enhanced by the on-going Arecibo L-Band Feed Array (ALFA) pulsar survey, which is expected to detect as many as 1000 ordinary pulsars and 50 millisecond ones \\citep{fgk04, cfl+05}."
    },
    "0512/astro-ph0512066_arXiv.txt": {
        "abstract": "We show that in theories with a nontrivial kinetic term the contribution of the gravitational waves to the CMB fluctuations can be substantially larger than that is naively expected in simple inflationary models. This increase of the tensor-to-scalar perturbation ratio leads to a larger B-component of the CMB polarization, thus making the prospects for future detection much more promising. The other important consequence of the considered model is a higher energy scale of inflation and hence higher reheating temperature compared to a simple inflation. ",
        "introduction": "The main consequence of inflation is the generation of primordial cosmological perturbations \\cite{mc} and the production of long wavelength gravitational waves (tensor perturbations)~\\cite{star}. The predicted slightly red-tilted spectrum of the scalar perturbations is at present in excellent agreement with the measurements of the CMB fluctuations \\cite{cmb}% . The detection of a small deviation of the spectrum from flat together with the observation of primordial gravitational waves would make us completely confident in early-time cosmic acceleration. The detection of primordial gravitational waves is not easy, but they can be seen indirectly in the B-mode of the CMB polarization (see, for example, \\cite{book}). In standard slow-roll inflationary scenarios \\cite{Chaot} the amplitude of the tensor perturbations can, in principle, be large enough to be discovered. However, it is only on the border of detectability in future experiments.  There is no problem to modify the inflationary scenarios in a way to suppress the tensor component produced during inflation. In particular, in models such as new inflation~\\cite{New} and hybrid inflation~\\cite{Hybrid}, tensor perturbations are typically small~\\cite{book}. Moreover, in the curvaton scenario \\cite{curva} and k-inflation \\cite{k-Inflation}, they can be suppressed completely.  An interesting question is whether the gravitational waves can be significantly enhanced compared to our naive expectations. Recently it was argued that the contribution of tensor perturbations to the CMB anisotropy can be much greater than expected \\cite{Bartolo:2005jg,Sloth:2005yx}. However, it was found in \\cite{lms} that in the models considered in \\cite% {Bartolo:2005jg,Sloth:2005yx} one cannot avoid the production of too large scalar perturbations and therefore they are in contradiction with observations. It is also possible to produce blue-tilted spectrum of gravitational waves in a so-called superinflation \\cite{baldi}, which is, however, plagued by graceful exit problem.  At present there exists no inflationary model with graceful exit where the B-mode of polarization would exceed that predicted by simple chaotic inflation. The purpose of our paper is to show that such models can be easily constructed even within the class of simple slow-roll inflationary scenarios if we allow a nontrivial dependence of the Lagrangian on the kinetic term. These models resemble k-inflation with only difference that here inflation is due to the potential term in the Lagrangian. ",
        "conclusions": "We have shown above that in theories where the Lagrangian is a nontrivial, nonlinear function of the kinetic term, the scale of inflation can be pushed to a very high energies without coming into conflict with observations. As a result, the amount of produced gravitational waves can be much larger than is usually expected. If such a situation were realized in nature then the prospects for the future detection of the B-mode of CMB polarization are greatly improved. Of course, the theories where this happens are somewhat \\textquotedblleft fine-tuned\\textquotedblright . Namely, the corresponding Lagrangian $p\\left( X,\\phi \\right) $ must generically satisfy the condition, $p_{,X}+2Xp_{,XX}\\ll p_{,X},$ during inflation. For the particular model (\\ref{sm1}) this means that $12\\pi \\alpha ^{2}-m^{2}\\ll12\\pi \\alpha ^{2}$. This fine-tuning of the parameters of the theory should not be confused, however, with the fine-tuning of the initial conditions. In fact, the true theory of nature is unique and, for example, Lagrangian (\\ref{sm1}), where $12\\pi \\alpha ^{2}$ is not very different from $m^{2}$ looks even more attractive because it has fewer free parameters. Therefore, future observations of the CMB fluctuations are extremely important since they will restrict the number of possible candidates for the inflaton. The other peculiar properties of the generalized slow-roll inflation will be considered in \\cite{mv}."
    },
    "0512/astro-ph0512299_arXiv.txt": {
        "abstract": "{}{We report a detailed attempt to model the ortho-to-para abundance ratio of $c$-C$_3$H$_2$ so as to reproduce observed values in the cores of the well-known source TMC-1. According to observations, the ortho-to-para ratios vary, within large uncertainties, from a low of near unity to a high of approximately three depending on the core. } {We used the  osu.2003 network of gas-phase chemical reactions augmented by reactions that specifically consider the formation, depletion, and interconversion of the $ortho$ and $para$ forms of the $c$-C$_{3}$H$_{2}$ and its precursor ion $c$-C$_{3}$H$_{3}^{+}$. We investigated the sensitivity of the  calculated ortho-to-para ratio for $c$-C$_{3}$H$_{2}$ to a large number of factors. } {For the less evolved cores C, CP, and D, we had no difficulty reproducing the observed ortho-to-para ratios of 1-2. In order to reproduce observed ortho-to-para ratios of near three, observed for the evolved cores A and B, it was necessary to include rapid ion-catalyzed interconversion processes. } {} \\keywords {Astrochemistry -- ISM: molecular abundances -- ISM: clouds -- ISM: molecules -- ISM: Individual objects: TMC1}   \\authorrunning{Park, Wakelam, \\& Herbst} \\titlerunning{Modeling the $o/p$ C$_{3}$H$_{2}$ ratio in TMC-1} ",
        "introduction": "Cyclopropenylidene, $c$-C$_3$H$_2$, is a widely-distributed abundant organic ring molecule in the interstellar medium \\citep{Ref:c3h2-10,mat85,Ref:c3h2-12}. Its ``linear'' isomer, propadienylidene ($l$-C$_3$H$_2$), has also been observed in a number of sources, albeit at lower abundance \\citep{Ref:c3h2-7,Ref:c3h2-5,teyssier,Ref:c3h2-6,Ref:c3h2-66,Ref:c3h2-8}. For each isomer, two equivalent H nuclei, of spin 1/2, couple to generate $ortho$ (nuclear spin = 1) and $para$ (nuclear spin = 0) species with spin statistical weights of 3 and 1, respectively. The nuclear spin wave functions of the $ortho$ states are symmetric to exchange of the protons whereas those of the $para$ states are anti-symmetric. Since the Pauli exclusion principle requires total wave functions to be antisymmetric to exchange of protons, the symmetry of the rotational wave function of the molecule to exchange must be antisymmetric for $ortho$ spin states and symmetric for $para$ spin states if the ground electronic state is symmetric to exchange. The dependence of the exchange symmetry of the rotational state on the rotational quantum numbers depends on the structure of the molecule and on the symmetry of the electronic state \\citep{townes}. Moreover, the conversion of molecular species from rotational levels with $ortho$ spin functions to those with $para$ spin functions does not occur  efficiently if at all by radiative or non-reactive collisional mechanisms in the gas, so that, in the absence of reaction (or strong binding with a surface),  the sets of rotational levels with $ortho$ and $para$ spin functions can be regarded as distinct species.   The cyclic isomer of C$_{3}$H$_{2}$ has been observed in the interstellar medium in transitions belonging to both spin states:  for $o$-C$_3$H$_2$, the observed lines include the $b$-type transitions $1_{10}$-1$_{01}$, 2$_{12}$-1$_{01}$, 3$_{12}$-3$_{03}$, 3$_{21}$-3$_{12}$ and $3_{30}$-3$_{21}$ at 18.3, 85.3, 83.0, 44.1, and 27.1 GHz respectively, whereas for $p$-C$_3$H$_2$, the observed $b$-type transitions include 2$_{02}$-1$_{11}$, 2$_{11}$-2$_{02}$, 2$_{20}$-2$_{11}$ and 3$_{22}$-3$_{13}$ lines at 82.1, 46.8, 21.6, and 84.7 GHz respectively \\citep{Ref:c3h2-7,mat85, Ref:c3h2-10,Ref:opC3H2,Ref:BR-madden,morisawa}. With this many transitions, ortho-to-para ratios of reasonably high accuracy can be obtained, allowing a detailed understanding of the chemical processes involved and possibly even the history of the source.  In the terrestrial laboratory one often refers to two-types of mixtures for $ortho$ and $para$ species involving two protons: the common $normal$ variety, in which the ortho-to-para ratio is three, and the $equilibrium$ variety, in which the ortho-to-para ratio reflects thermal equilibrium.  In the interstellar medium, however, the ortho-to-para  ratio at low temperatures is probably governed by kinetic rather than thermodynamical considerations, including processes on grains, and is likely to be time-dependent and may even reflect the original circumstances of the molecule formation \\citep{tine,Ref:opH2-2,Ref:opH2-3}. In fact, detailed estimates of the steady-state interstellar ortho-to-para H$_2$ ratio at 10 K show that it is approximately 10$^{-3}$ \\citep{opH2-ratio,opH2-non-thermal,cooling-H2}, far higher than the equilibrium value of $3 \\times 10^{-7}$.  Molecular hydrogen is unique in showing such a strong deviation between thermal and normal values for the $o/p$ ratio at temperatures in the vicinity of 10 K. For heavier species, this deviation is much smaller because more levels of each spin modification are thermally populated so that values of $\\approx 3.0$ for the thermal ratio pertain to quite low temperatures. For H$_{2}$CO, \\citep{Ref:H2CO-1,Ref:H2CO-2},  this ratio is reached at temperatures as low as $\\approx$ 15 K, while for $c$-C$_{3}$H$_{2}$, the ratio is reached under thermal conditions at temperatures under 5 K \\citep{Ref:opC3H2}. Nevertheless, the actual $o/p$ ratio in cold dense interstellar sources is determined by kinetic considerations and may well be significantly under three at 10 K clouds for heavy species with two equivalent protons, as detected for the species H$_2$CO, H$_2$CS, and H$_2$CCO \\citep{ Ref:H2CO-1,minh,ohishi}.  Although the ortho-to-para abundance ratio of $c$-C$_3$H$_2$ had been studied observationally in TMC-1 by \\cite{Ref:BR-madden} and \\cite{Ref:opC3H2} and in L1527 by \\citet{Ref:opC3H2}, this work is motivated by the recent observational work of \\citet{morisawa} (hereafter, MFK), who measured the ratio in a more comprehensive manner for the six different cores  within the TMC-1 ridge: A, B, C, D,  CP, and E. Their measured ratios range from 1.4$\\pm$0.7 (TMC-1D and CP) to 3.0$\\pm$1.5 (TMC-1A and B). Note that Core D sometimes refers to some substructure near the CP condensation, in which case, it can be referred to as CP-b \\citep{TMC-1-CP-b}, but  other authors have used Core D to be the same as Core CP \\citep{ Ref:ism-3-2}, which is the very well-studied core where the abundances of complex carbon chains peak \\citep{smith04}. In addition, TMC1-B is the so-called ammonia peak \\citep{age-howe}.  The current understanding is that some of the cores (e.g. Core B) are more evolved than others (e.g. Core CP) although the cause of the differential evolution is not fully understood \\citep{Ref:TMC-models,age-howe,markwick}.  Table \\ref{op-obs} shows the observed $o/p$ abundance ratios of cyclic $c$-C$_{3}$H$_{2}$ in the six cores \\citepalias{morisawa} along with  measured fractional abundances (X) of $o$-C$_{3}$H$_{2}$ with respect to H$_{2}$  and estimated ages of the cores. Although the determination of cloud ages is still contentious, there is a sufficient consensus concerning the chemical indicators of age to use the results in this paper. It can be seen that, as noted by \\citetalias{morisawa}, the ortho-to-para ratio seems to be generally correlated with the evolution of the core; i.e., for all cores other than Core E,  the older the core chemically, the higher the ortho-to-para ratio within large uncertainties, especially for the higher ortho-to-para ratios. \\begin{table*} \\caption[]{Some c-C$_{3}$H$_{2}$ observations of the cores in the TMC-1 ridge} $$ \\begin{array}{p{0.225\\linewidth}p{0.225\\linewidth}p{0.225\\linewidth}p{0.225\\linewidth}} \\hline Source     &  Observed $o/p$ Ratio$^b$       \t\t& Observed X($o$-C$_3$H$_2$) $\\times\\ 10^{-10}$ $^{b}$  & Estimated Age (yr) \\\\\\noalign{\\smallskip} \\hline\\noalign{\\smallskip} {\\it Cores with high $o/p$ ratio\\ $\\cdots$}   \t\t& \t        &              \t& \\\\\\noalign{\\smallskip} TMC-1A      & 3.0$^{+2.2}_{-1.3}$    \t    \t\t& 7.5$^{+12}_{-5.4}$\t& $\\sim$ 10$^6$ $^b$ \\\\\\noalign{\\smallskip} TMC-1B      & 3.0$^{+1.7}_{-1.1}$; 3.0$\\pm$0.16$^{c}$   & 7.9$^{+12}_{-5.4}$\t& $\\sim$ 10$^6$-10$^7$ $^h$ \\\\\\noalign{\\smallskip} TMC-1E      & 2.8$^{+2.1}_{-1.2}$   \t \t    \t& 4.6$^{+7.7}_{-3.4}$\t& $\\sim$ 10$^4$-10$^5$ $^{a,b}$ \\\\\\noalign{\\smallskip} {\\it Cores with low $o/p$ ratio\\ $\\cdots$}    \t&       &              \t& \\\\\\noalign{\\smallskip} TMC-1C      & 1.8$^{+0.9}_{-0.6}$; 2.4$\\pm$0.1$^{c}$    & 6.0$^{+8.4}_{-4.0}$; 13 $^{e}$\t& $\\sim$ 10$^5$ $^b$ \\\\  \\noalign{\\smallskip} TMC-1CP     & 1.4$^{+0.9}_{-0.5}$; 2.56$\\pm$0.31$^{d}$  & 16$^{+32}_{-13}$; 57 $^{d}$ \t& $\\sim$ 10$^5$ $^{f,g}$; $6 \\times 10^4$ $^h$ \\\\  \\noalign{\\smallskip} TMC-1D$^a$  & 1.4$^{+1.1}_{-0.7}$                      \t& 8.3$^{+19}_{-6.9}$ \t\t \t\t& $6 \\times 10^4$ $^h$ \\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\label{op-obs} \\begin{list}{}{} \\item  $^a$ \\citet{TMC-1-CP-b}. Core D is Core CP-b according to this source. \\item  $^b$ \\citetalias{morisawa} unless noted. The fractional abundances X of $o$- C$_{3}$H$_{2}$ are with respect to H$_{2}$.  Based on the MFK logarithmic observables, we carried out slightly different error calculations, because here $\\Delta$ logX is not much less than 1, so that two 1$\\sigma$ random errors $\\Delta{X}_{-}\\napprox\\Delta{X}_{+}$   should be calculated \\citep{wakelam} . \\item  $^c$ \\citet{Ref:opC3H2} \\item  $^d$ \\cite{madphd} \\item  $^e$ \\citet{cox-c3h2} \\item  $^f$ \\citet{age-D} \\item  $^g$ \\citet{smith04} \\item  $^h$ \\citet{Ref:ism-3-2} \\end{list} \\end{table*}  \\citetalias{morisawa} used a simple model of a small number of gas-phase chemical reactions to understand their measured abundance ratios in terms of the evolution of the cores. They showed that the ortho-para chemistry of $c$-C$_{3}$H$_{2}$ is rather simple compared with other heavy molecules that have been studied, so that a complete understanding of the $o/p$ abundance ratio is possibly obtainable. Our detailed model is based on their simple one, with, as discussed below, the important reactions for the formation and depletion of the $ortho$ and $para$ forms of $c$-C$_{3}$H$_{2}$ and its protonated ion precursor embedded in a large chemical network.\\\\  In this paper, we report the results of our detailed gas-phase chemical model of the ortho-to-para abundance ratio of $c$-C$_{3}$H$_{2}$. The remainder of this paper is divided as follows. In Section 2 we discuss the details of our gas-phase chemical model, while in Section 3, we illustrate our results for the ortho-to-para ratio of $c$-C$_{3}$H$_{2}$ as a function of time. We also compare our results with the model of \\citetalias{morisawa}. As will be seen, our standard results are not capable of producing $o/p$ ratios greater than 2.1, so we consider some additional direct chemical processes that can convert the $para$ form of the ring molecule to its $ortho$ form. With these processes, we are able to account for the possibly high $o/p$ ratio in the more evolved cores. We then subject our results to an uncertainty analysis. Section 4, the discussion, concludes the paper. ",
        "conclusions": "New observations by \\citetalias{morisawa} of the ortho-to-para abundance ratio for the cyclic isomer of C$_{3}$H$_{2}$ in the six cores of TMC-1 indicate that this ratio is under 2.0 for two of three cores thought to be chemically young and is closer to the thermal value of 3.0 for those cores thought to be chemically more evolved.  The parameter used to determine chemical age is the abundance ratio of NH$_{3}$ to that of CCS.  The $o/p$ ratio is determined by a chemistry in which the cyclic ion C$_{3}$H$_{3}^{+}$ is formed initially solely in its {\\it para} state, but by a process of dissociative recombination and re-protonation, a significant ortho/para abundance ratio, near the thermal value of 3.0,  can be built up given sufficient  time for high-metal abundances if competitive reactions are not dominant.  Although a simple model by \\citetalias{morisawa} indicates agreement with their observations for five out of six cores if high-metal abundances are used, our more detailed calculations with an extended version of the {\\sc osu.2003} network of reactions do not give such a definitive result, mainly because of competitive processes.  In fact, with standard parameters varied over rather wide ranges, we find it difficult to obtain $o/p$ ratios exceeding 2.0 at any time.  We are able to calculate higher $o/p$ ratios for C$_{3}$H$_{2}$ only if we assume that the $ortho$ and $para$ forms of this species can be converted into one another by exchange processes involving H$^{+}$ and protonating ions HX$^{+}$.  This assumption is based on some experiments involving H$_{3}^{+}$ by  \\citet{cordon}   but here runs up against the problem that both H$^{+}$ and HX$^{+}$ will likely react with C$_{3}$H$_{2}$ and the competition between reaction and $ortho/para$ conversion may well favor the former.  Thus, there is still no definitive chemical proof that gas-phase processes can produce $ortho/para$ ratios for cyclic C$_{3}$H$_{2}$ near the thermal value of 3.0.  If, then, the observational values in this range are sufficiently accurate to distinguish correctly between low and high ratios, either the exchange reactions we consider between $o$-C$_{3}$H$_{2}$ and $p$-C$_{3}$H$_{2}$ do occur and are actually favored over reaction, our branching fractions or other aspects of the gas-phase chemistry are in error, or one must consider gas-grain interactions.  Future observations with a lower uncertainty will help to determine which if any of these choices is the correct one.  Even if the $o/p$ ratio is reproducible by pseudo-time-dependent models of the type considered here, we must realize that the fractional abundance of the cyclic species tells a different story."
    },
    "0512/astro-ph0512250_arXiv.txt": {
        "abstract": "We present multi-frequency, multi-epoch radio imaging of the complex radio source B2151+174 in the core of the cluster, Abell~2390 ($z\\simeq0.23$). From new and literature data we conclude that the FRII-powerful radio source is the combination of a compact, core-dominated `medium-symmetric object' (MSO) with a more extended, steeper spectrum mini-halo. B2151+174 is unusual in a number of important aspects: i) it is one of the most compact and flat spectrum sources in a cluster core known; ii) it shows a complex, compact twin-jet structure in a north-south orientation; iii) the orientation of the jets is 45\\degr misaligned with apparent structure (ionization cones and dust disk) of the host galaxy on larger scales. Since the twin-jet of the MSO has its northern half with an apparent `twist', it might be that precession of the central supermassive black hole explains this misalignement. B2151+174 may be an example of the early stage (10$^{3-4}$yrs duration) of a `bubble' being blown into the ICM where the plasma has yet to expand. ",
        "introduction": "The cores of clusters of galaxies host many of the most luminous radio sources in the local Universe (e.g. Cygnus-A, Hercules-A, Hydra-A). The influence of such powerful sources of radiation and mechanical energy on the surrounding intracluster medium is clear in the core of the Perseus cluster \\cite{Bohetal93,Fabetal03} and may be a significant energy source regulating cooling flows \\cite{Soker01,Fabetal02,BrugK02,Ruszet02}. The majority of this energy is believed to be injected into the ICM through mechanical work done by radio jets but relatively few cooling flows have luminous, lobed radio emission and most are relatively compact and low luminosity. This is as would be expected if these radio sources have a duty-cycle of the order of 10$^8$yr \\cite{dlVetal04} but in a large sample of clusters a few luminous, compact sources should be found that are in the process of launching jets. With a view to investigating the influence of more compact radio sources on the ICM, we have undertaken an extensive, high resolution radio imaging campaign of the radio source B2151+174 (\\pcite{Broetal98}) that lies in the central galaxy of the rich cluster Abell 2390. This cluster has a high X-ray luminosity ($L=2.125 \\times 10^{45}$ erg~s$^{-1}$  (0.1-2.4 keV) % --- \\pcite{Ebeetal96}) and is known to contain a cooling flow ($\\la 300$ M$_{\\odot}$yr$^{-1}$ \\cite{Alletal01} from recent Chandra observations). % The central host galaxy of B2151+174 has extended optical emission lines \\cite{LeBetal91}, extended Lyman-$\\alpha$ \\cite{HutBal00}, strong dust absorption in the optical and sub-mm dust continuum emission \\cite{Edgetal99}.  B2151+174 is one of the most luminous radio sources in a BCG (FRII 1.4~GHz power of $10^{25.1}$~W/Hz) in the combined BCS/eBCS X-ray flux-limited sample of 300 clusters \\cite{Ebeetal98,Ebeetal00} and samples of radio sources in cluster cores  \\cite{BalBurLok93} but the radio properties of B2151+174 differentiate it from most canonical central cluster radio sources. It has an apparently self-absorbed radio spectrum \\cite{Edgetal99} implying that the majority of the radio emission is from a very compact region.   The kpc-scale radio structure of B2151+174 was unveiled for the first time by Augusto et al. (1998), from VLBA  and MERLIN 5 GHz maps, who studied it as a candidate compact-symmetric object (CSO). In fact, in this paper, we revise the classification and find the source to be larger than 1~kpc, making it a medium-sized symmetric object (MSO). CSOs (sizes  on 1--1000~pc) are radio sources with lobes on either side of a compact core which dominates at high frequencies (e.g.\\ \\pcite{Wiletal94}).  MSOs are similar  but larger (1--15~kpc). CSOs,  short and `twin-sided', are likely young sources ($10^3$--$10^4$ yrs); the ages of some have been determined by either kinematical studies (e.g.\\ \\pcite{OwsConPol98}) or synchrotron emission aging  \\cite{Reaetal96a}.     In Section~2 we present the data collection and describe its processing, including the problems found. In Section~3 we use all the data to make an extensive study of the radio properties of B2151+174, from maps, spectra (total and main components), variability (or lack of it), jets orientation, size, mini-halo evidence, etc. Finally, in Section~4, we conclude the paper with an overall discussion.  Throughout this paper, we assume H$_{0}=75$ km s$^{-1}$ Mpc$^{-1}$, q$_{0}=0.5$ and $\\Lambda=0$. ",
        "conclusions": "The radio source associated with the cD in the rich, X-ray luminous cluster A2390 (B2151+174) is unusual in a number of important aspects. First, it is one of the most compact and flat spectrum sources in a cluster core known. Second, it shows a complex, compact twin-jet structure. Third, the orientation of the jets is not aligned with apparent structure of the host galaxy on larger scales.  These properties are intriguing given the recent realisation that the influence of an AGN in the centre of a cluster could have a very significant impact on the intracluster medium in the core and the suppression of cooling. B2151+174 may be an example of the early stage of a `bubble' being blown into the ICM where the plasma is yet to expand. This initial phase will be relatively short (10$^{3-4}$yrs) compared to the later stages of the expansion (10$^{6-7}$yrs) so few `young' sources can be expected in any selection of clusters.  It is important to note that the 1.4~GHz radio power of B2151+174 ($10^{25.1}$~W/Hz) is relatively small compared to the events that are thought to have an appreciable effect on the ICM in the core of a cluster \\cite{dlVetal04} but the probability of observing the early stages of a sufficiently powerful event is small.  In samples of clusters of galaxies at least 300 are needed before the meaningful statistics of these `young' sources can be assessed as the ``outburst'' period is a small fraction of the duration of each injection event. Alternatively, establishing what fraction of bright, compact Giga-Hertz peaked spectrum radio sources (the probable appearance of such sources) are in cluster cores would also provide important constraints on these issues.  Perhaps the most puzzling aspect of B2151+174 is the gross misalignment of the smallest scale radio jet with the larger scale `ionization cones' and dust disk seen in HST imaging. This may simply result from precession of the central supermassive black hole between two distinct episodes of activity or indicate a complex change in the angular momentum of gas accreted in the inner part of the disk in this galaxy. High spatial resolution radio imaging of other systems in which a clear large scale alignment is found will be important in addressing this issue."
    },
    "0512/astro-ph0512176_arXiv.txt": {
        "abstract": "We analyzed RXTE archival observations of 4U 1907+09 between 17 February 1996 and 6 March 2002. The pulse timing analysis showed that the source stayed at almost {\\bf{constant}} period around August 1998 and then started to spin-down at a rate of $(-1.887\\mp 0.042)\\times 10^{-14}$ Hz s$^-1$  which is $\\sim$ 0.60 times lower than the long term ($\\sim 15$ years) spin-down rate (Baykal et al. 2001). Our pulse frequency measurements for the first time resolved significant spin-down rate variations since the discovery of the source. We also presented orbital phase resolved X-ray spectra during two stable spin down episodes during November 1996 - December 1997 and March 2001 - March 2002. The source has been known to have two orbitally locked flares. We found that X-ray flux and spectral parameters except Hydrogen column density agreed with each other during the flares.We interpreted the similar values of X-ray fluxes as an indication of the fact that the source accretes not only via transient retrograde accretion disc (in't Zand et al. 1998) but also via the stellar wind of the companion (Roberts et al. 2001), so that the variation of the accretion rate from the disc does not cause significant variation in the observed X-ray flux. Lack of significant change in spectral parameters except Hydrogen column density was interpreted as a sign of the fact that the change in the spin-down rate of the source was not accompanied by a significant variation in the accretion geometry.   {\\bf{Keywords:}} accretion, accretion discs -- stars:neutron -- X-rays:binaries -- X-rays:individual 4U 1907+09 ",
        "introduction": "The X-ray source 4U 1907$+$09 is an accretion powered X-ray pulsar accreting plasma from a blue supergiant companion star. Since its discovery in the Uhuru survey (Giacconi et al. 1971), it has been observed by various X-ray observatories including  Ariel V (Marshall \\& Ricketts 1980), Tenma (Makishima et al. 1984), EXOSAT (Cook \\& Page 1987), Ginga (Makishima \\& Mihara 1992; Mihara 1995), BeppoSAX (Cusumano et al. 1998), XMPC (Chitnis et al. 1993), IXAE (Mukerjee at al.2001) and RXTE (in't Zand, Baykal \\& Strohmayer 1998a; in't Zand, Strohmayer \\& Baykal 1997, 1998b; Roberts et al. 2001 ;Baykal et al. 2001). Marshall \\& Ricketts (1980) determined the orbital period to be 8.38d using Ariel V observations. They also found two flares, namely a primary and a secondary {\\bf{occuring near periastron and apastron}}, and EXOSAT and RXTE observations (in't Zand et al. 1998a,b) have clearly shown that these flares are locked to orbital phases separated by half an orbital period. Using Tenma observations, Makishima et al. (1984) also found the pulse period of the source to be 437.5 sec. Makishima et al. (1984) and Cook \\& Page (1987) discussed the possibility that the two flares are caused by an equatorial disc-like envelope (circumstellar disc) around the companion star. When the neutron star crosses the disc, the mass accretion rate onto the neutron star increases temporarily. We expect to see two peaks while the neutron star passes twice through the circumstellar disc of the companion star in every cycle. We also expect a transient accretion disc formation around the neutron star while the neutron star passes through this circumstellar disc. For the formation of the circumstellar disc around the companion star, the companion must most likely be a Be type star. Recent optical observations showed that  the optical companion of 4U 1907$+$09 is an O8/O9 supergiant with a dense stellar wind (Cox, Kaper, Mokiem 2005). Formation of a transient accretion disc around the neutron star as a result of disruption of the dense stellar wind of a massive companion is quite possible as well (Blondin et al. 1990).  The timing measurements from EXOSAT (Cook \\& Page 1987) and RXTE observations (in't Zand et al. 1998a,b; Baykal et al. 2001) showed continuous spin down of the neutron star.The source was found to be spinning down almost at a rate of ${\\dot \\nu}=(-3.54 \\pm 0.02)\\times {10^{-14}} {Hz s^{-1}}$ for more than 15 years (Baykal et al. 2001). Timing analysis of RXTE (in't Zand et al. 1998) and IXAE (Mukerjee et al. 2001) observations of 4U 1907-09 also revealed transient oscillations with periods  of about 18.2s and 14.4s respectively. These transient oscillations {\\bf{were}} suggested to be related to the presence of a transient retrograde accretion disc which is also responsible for the slowing down of the pulsar.  In this paper we present the results of timing and spectral studies of archival RXTE-PCA observations of 4U 1907+09,  and report a change in the spin down rate of this system. ",
        "conclusions": "The spin-down rate of 4U 1907+09 obtained from the pulse arrival times between March 2001 and March 2002 was found to be $\\sim$ 0.60 times lower than both the previous RXTE measurements of spin-down rate between November 1996 - December 1997 by Baykal et al., (2001) and the long term spin down trend of the source between 1983 and 1997 (See Figure 1 in Baykal et al. 2001).  In this and the previous RXTE observations, time span of the observations were similar ($\\sim $ over year) and in both observations posses low timing noise. However spin down rate is lowered by a factor of 0.6  in latter observations. It is also interesting to observe that the spin-down rate is consistent with zero around MJD 51000 before this significant change in spin-down rate. Our pulse frequency measurements for the first time resolved significant spin-down rate variations since the discovery of the source.  The steady spin-down rate (Baykal et al. 2001) and the presence of occasional transient oscillations (in't Zand et al. 1998a; Mukerjee et al. 2001) in 4U 1907+09 support the idea that the source accretes from retrograde transient accretion disc. The possibility of a prograde disc for which the magnetospheric radius should be close to the corotation radius so that the magnetic torque overcomes material torque is not likely to be the case for this system (in't Zand et al. 1998a). If the disc is retrograde and the material torque is dominant, using Ghosh\\& Lamb (1979) accretion disc model, a decrease in spin-down rate of the neutron star should be a sign of a decrease in the mass accretion rate coming from the accretion disc. If the disc accretion is the only accretion mechanism, this decrease in mass accretion rate should also lead to a decrease in X-ray flux of the system.  From the X-ray spectral analysis of the source, we found no clear evidence of a correlation between 3-25 keV flux and the change in spin-down rate (see Figure 3). The flux levels were about the same for latter observations with low spin-down rate compared to the rest of the observations. This was not our expectation if we only consider disc accretion and may be because of the fact that the total mass accretion rate is not only due to only disc accretion, but also accretion from stellar wind may contribute to the total mass accretion rate (see also Roberts et al. 2001) . In case of accretion from both transient disc and wind, the change in mass accretion rate from the accretion disc might not cause a substantial change in X-ray flux. However, it is important to note that the orbital coverage of the latter (proposal ID 60061) observations is poor and limited to the phase locked flares in orbital phases $\\sim 0.05$ and $\\sim 0.6$ as seen in Figure 3. Future observations of the source may be helpful to have a better understanding of a possible relation between spin-down rate and X-ray flux.  Examining spectral parameters plotted in Figure 3 especially for flare parts, we also found an evidence of an increase in Hydrogen column density ($n_H$) for the latter observation with lower spin-down rate. On the other hand, there is no clear evidence of a variation of the other spectral parameters in the flares after the spin-down rate decreased. Change in $n_H$ may be related to a change in the accretion geometry of the source. However, if there had been an accretion geometry change accompanied with the change in the spin-down rate of the source, we would expect other spectral parameters to vary as well, especially the power law index as in the RXTE observations of SAX J2103.5+4545 (Baykal et al. 2002), and 2S 1417-62 (Inam et al. 2004). New X-ray observations of the source will be useful to monitor any possible changes in the spin-down rate and its consequences.  {\\bf"
    },
    "0512/astro-ph0512495_arXiv.txt": {
        "abstract": "Relativistic neutron-loaded outflows in gamma-ray bursts are studied at their early stages, before deceleration by a surrounding medium. The outflow has four components: radiation, electrons, protons and neutrons. The components interact with each other and exchange energy as the outflow expands. The presence of neutrons significantly changes the outflow evolution. Before neutrons decouple from protons, friction between the two components increases their temperatures by many orders of magnitude. After the decoupling, the gradual neutron decay inside the outflow has a drag effect on the protons and reduces their final Lorentz factor. ",
        "introduction": "\\label{sec:introc4}  A neutron component in $\\gamma$-ray bursts (GRBs) was proposed by Derishev, Kocharovsky \\& Kocharovsky (1999a,b), and detailed calculations of nuclear composition show that free neutrons are inevitably present among ejected baryons (Beloborodov 2003b; hereafter B03b). Any plausible central engine of GRBs is dense and at least mildly degenerate, which leads to its neutronization. During the explosion, the expanding neutron-rich material may undergo nucleosynthesis: neutrons tend to recombine with protons to $\\alpha$-particles (Lemoine 2002; Pruet, Guiles \\& Fuller 2002; B03b). This recombination may be successful if the outflow is collimated. However, even in an extreme case of a very efficient recombination, a significant neutron component is left over in neutron-rich outflows because the formation of $\\alpha$-particles consumes equal numbers of neutrons and protons. The abundance of leftover neutrons may vary from 10\\% to more than 90\\% depending on the precise parameters of the burst.  The presence of neutrons changes the theoretical picture of GRB explosion. Firstly, they may develop a somewhat smaller Lorentz factor than protons. When such a decoupling takes place, the last $n$-$p$ collisions lead to emission of observable multi-GeV neutrinos (Derishev et al. 1999a, hereafter DKK99a; Bahcall \\& \\Mesz~2000; \\Mesz~ \\& Rees 2000a). Secondly, neutrons decay with time. The decay impacts the external blast wave at radii $r\\sim 10^{16}-10^{17}$~cm because even an exponentially small number of survived neutrons carry an energy much larger than the rest-energy of external medium (Beloborodov 2003a).  In the present paper, we study the dynamics of neutron-loaded outflows at early stages of their expansion, $r<10^{16}$~cm, before they are decelerated by an external medium. The neutrons decay gradually at all radii $r\\lsim 10^{17}$~cm, and at small $r$ the decay occurs {\\it inside} the GRB outflow. The decay turns out important at $r$ as small as $10^{12}$~cm.  Interesting effects also take place at small $r<10^{12}$~cm. Neutrons are initially accelerated together with protons because they are collisionally coupled, and the last $n$-$p$ collisions before decoupling cause a significant heating.  We assume in this paper a simple hydrodynamic picture of expansion driven by thermal pressure and study the basic dynamics of a uniform neutron-loaded outflow. We do not consider internal shocks or possible dynamical effects of magnetic fields. The paper is organized as follows. In section~\\ref{sec:pfire} we briefly review neutron-free outflows, which have been studied previously in detail (see Piran 2004 for a review). In section~\\ref{sec:npfire} we derive equations describing neutron-loaded outflows and calculate example numerical models. Results are discussed in section~\\ref{sec:discc4}. ",
        "conclusions": "\\label{sec:discc4}  In this paper we developed the theory of relativistic neutron-loaded outflows. The presence a neutron component significantly affects the early dynamics of GRB explosions. In particular, the plasma temperature is increased by many orders of magnitude, and this heating can compete with other heating mechanisms such as internal shocks or dissipation of magnetic fields in the outflow.  The effects of neutrons are pronounced in outflows with $\\eta>\\eta_*$ given by equation~(\\ref{eq:etastar}), whose typical value is several hundred. Then neutrons and protons develop a substantial relative Lorentz factor which leads to strong heating and momentum exchange when neutrons decay. We also note that neutron-rich outflows with very high $\\eta\\gg\\eta_*$ lose most of their energy at the photosphere: their thermal radiation carries most of the energy at the moment of transparency. Therefore, the neutron effects may be especially strong in GRBs with a significant thermal component in the radiation spectrum (a number of such GRBs have been identified recently, see e.g. Ghirlanda 2003; Ryde 2004).  The calculations in this paper focused on the simplest model of a uniform hydrodynamic outflow with weak magnetic fields. We did not consider, for instance, internal shocks (e.g. Rees \\& M\\'esz\\'aros 1994) and possible pair creation by nonthermal $\\gamma$-rays generated in the outflow. Pair creation may extend the optically thick stage of expansion, and the trapped radiation may convert its energy more efficiently into bulk kinetic energy of the plasma. A large-scale magnetic field may gradually collimate the outflow so that the conical geometry of expansion does not apply (Vlahakis, Peng \\& K\\\"onigl 2003).  The $\\beta$-decay is likely to affect the development of internal shocks in the outflow. The shocks are caused by a non-uniform profile of the Lorentz factor, and the drag effect of decayed neutrons tends to smoothen this profile. The fastest portions of the outflow are more effectively decelerated and the initial contrast of Lorentz factors may be substantially reduced already at $r\\sim 10^{14}$~cm. This effect constrains the dissipation efficiency of the shocks, which is sensitive to the contrast of Lorentz factors (see Fig.~3 in Beloborodov 2000). In addition, the high temperature of the outflow heated by $\\beta$-decay may prevent development of the shocks. We defer a detailed study of these effects to a future work.  The impact of neutrons on the prompt burst and its afterglow provides a unique opportunity to link the observed emission with physical conditions in the central engine of the explosion. The very presence of neutrons is a signature of an extremely hot and dense engine. Observable effects of neutrons may shed light on the mechanism of GRB trigger."
    },
    "0512/astro-ph0512389_arXiv.txt": {
        "abstract": "Modeling the multiwavelength emission of successive regions in the jet of the quasar PKS 1136--135 we find indication that the jet suffers deceleration near its end on a (deprojected) scale of $\\sim$400 kpc. Adopting a continuous flow approximation we discuss the possibility that the inferred deceleration from a Lorentz factor $\\Gamma =6.5$ to $\\Gamma=2.5$ is induced by entrainment of external gas. Some consequences of this scenario are discussed. ",
        "introduction": "Despite decades of intense efforts, the present knowledge of the physical processes acting behind the phenomenology shown by relativistic jets is still rather poor. Basic questions about matter content, transported power, the role of the magnetic field, the dissipation mechanisms are awaiting answers (e.g. Blandford 2001). Among these problems, one of the most fundamental concerns the speed of the flow. While it is widely accepted that the rapid variability and the extreme compactness displayed by blazars (e.g. Maraschi \\& Tavecchio 2005; Sikora \\& Madejski 2001) and the observation of VLBI superluminal components (e.g. Kellermann et al. 2004 ) imply relativistic bulk flows (with $\\Gamma \\sim 10-20$) for the innermost portions of the jet ($0.1$ pc $ < d < 100$ pc), it is not clear whether the large scale jets still have relativistic speeds and how/where deceleration takes place. The present evidence suggests that FRI jets become trans-relativistic quite early, within few kiloparsecs (e.g. Bridle \\& Perley 1984), while the situation for FRII jets appears more ambiguous. Mildly relativistic speeds ($\\beta < 0.95$) are suggested by studies of high luminosity, lobe dominated sources based on the jet to lobe prominence (Wardle \\& Aaron 1997), while the interpretation of multiwavelength observations of extended jets in QSOs points toward highly relativistic speeds ($\\Gamma \\sim 10$) even at very large scales ($\\sim $ 100 kpc; e.g., Tavecchio et al. 2000, Celotti et al. 2001, Sambruna et al. 2002,2004, Siemiginowska et al. 2002, Marshall et al. 2005).  Deceleration of FRI jets is commonly believed to be the result of entrainment of external gas by the flow (e.g. Komissarov 1994, Bicknell 1994, B94 hereafter). In this framework it is conceivable that FRII jets, characterized by larger powers, are not substantially perturbed by entrainment: the fundamental division between FRI and FRII jets could basically depend on the difference in the mass/energy flux of the two types of jets (Bicknell 1995, Bowman et al. 1996). An important unknown parameter is the efficiency with which the jet can entrain the external gas. Efficient entrainment can be assured by the onset of fully-developed turbulence at transonic speeds (e.g. B94). In this scheme, supersonic flows are thought to be less efficient at collecting material and entrainment could affect only the external regions of the jet, resulting in the formation of a slow layer surrounding a fast, unperturbed ``spine'' (e.g. de Young 1986). Nevertheless, even low efficiency entrainment could become important if it is maintained over very long scales.  The possibility of ``observing''  the gradual slowing-down of a jet could in principle provide precious information on the physical processes at work. This approach was successfully developed in great detail for a few FRI jets where the morphology could be well studied thanks to the large angular scale (Laing \\& Bridle 2002, 2004).  For FR II jets the stronger asymmetry and smaller angular scale has generally prevented comparable studies.  Of particular interest are therefore the cases of FRII jets hosted by intermediate-$z$ QSOs in which multiwavelength observations have recently shown that the radio to X-ray flux ratio increases dramatically along the jet. This feature, interpreted in the framework of the synchrotron-IC/CMB emission model, suggests that the jet undergoes deceleration. The possible role played by deceleration was first discussed for the jet of 3C273 (Sambruna et al. 2001). Other cases of jets showing an increasing radio to X-ray flux ratio were presented in the survey discussed in Sambruna et al. (2002) and Sambruna et al. (2004). Georganopoulos \\& Kazanas (2004; hereafter GK2004) proposed an analytic framework to describe the effects of deceleration on the physical parameters of the jet on the basis of an adiabatic assumption, remarking that, if the jet decelerates, the magnetic field as well as the particle energy density will increase due to compression causing a large increase of the radio synchrotron emission, while the X-ray flux will decrease because of a reduced beaming of the CMB radiation field. In Sambruna et al. (2005; Paper I in the following) we report the results of deep {\\it Chandra} and multiwavelenth observations of PKS~1136-135, which allowed us to image in detail the emission along the jet.  Modelling the observed multifrequency emission we derived the profiles of the basic physical quantities along the jet, which support the deceleration scenario discussed above. In this paper on the basis of the observational evidence collected for the jet of 1136-135, already presented in Paper I, we discuss in more depth deceleration in terms of physical models and in particular entrainment of external gas.  The plan of the paper is the following: in Sect.2 we derive the profile of the relevant physical quantities, suggesting deceleration of the flow. In Sect.3 we discuss the physical origin of the deceleration. In Sect.4 we discuss in detail the possibility that deceleration is induced by the loading of the jet through entrainment. Finally in Sect.5 we conclude. Throughout this work we use the current ``concordance'' cosmological parameters: $\\rm H_0=71\\; Km\\; s^{-1}\\; Mpc^{-1} $, $\\Omega _{\\Lambda}=0.73$, $\\Omega _m = 0.27$ (Bennett et al. 2003). ",
        "conclusions": "The proposed scenario can explain in a plausible way the deceleration inferred for the jet of 1136-135 (and possibly in the other sources showing the same radio-to-X-rays increasing trend, GK2004).  A more detailed understanding of the processes at work and the comparison with the observed properties of jets necessarily involves several still poorly-known physical issues.  It is worth remarking that the evidence for deceleration in the jet of 1136-135 is based on the observed run of radio and X-ray fluxes along the jet and on the use of a specific emission model to interpret the multiwavelength observations (for a discussion of alternative emission models see Paper I). In this respect, the conclusion about deceleration is not completely model-independent. Other specific assumptions made in the modelling, such as the equipartion condition, affect (although in a minor way) the conclusion. Moreover, apart from the application of a specific emission model, one has to bear in mind that the result is based on the hypothesis that the emission comes from a not-structured flow. At least part of the increase of the radio flux close to the terminal part of the jet could be due to the contribution from another component (e.g. the plasma backflowing after the compression at the hot-spot).  The approach that we used to study the deceleration is based on a pure hydrodynamical treatment.  The contribution of all the components of the flow (protons, electrons, magnetic field) is included in just one parameter, the total pressure. Likely, much of the energy dissipated during the deceleration will be stored in the heated proton component. To follow the evolution of the other components (thermal and non-thermal electrons, magnetic field) we should specify the actual mechanisms coupling each component to the others. The determination of the coupling between protons, magnetic field and electrons, determining the fraction of dissipated energy that goes into the magnetic and non-thermal electron components, is a longstanding problem in astrophysics, clearly beyond the scope of this work. The derived profiles are in agreement with the possibility that the increase of the magnetic field and of the density of the relativistic electrons just follows the adiabatic heating induced by the deceleration. However, given the large uncertainties, this evidence is not conclusive. The contribution of non-thermal electrons and magnetic field to the total pressure is small, around 10\\%. Therefore, a large portion of the dissipated kinetic energy resides in the hot protons and is not radiated away.  Another interesting point is related to the nature of the mechanism able to induce the assumed entrainment. It is widely assumed that turbulence can effectively work only in slow (transonic or subsonic) jets.  Apart from fully developed turbulence, it is also possible that entrainment is mediated by viscosity effects induced by small-scale turbulence.  It is worth noting that in the calculation we have presented it is supposed that the entire jet is decelerated under the effect of the entrainment. More realistically, these processes will induce a radial structure in the jet, with the development of a sheared flow (e.g. De Young 1996).  In the simple scheme presented here, the jet starts to be decelerated when the amount of collected gas is of the order of 1/$\\Gamma $ of the mass transported by the jet. In this framework, jets characterized by different mass fluxes will experience different behaviours. Large mass fluxes will assure that, under the same conditions of external gas density and entrainment rate, the jet will reach its hotspot almost unperturbed. On the other hand, jets with a small mass flux will be decelerated soon. It is tempting to further speculate along these lines, connecting the deceleration of the jet to other properties of the central AGN. To this aim we indicate the jet mass flux as $\\dot{M}_{\\rm out}$, and we note that $\\dot{M}_{\\rm out}$ is the mass output rate of the central engine, which is probably some fraction of the accretion mass rate $\\dot{M}_{\\rm out}\\propto \\dot{M}_{\\rm acc}$ (e.g. Ghisellini \\& Celotti 2002b). From the condition expressed by Eq.(\\ref{fent}) and the definitions above we can find that the deceleration length $x_1$ is $x_1 \\propto \\dot{M}_{\\rm acc}/r_o$, (the exact dependence is related to the jet and the external density profile) where we have supposed that all the jet, before the deceleration, have a similar bulk Lorentz factor. If the initial radius of the jet scales as the gravitational radius of the central black hole $r_o \\propto M_{BH}$, we finally find $x_1 \\propto \\dot{m} _{\\rm acc}$, where $\\dot{m} _{\\rm acc}$ is the accretion rate in Eddington units. Thus it is possible to identify sources for which the jet is effectively decelerated within short scales with sources with low $\\dot{m} _{\\rm acc}$ and {\\it viceversa}, a situation reminiscent of the FRI/FRII division."
    },
    "0512/astro-ph0512030_arXiv.txt": {
        "abstract": "We discuss the design and implementation of HYDRA\\_OMP a parallel implementation of the Smoothed Particle Hydrodynamics--Adaptive \\p3m (SPH-A\\p3m) code HYDRA. The code is designed primarily for conducting cosmological hydrodynamic simulations and is written in Fortran77+OpenMP. A number of optimizations for RISC processors and SMP-NUMA architectures have been implemented, the most important optimization being hierarchical reordering of particles within chaining cells, which greatly improves data locality thereby removing the cache misses typically associated with linked lists. Parallel scaling is good, with a minimum parallel scaling of 73\\% achieved on 32 nodes for a variety of modern SMP architectures. We give performance data in terms of the number of particle updates per second, which is a more useful performance metric than raw MFlops. A basic version of the code will be made available to the community in the near future.\\\\ ",
        "introduction": "The growth of cosmological structure in the Universe is determined primarily by (Newtonian) gravitational forces. Unlike the electrostatic force, which can be both attractive and repulsive and for which shielding is important, the ubiquitous attraction of the gravitational force leads to extremely dense structures, relative to the average density in the Universe. Galaxies, for example, are typically $10^4$ times more dense than their surrounding environment, and substructure within them can be orders of magnitude more dense.  Modelling such large density contrasts is difficult with fixed grid methods and, consequently, particle-based solvers are an indispensable tool for conducting simulations of the growth of cosmological structure. The Lagrangian nature of particle codes makes them inherently adaptive without requiring the complexity associated with adaptive Eulerian methods. The Lagrangian Smoothed Particle Hydrodynamics (SPH,\\cite{GM77}) method also integrates well with gravitational solvers using particles, and because of its simplicity, robustness and ability to easily model complex geometries, has become widely used in cosmology. Further, the necessity to model systems in which orbit crossing, or phase wrapping, occurs (either in collisionless fluids or in collisional systems) demands a fully Lagrangian method that tracks mass. While full six-dimensional (Boltzmann) phase-space models have been attempted, the resolution is still severely limited on current computers for most applications.  Particle solvers of interest in cosmology can broadly be divided into hybrid direct plus grid-based solvers such as Particle-Particle, Particle-Mesh methods (\\p3m,\\cite{He81}) and ``Tree'' methods which use truncated low order multipole expansions to evaluate the force from distant particles \\cite{BH86}. Full multipole methods \\cite{GR87}, are slowly gaining popularity but have yet to gain widespread acceptance in the cosmological simulation community. There are also a number of hybrid tree plus particle-mesh methods in which an efficient grid-based solver is used for long-range gravitational interactions with sub-grid forces being computed using a tree. Special purpose hardware \\cite{Su90} has rendered the direct PP method competitive in small simulations (fewer than 16 million particles), but it remains unlikely that it will ever be competitive for larger simulations.  The \\p3m algorithm has been utilized extensively in cosmology. The first high resolution simulations of structure formation were conducted by Efstathiou \\& Eastwood \\cite{EE81} using a modified \\p3m plasma code. In 1998 the Virgo Consortium used a \\p3m code to conduct the first billion particle simulation of cosmological structure formation \\cite{Ev02}. The well-known  problem of slow-down under heavy particle clustering, due to a rapid rise in the number of short-range interactions, can be largely solved by the use of adaptive, hierarchical, sub-grids \\cite{C91}. Only when a regime is approached where multiple time steps are beneficial does the adaptive \\p3m (\\ap3m) algorithm become less competitive than modern tree-based solvers. Further, we note that a straightforward multiple time-step scheme has been implemented in \\ap3m with a factor of 3 speed-up reported \\cite{DE93}.  \\p3m has also been vectorized by a number of groups including Summers \\cite{SU93}. Shortly after, both Ferrell \\& Bertschinger \\cite{FB94} and Theuns \\cite{T94} adapted \\p3m to the massively parallel architecture of the Connection Machine. This early work highlighted the need for careful examination of the parallelization strategy because of the load imbalance that can result in gravitational simulations as particle clustering develops. Parallel versions of \\p3m that use a 1-dimensional domain decomposition, such as the P4M code of Brieu \\& Evrard \\cite{BE00} develop large load imbalances under clustering rendering them useful only for very homogeneous simulations. Development of vectorized treecodes \\cite{H90,HK89} predates the early work on \\p3m codes and a discussion of a combined TREE+SPH (TREESPH) code for massively parallel architectures is presented by Dav\\'{e} \\etal\\, \\cite{D97}. There are now a number of combined parallel TREE+SPH solvers \\cite{W03,K03,V01,L00} and TREE gravity solvers \\cite{MM00,UA01,D96}. Pearce \\& Couchman \\cite{PC97} have discussed the parallelization of \\ap3m+SPH on the Cray T3D using Cray Adaptive Fortran (CRAFT), which is a directive-based parallel programming methodology. This code was developed from the serial HYDRA algorithm \\cite{CT95} and much of our discussion in this paper draws from this first parallelization of \\ap3m+SPH.  A highly efficient distributed memory parallel implementation of \\p3m using the Cray SHMEM library has been developed by MacFarland \\etal\\, \\cite{M98}, and further developments of this code include a translation to MPI-2, the addition of \\ap3m subroutines and the inclusion of an SPH solver \\cite{T03}. Treecodes have also been combined with grid methods to form the Tree-Particle-Mesh solver \\cite{Jim,BO00,W02,B02,D03,VS05}. The algorithm is somewhat less efficient than \\ap3m in a fixed time-step regime, but its simplicity offers advantages when multiple time-steps are considered \\cite{VS05}. Another interesting, and highly efficient N-body algorithm is the Adaptive Refinement Tree (ART) method \\cite{KK97} which uses a short-range force correction that is calculated via a multi-grid solver on refined meshes.  There are a number of factors in cosmology that drive researchers towards parallel computing. These factors can be divided into the desire to simulate with the highest possible resolution, and hence particle number, and also the need to complete simulations in the shortest possible time frame to enable rapid progress. The desire for high resolution comes from two areas. Firstly, simultaneously simulating the growth of structure on the largest and smallest cosmological scales requires enormous mass resolution (the ratio of mass scales between a supercluster and the substructure in a galaxy is $>10^9$). This problem is fundamentally related to the fact that in the currently favoured Cold Dark Matter \\cite{B84} cosmology structure grows in a hierarchical manner. A secondary desire for high resolution comes from simulations that are performed to make statistical predictions. To ensure the lowest possible sample variance the largest possible simulation volume is desired.  For complex codes, typically containing tens of thousands of lines, the effort in developing a code for distributed-memory machines, using an API such as MPI \\cite{MPI}, can be enormous. The complexity within such codes arises from the subtle communication patterns that are disguised in serial implementations. Indeed, as has been observed by the authors, development of an efficient communication strategy for a distributed memory version of the \\p3m code has required substantially more code than the \\p3m algorithm itself (see \\cite{M98}). This is primarily because hybrid, or multi-part solvers, of which \\p3m is a classic example, have data structures that require significantly different data topologies for optimal load balance at different stages of the solution cycle. Clearly a globally addressable work space renders parallelization a far simpler task in such situations. It is also worth noting that due to time-step constraints and the scaling of the algorithm with the number of particles, doubling the linear resolution along an axis of a simulation increases the computational work load by a factor larger than 20; further doubling would lead to a workload in excess of 400 times greater.  The above considerations lead to the following observation: modern SMP servers with their shared memory design and superb performance characteristics are an excellent tool for conducting simulations requiring significantly more computational power than that available from a workstation. Although such servers can never compete with massively parallel machines for the largest simulations, their ease of use and programming renders them highly productive computing environments. The OpenMP (http://www.openmp.org) API for shared-memory programming is simple to use and enables loop level parallelism by the insertion of pragmas within the source code. Other than their limited expansion capacity, the strongest argument against purchasing an SMP server remains hardware cost. However, there is a trade-off between science accomplishment and development time that must be considered above hardware costs alone.  Typically, programming a Beowulf-style cluster for challenging codes takes far longer and requires a significantly greater monetary and personnel investment on a project-by-project basis.  Conversely, for problems that can be efficiently and quickly parallelized on a distributed memory architecture, SMP servers are not cost effective. The bottom line remains that individual research groups must decide which platform is most appropriate.  The code that we discuss in this paper neatly fills the niche between workstation computations and massively parallel simulations. There is also a class of simulation problems in cosmology that have particularly poor parallel scaling, regardless of the simulation algorithm used (the fiducial example is the modelling of single galaxies, see \\cite{TC00}). This class of problems corresponds to particularly inhomogeneous particle distributions that develop a large disparity in particle-update timescales (some particles may be in extremely dense regions, while others may be in very low density regions). Only a very small number of particles\\-insufficient to be distributed effectively across multiple nodes\\-will require a large number of updates due to their small time-steps. For this type of simulation the practical limit of scalability appears to be order 10 PEs.  The layout of the paper is as follows: in section 2 we review the physical system being studied. This is followed by an extensive exposition of the \\p3m algorithm and the improvements that yield the \\ap3m algorithm.  The primary purpose of this section is to discuss some subtleties that directly impact our parallelization strategy. At the same time we also discuss the SPH method and highlight the similarities between the two algorithms. Section 2 concludes with a discussion of the serial HYDRA code. Section 3 begins with a short discussion of the memory hierarchy in RISC (Reduced Instruction Set Computer) systems, and how eliminating cache-misses and ensuring good cache reuse ensures optimal performance on these machines. This is followed by a discussion of a number of code optimizations for RISC CPUs that also lead to performance improvements on shared memory parallel machines (primarily due to increased data locality). In particular we discuss improvements in particle bookkeeping, such as particle index reordering. While particle reordering might be considered an expensive operation, since it involves a global sort, it actually dramatically improves run time because of bottlenecks in the memory hierarchy of RISC systems.  In section 4 we discuss in detail the parallelization strategies adopted in HYDRA\\_OMP. To help provide further understanding we compare the serial and parallel call trees. In section 5 we consolidate material from sections 3 \\& 4 by discussing considerations for NUMA machines and in particular the issue of data placement.  Performance figures are given in Section 6, and we present our conclusions in section 7. ",
        "conclusions": "Conducting high resolution simulations of cosmological structure formation necessitates the use of parallel computing. Although distributed architectures provide an abundance of cheap computing power, the programming model for distributed systems is fundamentally complex. Shared memory simplifies parallel programming greatly since the shared address space means that only the calculation itself need be distributed across nodes. In this paper we have discussed a code for parallel shared memory computers that exhibits only marginally higher complexity than a serial version of the code and which also exhibits excellent performance. Additional constructs for parallel execution introduce only a small (10\\%) penalty for running on 1 node compared to the serial code.  Although the code does have some problems with regards load balancing, in particular a deficit in performance occurs when a refinement is too large to be calculated as part of the task farm but is not large enough to be efficient across the whole machine, these situations are comparatively rare. The poor scaling of SPH under heavy clustering is the most significant cause of load imbalance. In particular, if the heavy calculational load is confined to one refinement that is part of the task farm all threads will block until this refinement is completed. The most satisfactory solution to this problem is to substitute an alternative algorithm for the SPH in high density regions. We will present details of an algorithm that improves the SPH cycle time for high density regions elsewhere (Thacker \\etal in prep).  Most of the performance limitations can be traced to applying a grid code in a realm where it is not suitable. As has been emphasized before, treecodes are particularly versatile, and can be applied to almost any particle distribution. However, for periodic simulations they become inefficient since Ewald's method must used to calculate periodic forces. FFT-based grid methods calculate the periodic force implicitly, and exhibit particularly high performance for homogeneous particle distributions under light to medium clustering. Highly clustered (or highly inhomogeneous) particle distributions are naturally tailored to the multi-timestepping capability of treecodes. Although we see scope for introducing a multi-time stepping version of \\ap3m where sub-grids are advanced in different time step bins it is unclear in details what efficiencies could be gained. There are clearly parts of the algorithm, such as mass assignment, that are unavoidably subject to load imbalances. We expect that since the global grid update would be required infrequently the global integrator can still be made efficient. An efficient implementation of multiple time-steps is the last area where an order of magnitude improvement in simulation time can be expected for this class of algorithm.  In terms of raw performance, the code speed is high relative to the values given by Dubinski et al. On the GS1280 the full solution time for the unclustered distribution even exceeds that of the PM solution quoted for GOTPM on 64 processors. \\ap3m has been criticized previously for exhibiting a cycle time that fluctuates depending upon the underlying level of clustering. The data we have presented here shows the range in speeds is comparatively small (a factor of 4). We would also argue that since the cost of the short range correction is so small at early times, this criticism is misplaced. While recent implementations of Tree-PM have an approximately constant cycle time irrespective of clustering, the large search radius used in the tree correction leads to the tree part of the algorithm dominating execution time for all stages of the simulation. Conversely, only at the end of the simulation is this true for HYDRA.  Arguments have also been presented that suggest the PM cycle introduces spurious force errors that can only been corrected by using a long range PP correction (out to 5 PM cells). It is certainly true that PM codes implemented with the so called `Poor Man's Poisson solver' \\cite{BR69}, and Cloud-in-cell interpolation do suffer from large ($\\sim$50\\%) directional errors in the force around 2-3 grid spacings. However, as has been shown, first by Eastwood (see \\cite{He81} for references) and more recently by Couchman, a combination of higher order assignment functions, Q-minimized Green's functions, and directionally optimized differencing functions can reduce errors in the inter-particle forces to sub 0.3\\% levels (RMS). Surprisingly, although CIC gives a smooth force law (as compared to NGP), it does not reduce the angular isotropy of the mesh force. Indeed, in two dimensions, moving from CIC to TSC interpolation  reduces directional errors from 50\\% to 6\\% and Q-minimization of the Green's function reduces the anisotropy to sub 0.5\\% levels \\cite{E74}. Furthermore, the technique of interlacing can be used to improve the accuracy of the PM force still further, but the additional FFTs required for this method rapidly lead to diminished returns.  To date we have used this code to simulate problems ranging from galaxy formation to large-scale clustering. As emphasized in the introduction, the simple programming model provided by OpenMP has enabled us to rapidly prototype new physics algorithms which in turn has lead to the code being applied across a diverse range of astrophysics. Developing new physics models with this code takes a matter of hours, rather than the days typical of MPI coding.  We plan to make a new version of the code, incorporating more streamlined data structures and minor algorithmic improvements, publically available in the near future."
    },
    "0512/astro-ph0512206_arXiv.txt": {
        "abstract": "We present new analytic calculations of the coupling between ultraviolet resonance photons and the population of the hyperfine states in the light elements (H, D, ${\\rm ^3He^+}$) which include several previously neglected physical processes. Among these are the backreaction of resonant scattering on the pumping radiation, the scattering of Ly$\\beta$ photons and the effect of local departure from pure Hubble flow. The application of the new treatment to the redshifted hydrogen 21 and deuterium 92 cm lines from the high-redshift universe results in an amplitude correction of up to an order of magnitude. We further show that the standard assumption that ultraviolet pumping drives the spin temperature towards the kinetic temperature does not hold for deuterium, whose spin temperature is generally negative. ",
        "introduction": "Introduction} The hyperfine transitions in hydrogen, deuterium and helium-3 provide an opportunity to study the universe during its ``dark ages'' and the epoch of reionization that followed. From the radio continuum produced by redshifted lines of these elements, which can be observed in absorption or in emission against the cosmic microwave background (CMB),  information about the thermal evolution, density perturbation spectrum and ionization history of the universe can be obtained. However, the correct interpretation of the measured signal requires knowledge of the processes which decouple the spin temperature of the atoms, $T_{\\rm s}$, from the CMB temperature, $T_{\\rm CMB}$. Of these, one of the most important during the epoch of reionization ($6\\lsim z\\lsim 20$) is ultraviolet pumping. Resonant photons, such as Ly$\\alpha$, absorbed by the atom in the ground state may, after the decay of the excited state, leave the electron in a different hyperfine level. This process couples the spin temperature to the color temperature of the resonant photons, $T_\\alpha$ \\cite{f8} \\begin{eqnarray} \\label{Tsp} T_{\\rm s}=\\frac{T_{\\rm CMB}+y_\\alpha T_\\alpha+y_{\\rm c}T_{\\rm k}}{1+y_\\alpha+y_{\\rm c}}, \\end{eqnarray} where $T_{\\rm k}$ is the kinetic temperature, $y_{\\rm c}$ is a constant proportional to the collisional excitation rate and $y_\\alpha$ is given by \\begin{eqnarray} y_{\\alpha}=\\frac{P_{10}T_*}{A_{10}T_{\\alpha}} \\end{eqnarray} where $T_*=h\\nu_{\\rm hyp}/k$, $\\nu_{\\rm hyp}$ is the frequency of the hyperfine transition, $A_{10}$ is the spontaneous emission coefficient of the hyperfine transition and $P_{10}$ is the radiative deexcitation rate for the upper hyperfine level due to Ly$\\alpha$ photons.  As we show in this paper, the correct application of the above expressions requires us to take into account several physical processes which have previously been neglected. The first of these is the effect of the hyperfine splitting on the photon color temperature and intensity. In their pioneering papers, Wouthuysen (1952) and Field (1958, 1959) argued that in the limit of large optical depth,  $T_\\alpha$  must approach the kinetic temperature of the atoms, to which the photons are coupled by recoil. As a result, since the  Ly$\\alpha$ resonance scattering depth for hydrogen atoms in the neutral intergalactic medium (IGM) is extremely large, it has been customary ever since then to replace $T_\\alpha$ with $T_{\\rm k}$ in eqs. 1 and 2 above. That argument, however, did not take into account the effect of the hyperfine splitting on the change of photon energy during scattering.  Just as the spin temperature of the atoms is affected by the spectrum of the UV photons in the resonance line, so the UV spectrum around the resonance is in turn affected by the spin temperature. This backreaction shifts $T_\\alpha$ towards $T_{\\rm CMB}$, thus reducing the coupling between the kinetic and spin temperatures. In addition, we consider the effect of scatterings on the radiation intensity near the resonance. Chen \\& Miralda-Escude (2004) showed that, in a homogeneously expanding cosmological medium, the intensity of the continuum radiation drops around the resonance, thus reducing the value of $y_\\alpha$. Here we derive an analytical formula for the amplitude of this drop. We find significant discrepancies between our analytical calculations and some of their numerical results. Furthermore, we extend their analysis to consider departures from pure Hubble flow, associated with cosmic structure formation, such as a contracting medium which is appropriate for overdense regions. Finally, we also consider scattering by resonant photons other than Ly$\\alpha$. Ly$\\alpha$ photons have the highest scattering cross-section and, unlike higher resonances like Ly$\\beta$, they are not destroyed by cascade after just a few scatterings. For these reasons, the higher resonances are generally assumed to be unimportant. While correct for hydrogen, this assumption proves to be a very gross error in the case of deuterium, for which the role of Ly$\\beta$ photons is crucial.  The corrections derived here to the hydrogen spin temperature which results from Ly$\\alpha$ pumping in the early universe can have a significant effect on predictions of the 21-cm brightness temperature and its fluctuations.  As we shall show their importance depends on the intensity of the UV background and the temperature of the IGM. An approximation which is often made in recent theorectical predictions of the 21-cm background from the epoch reionization involves the assumption that the UV pumping background is high enough to make $T_s$ equal $T_k$, while some other process (e.g. X-ray background) heats the IGM without ionizing it, to $T_k\\gg T_{CMB}$ \\cite{MMR,T,C,CM,ZFH,FSH,BL,Aal,Sal,Mal}. In that limit the 21-cm brightness temperature becomes independent of the actual value of $T_s$, so the calculations are greatly simplified. However, the period either before the UV background reaches such high intensity or the neutral IGM is heated is not well represented by this limiting approximation. This includes the epoch when the 21-cm signal is in absorption and potentially at its strongest. That is the case in which the results presented here will be important.  While this paper was in preparation, a preprint of Hirata (2006) was posted in which some of the above effects were calculated numerically. We find an excellent agreement between our analytic solutions and his numerical results.  In \\S 2, we estimate the effect of the hyperfine transition on the color temperature of Ly$\\alpha$ photons. In \\S 3, we calculate the change of intensity of Ly$\\alpha$ photons due to their scattering by atoms. In \\S 4, we estimate the effect of these processes on the radio emission signals from D and ${\\rm ^3He^+}$. In \\S 5, we discuss the implications of our results for the interpretation of hydrogen and deuterium radio signals from the epoch of reionization and the end of the cosmic dark ages. ",
        "conclusions": "We have shown that the accurate modeling of the photon spectrum around the hydrogen Ly$\\alpha$ resonance results in a significant correction to the radiative pumping efficiency. For the cosmological parameters given by the latest WMAP results, $\\Omega_{b0} h_0=0.03$ and $\\Omega_{m0}=0.25$, \\cite{Sp} and unperturbed Hubble expansion, $T_{\\rm s}$ is correctly given by eq. \\ref{Tf} if $y_{\\alpha,eff}$ is given by \\begin{eqnarray} \\label{cor} \\frac{y_\\alpha,eff}{y_{\\alpha,0}}=e^{-0.37(1+z)^{1/2}T_{\\rm k}^{-2/3}}\\left(1+\\frac{0.4}{T_{\\rm k}}\\right)^{-1}, \\end{eqnarray} where $y_{\\alpha,0}$ is the previously used coupling constant, given by \\begin{eqnarray} y_{\\alpha,0}=\\frac{16\\pi^2 T_*e^2 f_{12}J_0}{27 A_{10}T_{\\rm k} m_{\\rm e}c}, \\end{eqnarray} where $f_{12}=0.416$ is the oscillator strength of the Ly$\\alpha$ transition and $J_0$ is the intensity at Ly$\\alpha$ resonance, when the backreaction caused Ly$\\alpha$ scattering by is neglected .  Prior to reionization the mean differential brightness temperature of the hydrogen 21 cm signal is given by \\begin{eqnarray} \\delta T_{\\rm b}=0.03\\; {\\rm K} \\left(\\frac{T_{\\rm s}-T_{\\rm CMB}}{T_{\\rm s}}\\right) \\times \\hspace{2cm} \\nonumber \\\\ \\left(\\frac{\\Omega_{b} h_{100}}{0.03}\\right) \\left(\\frac{\\Omega_{m}}{0.25}\\right)^{-1/2}  \\left(\\frac{1+z}{10}\\right)^{1/2}. \\end{eqnarray} From the above equation it follows that the strongest signal is obtained when the spin temperature is well below $T_{\\rm CMB}$, in which case $T_{\\rm b}\\propto T_{\\rm s}^{-1}$. When radiation sources begin to form, $T_{\\rm s}$ indeed starts to drop as it decouples from $T_{\\rm CMB}$ and approaches $T_{\\rm k}$ instead. However, if X-ray photons are present, with an intensity which increases simultaneously with those in the UV, these would at the same time heat the gas, so that, eventually, $T_{\\rm k}$ and $T_{\\rm s}$ become greater than $T_{\\rm CMB}$.  Chen \\& Miralda-Escude (2004) estimated that in this case the strongest signal ($\\delta T_{\\rm b}\\sim -20$ mK) arises when the gas is still in absorption and $T_{\\rm s}$ is not yet fully decoupled from $T_{\\rm CMB}$, i.e., when the signal is still very sensitive to the changes in the coupling strength,  $y_{\\alpha}$.  \\begin{figure} \\resizebox{\\columnwidth}{!} {\\includegraphics{f5.eps}} \\caption{\\label{T1}Hydrogen spin temperature at $z=12$ for different radiation intensities, $y_{\\alpha,0} T_{\\rm k}=1$, 10, $10^2$, $10^3$ K (lower curves correspond to higher values of  $y_{\\alpha,0}$). The dashed and the dashed-dotted lines are calculated, respectively, with and without the correction to the coupling constant. The solid lines are $T_{\\rm CMB}$ and $T_{\\rm k}$.} \\end{figure}   \\begin{figure} \\resizebox{\\columnwidth}{!} {\\includegraphics{f6.eps}} \\caption{\\label{T2}The correction factor to the gas absorption signal for different radiation intensities $y_{\\alpha,0}T_{\\rm k}=1$ (solid), 10 (dashed), $10^2$ (dotted) and  $10^3$ (dashed-dotted line).} \\end{figure}   A further complication arises from the small-scale density fluctuations. Even though the resolution of planned future experiments is currently limited to objects of at least $\\sim 1$ Mpc size, which are still mostly in the linear regime at the relevant epoch, the spectrum around the resonance is determined by conditions on much smaller ($<1$ kpc) scale, where nonlinear departures from Hubble flow are important. To illustrate the signal dependence on the small scale-structure we have plotted the spin temperature  of the adiabatically evolving hydrogen gas at $z=12$ at various overdensities (Figs. \\ref{Ga}-\\ref{T2}). In general, there is no one-to-one relationship between the density and the expansion rate, so to make this figure we assumed that all over/underdensities evolve as spherical ``top-hat'' density perturbations (which explains the break at $\\delta=4.6$ where the gas switches from expansion to contraction). This yields a unique dependence of $\\tau_{\\rm GP}$ on local overdensity for a given redshift (Fig. \\ref{Ga}). The results are shown in Figure \\ref{T1} for different values of $y_{\\alpha,0}T_{\\rm k}$ (i.e. different intensities), which can be related to the ratio of photons in the UV background in the frequency range between Ly$\\alpha$ and  Ly-limit per atom, $N_{\\alpha}$, according to $ y_{\\alpha,0}T_{\\rm k}\\sim 0.2N_{\\alpha}(1+z)^3$. The spin temperature changes significantly when the corrections to the coupling constant (eq. \\ref{cor}) are taken into account, except when the radiation intensity is very high. Figure \\ref{T2} shows how the local contributions to the overall 21 cm signal change after including this correction. The net effect of departures from uniform Hubble expansion on the mean 21 cm signal depends on the weighted contributions from gas at different overdensities, and, since this dependence is non-linear, the contributions of over- and under-dense regions do not cancel each other out.  For hyperfine transition of deuterium, we found that a drastic difference is made by inclusion of Ly$\\beta$ photons in the analysis, which practically do not affect hydrogen. Thus the ratio between the deuterium and hydrogen radio signals is not constant, but depends sensitively on the UV spectrum. Unfortunately, this eliminates the possibility of the precision radio measurements of D/H ratio suggested by Sigurdson \\& Furlanetto \\cite{SF}, except for the period before the first UV sources turn on ($z\\gsim 40$)."
    },
    "0512/astro-ph0512083_arXiv.txt": {
        "abstract": "{ We investigate possible formation sites of the cannonballs (in the gamma ray bursts context) by calculating their physical parameters, such as density, magnetic field and temperature close to the origin. Our results favor scenarios where the cannonballs form as instabilities (knots) within magnetized jets from hyperaccreting disks. These instabilities would most likely set in beyond the light cylinder where flow velocity with  Lorentz factors as high as 2000 can be achieved. Our findings challenge the cannonball model of gamma ray bursts if these indeed form inside core-collapse supernovae (SNe) as suggested in the literature; unless hyperaccreting disks and the corresponding jets are common occurrences in core-collapse SNe. ",
        "introduction": "It has been argued in the literature that, as an alternative to the fireball scenario (e.g. Piran \\cite{piran99}, and references therein), the so-called cannonball (CB) model provides a good fit to the observed GRB flux and temporal variations (\\cite{dar04}). For example, to explain GRBs, CBs must be created in supernova explosions and accelerated to  high Lorentz factors, $\\Gamma_{\\rm CB} \\sim 1000$. However, the origin of these highly relativistic ``balls'' of matter has not yet been investigated and is the subject of much debate and controversy. In order to  shed some light on the still open questions of their formation and early evolution we investigate, in this paper, the CB physical conditions at the origin given their features at the distance when they become transparent to their enclosed radiation as required to explain GRBs. Our proposal is that the conditions within the CB as we scale the distance down along its path to the origin should be an indication of their formation site. This, despite the simplicity of our approach, we hope might help elucidate some questions related to the origin/existence of these CBs. We start in Sect. 2 by a brief introduction to the CB model as described in  Dar \\& De R{\\'u}jula (2004). In particular CBs conditions at infinity which best fits GRB lightcurves are isolated. In Sect. 3 we present the methods we adopted to extrapolate back to the CB source.  In Sect. 4, given the conditions at the origin, we study possible formation sites and explore  formation mechanisms. Sect. 5 is devoted to the study of mechanisms capable of acceleration CBs to Lorentz factors as high as $\\sim 1000$. We summarize our results and conclude in Sect. 6. ",
        "conclusions": "Assuming that CBs move and expand with a constant velocity we have estimated the CB conditions as close as possible to their origin. CBs require extremely high internal magnetic fields when they are formed with field strength exceeding $\\sim 10^{15}$ G. The temperature was found to be of the order of $10^{11} - 10^{12}$ K. The physical parameters of the CBs at the origin are, within an order of magnitude estimates, indicative of hyperaccreting disks. However, if formed in the accretion disk we find it challenging to accelerate the CBs to the high Lorentz factors. The coronal origin is ruled out because of the unrealistically high coronal magnetic flux necessary to form the CBs. Our results instead hint to a jet origin of CBs. The radius ($< D_{\\rm lc}$) and mass flow ( $10^{-5} M_\\odot/s$) in a jet from a hyperaccreting can account for the CB mass and density. Furthermore, this outflow can be accelerated to $\\Gamma\\sim 2000$ by MHD processes (Fendt \\& Ouyed 2004). Any instability in this outflow beyond the light cylinder could lead to CB formation.  We thus suggest that CBs form as instabilities in ultra-relativistic jets emanating from the surface of hyperaccretion disks. The tight link between SNe and the CB model for GRB requires that all (or almost all) core collapse SNe will produce CBs. Our work, within its limitations, implies that hyperaccretion disks must be a common occurrence in core collapse SNe to accommodate the CB model - a notion which remains to be confirmed."
    },
    "0512/astro-ph0512426_arXiv.txt": {
        "abstract": "Future weak lensing measurements of cosmic shear will reach such high accuracy that second order effects in weak lensing modeling, like the influence of baryons on structure formation, become important.  We use a controlled set of high-resolution cosmological simulations to quantify this effect by comparing pure N-body dark matter runs with corresponding hydrodynamical simulations, carried out both in non-radiative, and in dissipative form with cooling and star formation.  In both hydrodynamical simulations, the clustering of the gas is suppressed while that of dark matter is boosted at scales $k>1\\hmpci$. Despite this counterbalance between dark matter and gas, the clustering of the total matter is suppressed by up to 1 percent at $1\\la k \\la 10\\hmpci$, while for $k \\approx 20\\hmpci$ it is boosted, up to 2 percent in the non-radiative run and 10 percent in the run with star formation. The stellar mass formed in the latter is highly biased relative to the dark matter in the pure N-body simulation.  Using our power spectrum measurements to predict the effect of baryons on the weak lensing signal at $100<l<10000$, we find that baryons may change the lensing power spectrum by less than 0.5 percent at $l<1000$, but by 1 to 10 percent at $1000<l<10000$. The size of the effect exceeds the predicted accuracy of future lensing power spectrum measurements and will likely be detected.  Precise determinations of cosmological parameters with weak lensing, and studies of small-scale fluctuations and clustering, therefore rely on properly including baryonic physics. ",
        "introduction": "Weak gravitational lensing directly measures the projected mass distribution and is emerging as one of the most powerful and robust probes of the large scale structure of the universe, and the nature of dark matter, dark energy and gravity \\citep{Huterer02,Hu02,Jain03,Takada04,Ishak05,Knox05}. Ongoing and upcoming surveys such as CFHTLS\\footnote{http://www.cfht.hawaii.edu/Science/CFHLS}, DES\\footnote{http://www.darkenergysurvey.org/}, LSST\\footnote{http://www.lsst.org/}, Pan-STARRS\\footnote{http://pan-starrs.ifa.hawaii.edu/}, SKA\\footnote{http://www.skatelescope.org/} and SNAP\\footnote{http://snap.lbl.gov/} will significantly reduce statistical errors in lensing power spectrum measurements to the $<1\\%$ level at $l\\sim 1000$.  As new analysis techniques enable a significant reduction of systematic errors in cosmic shear \\citep[][and references therein]{Jarvis04,Jain05,Heymans05} and cosmic magnification measurements \\citep{Menard02,Scranton05,Zhang05a,Zhang05b}, weak lensing measurement is quickly entering the precision era. To match this accuracy, many simplifications in theoretical modeling have to be scrutinized in detail \\citep{Bernardeau98,Schneider98,Dodelson05a,Dodelson05b,White05}.   A widely adopted simplification in weak lensing modeling is the assumption that baryons trace dark matter perfectly. With this simplification, weak lensing involves only gravity and collision-less dark matter dynamics, allowing it to be accurately predicted with the aid of N-body simulations.  However, on small scales, baryons {\\em do not follow} the dark matter distribution.  Recently, \\citet{white04} and \\citet{zhanknox04} estimated the effects of cooling gas and intra-cluster gas on the lensing power spectrum, respectively.  These two components of baryons can both have an effect of a few percent on the lensing power spectrum at $l\\sim 3000$, but with opposite signs. Such effects exceed future measurement errors and are certainly relevant in order to exploit the full power of precision measurements of weak lensing.  However, analytical models are simplified with {\\it ad hoc} parameters and lack the ability to deal with back-reactions of baryons on the dark matter. They can not robustly answer several key questions such as (1) to what level the two effects cancel, (2) at what scales each of them dominates, and (3) how large the effect is for the non-virialized intergalactic medium where most of the baryons reside.  To quantify the effect of baryons on lensing statistics accurately, it is hence necessary to use hydrodynamical simulations with all relevant gas physics included.  In this {\\it letter}, we analyze a controlled set of simulations to address this issue. ",
        "conclusions": "In this {\\it Letter}, we used a controlled set of high resolution N-body and hydrodynamical/N-body simulations to study the influence of baryons on the clustering of cosmic matter. Since the three simulations we used have identical initial condition, simulated power spectra suffer essentially the same sample variance. Since we quantify the baryonic effect as ratios of corresponding power spectra, the result presented in this {\\it Letter} is effectively free of sample variance. In both the non-radiative simulation and the simulation with gas cooling and star formation, the clustering of the gas is suppressed while that of dark matter is boosted at $k>1\\hmpci$. The stellar mass is highly biased relative to the dark matter in the pure N-body simulation.  Despite a partial counterbalance between the dark matter and the gas, the clustering of the total matter is suppressed by up to 1\\% at $1\\la k \\la 10\\hmpci$, and is boosted up to 2\\% in the non-radiative run, and 10\\% in the star formation run at $k \\approx 20\\hmpci$. Using these power spectrum measurements to study the baryonic effect on the weak lensing shear measurement at $100<l<10000$, we found that baryons can change the shear power spectrum by less than 0.5\\% at $l<1000$, but by 1\\% to 10\\% at $1000<l<10000$. Therefore, the influence of baryons on the clustering of cosmic matter will be detected in future weak lensing surveys. Understanding these baryonic effects is not only important for galaxy formation, but also crucial for accurately determining cosmological parameters with cosmic shear, and for constraining the initial fluctuations on small scales."
    },
    "0512/astro-ph0512610_arXiv.txt": {
        "abstract": "We found that in regions of high mass star formation the CS emission correlates well with the dust continuum emission and is therefore a good tracer of the total mass while the N$_2$H$^+$ distribution is frequently very different. This is opposite to their typical behavior in low-mass cores. The behavior of other high density tracers varies from source to source but most of them are closer to CS. Radial density profiles in massive cores are fitted by power laws with indices about $-1.6$, as derived from the dust continuum emission. The radial temperature dependence on intermediate scales is close to the theoretically expected one for a centrally heated optically thin cloud. The velocity dispersion either remains constant or decreases from the core center to the edge. Several cores including those without known embedded IR sources show signs of infall motions. They can represent the earliest phases of massive protostars. There are implicit arguments in favor of small-scale clumpiness in the cores. ",
        "introduction": "Many important aspects of high mass star formation (HMSF) remain unknown or poorly known. This is caused by the much more complicated appearance of HMSF regions in comparison with their low-mass counterparts and by the relative difficulties of their studies due to larger distances, etc. Therefore, detailed investigations of the structure and kinematics of HMSF cores are of great importance. For example, different models of star formation predict different density profiles in the cores and different frequency of outflow occurrence.  An important question in this respect is a selection of reliable tracers of physical parameters in HMSF regions. Chemical effects can seriously distort their appearance. It is now well established that the central parts of dense low mass cloud cores suffer strong depletion of many molecules (in particular, CO, HCO$^+$ and CS) onto dust grains (\\cite[Caselli, Walmsley, Tafalla, \\etal\\ 1999]{Caselli99}; \\cite[Caselli, Walmsley, Zucconi, \\etal\\ 2002]{Caselli02}; \\cite[Kramer, Alves, Lada, \\etal\\ 1999]{Kramer99}; \\cite[Willacy, Langer, \\& Velusamy 1998]{Willacy98}; \\cite[Jessop \\& Ward-Thompson 2001]{Jessop01}; \\cite[Tafalla, Myers, Caselli, \\etal\\ (2002)]{Tafalla02}; \\cite[Bergin, Ciardi, Lada, \\etal\\ 2001]{Bergin01}), suggesting that CS (so far considered a high density tracer) does not actually probe the central core regions. On the other hand, N$_2$H$^+$ is usually an excellent tracer of dust continuum emission (\\cite[Caselli \\etal\\ 2002]{Caselli02}), implying that this species depletes at higher densities compared to CO, despite of the same binding energies (\\cite[\\\"Oberg, van Broekhuizen, Fraser, \\etal\\ 2005]{Oberg05}).  The situation is less clear in HMSF cores where the mechanisms of chemical differentiation are probably different (e.g. not linked to molecular freeze-out). There were several studies which show chemical variations in a few similar regions (\\cite[Ungerechts, Bergin, Goldsmith, \\etal\\ (1997)]{Ungerechts97}; \\cite[Bergin, Goldsmith, Snell, \\etal\\ 1997]{Bergin97}). Here, we present observations of a statistically significant sample of HMSF cores. ",
        "conclusions": "\\label{sec:concl} We have shown that effects of chemical differentiation in warm massive cores are very different from those in cold low mass cores. In particular, CS is a good tracer of total mass here, while N$_2$H$^+$ is not. A relatively high N$_2$H$^+$/CS ratio can indicate the earliest phases of massive protostars. The internal structure of HMSF clumps is more or less consistent with the standard theory of star formation."
    },
    "0512/hep-ph0512107_arXiv.txt": {
        "abstract": "The electron-neutrino bremsstrahlung process has been considered in the framework of electro-weak theory. The scattering cross section has been calculated in the center of mass frame and approximated to extreme relativistic as well as non-relativistic case. The rate of energy-loss via this type of bremsstrahlung process has been obtained both in non-degenerate and degenerate region. The effect of this electron-neutrino bremsstrahlung process in different ranges of temperature and density characterizing the late stages of stellar evolution has been discussed. It is found from our study that this bremsstrahlung process is highly important in the non-degenerate region, although it might have some significant effect in the extreme relativistic degenerate region. \\vspace{0.5cm}\\\\\\noindent{PACS :\\hspace{0.2cm}}{\\it 12.15.-y; 13.85.Hd; 13.85.Lg; 14.60.Lm; 14.70.-e; 97.60.-s} ",
        "introduction": "It is known that neutrino emission processes have significant contribution in the late stages of the stellar evolution; when the core of the stars collapses, the neutrino emission occurs enormously in the temperature range $> 10^{9}$ K. In that range the neutrino emission process is important due to extremely large mean free path of the neutrino. It is believed that the stellar matter (even under some extreme conditions as in white dwarves or in neutron stars) is almost transparent to the neutrinos, such in contrast with its behavior with respect to photons. Once neutrinos are produced inside the star the rate of energy loss is higher resulting faster evolution. There are several processes functioning the major role in the energy loss from star in the late stage through the emission of neutrino-antineutrino pair. First time it was pointed out by Gamow and Schoenberg \\cite{Gamow} that neutrino might be emitted from star via $\\beta$ decay, which is termed as Urca process. Poentecorvo \\cite{Pontecorvo} showed the theoretical possibility of the formation of neutrino pairs in collisions between electrons and nuclei and the process was investigated with possible application in astrophysics by Gandel'man and Pinaev \\cite{Gandel'man}. The reaction rate of this process was calculated in detail by Chiu and his collaborators \\cite{Chiu1960, Chiu1961, Chiu}. Dicus \\cite{Dicus1972} drew a brief outline about some neutrino emission processes from star, though there might exist a few more processes having remarkable effect during the later stages of the stellar evolution. Earlier most of the were studied in the framework of the V-A interaction theory \\cite{Feynman}. Later, photon-neutrino weak coupling theory was introduced and some of those processes such as neutrino synchrotron process \\cite{Raychaudhuri1970}, photo-coulomb neutrino process \\cite{Raychaudhuri} were studied in this framework. The advancement of the Standard Model added a new dimension to study several neutrino emission processes such as pair annihilation \\cite{Dicus1972}, photo production \\cite{Dicus1972}, photon-photon scattering \\cite{Dicus1972,Dicus1993,Abbasabadi1998,Abbasabadi2000}, photo-coulomb neutrino process \\cite{Rosenberg, Bhattacharyya} etc. Recently Itoh et al. \\cite{Itoh} reviewed a number of neutrino emission processes and discussed their significance from astrophysical point of view. A minor extension of the Standard Model was done due to the existence of neutrino mass resulted from the `Solar Neutrino Problem' and `Atmospheric Neutrino Anomaly' \\cite{Santo}. It should be noted that a new theory of weak interaction is yet to be developed by introducing the neutrino mass. In the calculations of some weak processes the effect of this neutrino mass, whatever small it may be, may play an important role, for example photon-neutrino interaction \\cite{Dicus2000,Dodelson}.\\\\\\indent In this paper we have studied the `electron-neutrino bremsstrahlung process' given by $$e^{-}+e^{-}\\longrightarrow e^{-}+e^{-}+\\nu+\\bar{\\nu}$$ according to the electro-weak interaction theory. Previously it was considered by Cazzola and Saggion \\cite{Cazzola} while calculating the energy-loss rate by using Monte Carlo Integration method without evaluating the scattering cross-section explicitly. But the scattering cross-section, depending on the energy of the incoming electron, can give a very clear idea about the nature of this electro-weak process. It is to be approximated for extreme relativistic as well as non relativistic limit. Cazzola and Saggion \\cite{Cazzola} considered only the non-degenerate case though many stars in the later phases such as white dwarves, neutron stars etc. are degenerate. We cannot ignore the possibility of occurring the electron-neutrino bremsstrahlung process in degenerate star. We have considered all such cases separately and discussed all possible outcomes of this bremsstrahlung process. We also like to visualize a picture under what circumstances the process will have some significant effect. It cannot be denied that due to some approximations a little bit deviation may occur from the original result, but that will not deter to predict the physical picture. The role of this process has been studied thoroughly at different temperature and density ranges that characterize the late stages of the stellar evolution. It has also been pointed out in which range this process has significant effect. ",
        "conclusions": "In the calculations of scattering cross-section we have used an approximation in the equation (2.16), but the Table-1 shows our result is very close to that generated by the software. It strongly supports the approximation we have used in the equation (2.16). The scattering cross section obtained under the frame-work of electro-weak theory is very small, especially in the non-relativistic case, but as the mean free path of the neutrinos is much longer than the scale of stellar radius the electron-neutrino bremsstrahlung process may have some effect to release energy from star at high temperature and density. The relativistic effect comes into play when the temperature exceeds $6\\times 10^{9}$ K. It is evident from our work that the electron-neutrino bremsstrahlung process yields a large amount of energy loss from the stellar core when core temperature $\\geq 10^{10}$ K, both in non-degenerate as well as degenerate region. In that temperature range the radiation pressure is so dominating that the gas pressure has negligible effect \\cite{Chandrasekhar}. In this extreme relativistic region the process contributes significantly when the electron gas is non-degenerate. That was also shown by Cazzola and Saggion \\cite{Cazzola}. But they did not calculate the energy loss rate in the degenerate region; though they indicated that the electron neutrino bremsstrahlung process might be highly significant in that region. We have calculated the energy-loss rate to obtain an analytical expression for the energy loss rate when the electrons are strongly degenerate. A typical example of the extreme-relativistic degenerate stellar object is the newly born neutron star, which is the result of type-II Supernova. Our study reveals that the energy loss rate in the non-degenerate region is higher than that calculated in the degenerate region. This clearly indicates that though during the neutron star cooling electron-neutrino bremsstrahlung plays a significant role, but the process becomes more important to carry away the energy from the core of pre-Supernova star, which is a relativistic non-degenerate stellar object.\\\\\\indent Non-relativistically the process becomes insignificant unless the temperature is sufficiently high, at least the temperature should attain $10^{8}$ K. At this temperature the burning of helium gas in the stellar core takes place \\cite{Hayashi}. In the temperature range $10^{8}- 10^{9}$ K the gas pressure is dominating over the radiation pressure; though the effect of radiation pressure cannot be neglected in this region. In addition, the region will be non-degenerate if density $<2\\times 10^{6}$ gm/cc. The electron neutrino bremsstrahlung process may have some effect in this region though the energy loss rate is not so high as it is in the extreme relativistic case. In the low density the energy loss rate increases rapidly with rising core temperature. Eventually the process contributes in both degenerate and non-degenerate cases, whereas degenerate electrons participate only when the density of the medium is very high.  Hence, the electron-neutrino bremsstrahlung process is an important energy-generation mechanism during the evolution of stars, particularly in the later stages."
    },
    "0512/astro-ph0512604_arXiv.txt": {
        "abstract": "Astrophysical neutrinos at $\\sim$EeV energies promise to be an interesting source for astrophysics and particle physics. Detecting the predicted cosmogenic (``GZK'') neutrinos at 10$^{16}$ - 10$^{20}$ eV would test models of cosmic ray production at these energies and probe particle physics at $\\sim$100~TeV center-of-mass energy. While IceCube could detect $\\sim$1 GZK event per year, it is necessary to detect 10 or more events per year in order to study temporal, angular, and spectral distributions. The IceCube observatory may be able to achieve such event rates with an extension including optical, radio, and acoustic receivers.  We present results from simulating such a hybrid detector. ",
        "introduction": "Detecting and characterizing astrophysical neutrinos in the 10$^{16}$~eV to 10$^{20}$~eV range is a central goal of astro-particle physics.  The more optimistic flux models in this range involve discovery physics including topological defects and relic neutrinos.  Detecting the smaller flux of cosmogenic (or Greisen, Zatsepin, and Kusmin, ``GZK'') neutrinos produced via ultra-high energy cosmic ray interaction with the cosmic microwave background would test models of cosmic ray production and propagation and of particle physics at extreme energies. With $\\sim$100 detected events, their angular distribution would give a measurement of the total neutrino-nucleon cross section at $\\sim$100 TeV center of mass, probing an energy scale well beyond the reach of the LHC.  Hence, as a baseline, a detector capable of detecting $\\sim$10 GZK events per year has promising basic physics potential.  If any of the more exotic theories predicting greater EeV neutrino fluxes is correct, the argument in favor of such a detector is even stronger.  To detect $\\sim$10 GZK events per year, a detector with an effective volume of $\\sim$100~km$^3$ at EeV energies is necessary. In addition to the possibility of identifying neutrino-induced air showers, there are three methods of ultra-high energy neutrino detection in solid media: optical, radio, and acoustic.  Optical Cherenkov detection is a well-established technique that has detected atmospheric neutrinos up to 10$^{14}$ eV and set limits up to 10$^{18}$~eV \\cite{Chirkin}.  Radio efforts have produced steadily improving upper limits on neutrino fluxes from 10$^{16}$ to 10$^{25}$~eV \\cite{radio}. Acoustic detection efforts are at an earlier stage, with one limit published thus far from 10$^{22}$ to 10$^{25}$~eV \\cite{Vandenbroucke05}.  The currently planned 1~km$^3$ optical neutrino telescopes expect a GZK event rate of $\\sim$1 per year. It is possible to extend this by adding more optical strings for a modest additional cost \\cite{Halzen}, but it's difficult to imagine achieving 10 or more events per year with optical strings alone.  The radio and acoustic methods have potentially large effective volumes with relatively few receivers, but the methods are unproven in that they have never detected a neutrino.  Indeed, if radio experiments claim detection of a GZK signal, it may be difficult to confirm that it is really a neutrino signal. However, it may be possible to bootstrap the large effective volumes of radio and acoustic detection with the optical method, by building a hybrid detector that can detect a large rate of radio or acoustic events, a fraction of which are also detected by an optical detector. A signal seen in coincidence between any two of the three methods would be convincing. The information from multiple methods can be combined for hybrid reconstruction, yielding improved angular and energy resolution.  We simulated the sensitivity of a detector that could be constructed by expanding the IceCube observatory currently under construction at the South Pole.  The ice at the South Pole is likely well-suited for all three methods:  Its optical clarity has been established by the AMANDA experiment \\cite{Chirkin}, and its radio clarity and suitability for radio detection in the GZK energy range has been established by the RICE experiment \\cite{radio}. Acoustically, the signal in ice is ten times greater than that in water. Theoretical estimates indicate low attenuation and noise \\cite{Price05}, and efforts are planned to measure both \\cite{SPATS05} with sensitive transducers developed for glacial ice \\cite{Nahnhauer05}. Here we estimate the sensitivity of such a detector by exposing all three components to a common Monte Carlo event set and counting events detected by each method alone and by each combination of multiple methods. ",
        "conclusions": "Ten-thousand events were generated at each half-decade in neutrino energy in a cylinder of volume 942~km$^3$.  For each method and combination of methods, the number of detected events was used to calculate effective volume as a function of neutrino energy (Fig. \\ref{fig2}).  This was folded with the GZK flux model of \\cite{ESS,Seckel} and the cross-section parametrizations of \\cite{Gandhi} to estimate detectable event rates (Fig. \\ref{fig2}).  We use a flux model which assumes source evolution according to $\\Omega_{\\Lambda}=$~0.7.  This model is a factor of $\\sim$2 greater than that for $\\Omega_{\\Lambda}=$~0 evolution; it is unclear which model is correct \\cite{Seckel}. For radio and acoustic, and their combination, all flavors and both interactions were included.  For those combinations including the optical method, only the muon channel has been simulated thus far; including showers will increase event rates for these combinations.  It may be possible to build an extension like that considered here for a relatively small cost. Holes for radio antennas and acoustic transducers can be narrow and shallow, and both devices are simpler than photo-multiplier tubes. The necessarily large acoustic channel multiplicity is partially offset by the fact that the acoustic signals are slower by five orders of magnitude, making data acquisition and processing easier.  The IceCube observatory will observe the neutrino universe from 10's of GeV to 100's of PeV. Our simulations indicate that extending it with radio and acoustic strings could produce a neutrino detector competitive with other projects optimized for high-statistics measurements of GZK neutrinos but with the unique advantage of cross-calibration via coincident optical-radio, optical-acoustic, and radio-acoustic events."
    },
    "0512/astro-ph0512268_arXiv.txt": {
        "abstract": "We present infrared echelle spectroscopy of three Herbig-Haro (HH) driving sources (SVS 13, B5-IRS~1 and HH~34~IRS) using Subaru-IRCS. The large diameter of the telescope and wide spectral coverage of the spectrograph allowed us to detect several H$_2$ and [Fe II] lines in the $H$- and $K$-bands. These include H$_2$ lines arising from $v$=1--3 and $J$=1--11, and [Fe II] lines with upper level energies of $E/k=1.1-2.7\\times10^4$~K. For all objects the outflow is found to have two velocity components: (1) a high-velocity ($-$70 to $-$130~km\\,s$^{-1}$) component (HVC), seen in [Fe II] or H$_2$ emission and associated with a collimated jet; and (2) a low-velocity ($-$10 to $-$30~km\\,s$^{-1}$) component (LVC), which is seen in H$_2$ emission only and is spatially more compact. Such a kinematic structure resembles optical forbidden emission line outflows associated with classical T Tauri stars, whereas the presence of H$_2$ emission reflects the low-excitation nature of the outflowing gas close to these protostars. The observed H$_2$ flux ratios indicate a temperature of $2-3\\times10^3$~K, and a gas density of $10^5$~cm$^{-3}$ or more, supporting shocks as the heating mechanism. B5-IRS~1 exhibits faint extended emission associated with the H$_2$-LVC, in which the radial velocity slowly increases with distance from the protostar (by $\\sim 20$~km\\,s$^{-1}$ at $\\sim$~500~AU). This is explained as warm molecular gas entrained by an unseen wide-angled wind. The [Fe II] flux ratios indicate electron densities to be $\\sim$10$^4$~cm$^{-3}$ or greater, similar to forbidden line outflows associated with classical T Tauri stars. Finally the kinematic structure of the [Fe II] emission associated with the base of the B5-IRS~1 and HH~34~IRS outflows is shown to support disk-wind models. ",
        "introduction": "Mass accretion and ejection are fundamental processes for the evolution of young stellar objects (YSOs). Understanding their mechanisms, and also physical conditions which determine stellar masses, is one of the major challenges of contemporary astrophysics. Mass accretion cannot proceed without removing angular momentum from the surrounding material. Theories predict that outflowing gas plays an important role in this process (see e.g. Blandford \\& Payne 1982; Shu et al. 2000; K\\\"onigl \\& Pudritz 2000). However, there are two crucial problems which hinder the investigation of this scenario through observations. One is the limited spatial resolution of telescopes, which has not been sufficient to fully resolve the jet/wind launching region, i.e. where the angular momentum transfer may occur. The other is the fact that such a region is often deeply embedded within a massive envelope, making it difficult to observe at optical to near-IR wavelengths.  Even so, classical T-Tauri stars, a class of low-mass YSOs, are directly observable at optical-UV wavelengths, allowing the study of mass accretion and ejection to within 100~AU of the central star for nearby sources. Studies of UV excess continuum and permitted emission lines strongly suggest that stellar bipolar magnetic fields are coupled to the inner edge of the disk at a few stellar radii, regulating stellar rotation to well below the break-up velocity (see, e.g., Calvet et al. 2000 and Najita et al. 2000 for reviews). Disk material seems to lose a significant amount of angular momentum at the inner edge of the disk, before falling down magnetic field lines and being shocked near the stellar surface. Such ``magnetospheric accretion flow'' is indeed observed as redshifted absorption in permitted lines (see, e.g., Najita et al. 2000). At the same time, a jet and/or wind seems to be launched within a few AU of the disk or at the inner disk edge. A physical link between these two phenomena is indicated by the fact that disk accretion rates, measured from the veiling continuum, are correlated with mass outflow signatures such as forbidden line luminosities (Hartigan, Edwards, \\& Ghandour 1995; Calvet 1997). Using the {\\it Hubble Space Telescope (HST)} Bacciotti et al. (2002), Coffey et al. (2004) and Woitas et al. (2005) show that the internal motion of forbidden line outflows is consistent with flow rotation. These results agree with the scenario that the flow removes angular momentum from the disk, thereby allowing mass accretion to occur. Studying the kinematics of outflowing gas at AU scales, which cannot be directly resolved by the {\\it HST} or adaptive optics with 8-10 m telescopes, could provide more crucial information for understanding their ejection. This spatial scale is being explored using the technique of spectro-astrometry which allows the spatial structure of emission/absorption features to be studied down to milliarcsecond scales (e.g., Bailey 1998; Takami et al. 2001; Takami, Bailey \\& Chrysostomou 2003; Whelan et al. 2004).  However, there is growing evidence that the stellar mass has almost been determined at younger evolutionary phases (Class 0--I phases), and mass accretion at the T Tauri phase is responsible for only a tiny fraction of the entire stellar mass (of order 1\\% --- see Calvet et al. 2000). It is often assumed, but not confirmed, that disk accretion within 100~AU of Class I protostars occurs by the same mechanism as for T Tauri stars. Br$\\gamma$ line profiles observed towards Class I protostars suggest that these stars are associated with magnetospheric accretion columns like T Tauri stars (see e.g., Davis et al. 2001). However, measured Br$\\gamma$ luminosities suggest that steady magnetospheric mass accretion can explain only a few percent of the total stellar mass, requiring episodic and violent mass accretion (see Calvet et al. 2000).  Class 0--I protostars also exhibit signatures of energetic mass ejection as optical jets and molecular bipolar outflows (see e.g., Reipurth \\& Bally 2001 and Richer et al. 2000 for reviews). While the extended parts of these outflows have been observed over decades, kinematics of outflowing gas within a few hundred AU has just begun to be revealed by near-IR high-resolution spectroscopy and spectro-imaging. Davis et al. (2001) observed H$_2$ 2.122~$\\mu$m emission at several HH driving sources, showing the presence of Molecular Hydrogen Emission Line (MHEL) regions, analogous to the forbidden emission line (FEL) regions associated with T Tauri stars. The H$_2$ emission shows a variety of profiles with multiple velocity peaks similar to FELs in T-Tauri stars. The kinematics of the H$_2$ 2.122~$\\mu$m emission differs from that of the [Fe II] 1.644~$\\mu$m emission, revealing complicated kinematic structure (Davis et al. 2003). At the same time Fabry-Perot observations (Davis et al. 2002) show the presence of small scale jets traced by H$_2$ emission, analogous to forbidden line jets associated with classical T Tauri stars. Evidence for H$_2$ emission from cavity walls is also seen for a few objects, suggesting the presence of a wide-angled wind.  Detailed studies of mass ejection in these objects could lead us to understand the mechanism of mass accretion, which is responsible for determining the stellar mass. We thus performed an extensive study of MHEL regions and [Fe II] emission line regions associated with three HH driving sources using Subaru-IRCS.  The rest of the paper is organised as follows. In \\S 2 we describe the observations and data reduction methods. In \\S 3 we present the observed kinematics seen in the H$_2$ 2.122~$\\mu$m and [Fe II] 1.644~$\\mu$m emission, the brightest H$_2$ and [Fe II] emission in our spectra. In \\S 4 we derive the extinction, density and temperature inferred from H$_2$ and [Fe II] line ratios. In \\S 5 we compare our results with forbidden line outflows associated with T Tauri stars, and investigate the mechanism of driving jets/winds for low-mass YSOs. In Paper II we show the results for other emission lines including CO and atomic lines. ",
        "conclusions": "We detect a number of infrared H$_2$ and [Fe II] lines towards three Herbig-Haro driving sources: SVS 13,  B5-IRS~1 and HH~34~IRS. These line include H$_2$ transitions with $v$=1--3 and $J$=1--11, and several [Fe II] lines with upper energy levels $E/k=1.1-2.7\\times10^4$~K. Outflowing gas in all objects show the presence of two blueshifted components at high ($-$70 to $-$130~km\\,s$^{-1}$) and low velocities ($-$10 to $-$30~km\\,s$^{-1}$), characteristically similar to forbidden line emission associated with T Tauri stars but without a low velocity peak for the forbidden lines, which is observed in the H$_2$ emission. In SVS 13, the HVC also appears in H$_2$ emission. This reflects the low-excitation nature of the outflowing gas close to these HH driving sources.  The observed H$_2$ flux ratios indicate a temperature and density of $2-3\\times10^3$~K and $\\ge 10^5$~cm$^{-3}$, respectively. In particular the temperature obtained from the $I_{2-1 S(1)}$/$I_{1-0 S(1)}$ ratio remains quite constant at $2\\times10^3$~K, lending support to shock heating for the excitation of the gas. In the jet of SVS 13, the gas density is seen to decrease further downstream. The observed level population of H$_2$ indicates a hydrogen number density $n_{H} >10^5 - 10^8$~cm$^{-3}$ depending on whether the excitation is dominated by H-H$_2$ or H$_2$-H$_2$ collisions, respectively. The observed electron density and temperature in [Fe II] emission line regions are $\\sim 10^4$~cm$^{-3}$ or greater, similar to FEL regions associated with T Tauri stars.  The absence of an LVC in [Fe II] emission in all three HH driving sources matches models based on magneto-centrifugal winds. Our results in B5-IRS~1 and HH~34~IRS so far support disk-wind models, which predict that the jet/wind originates from the disk surface at radii of a few AU. Currently the X-wind model, which predicts that the jet/wind originates from the inner disk edge within 0.1~AU of the star, cannot explain the absence of redshifted emission at the driving source. More detailed information regarding jet launching could be extracted if accurate collisional rate coefficients for [Fe II] $^2P$ states are obtained. In B5-IRS~1 the faint extended component of H$_2$-LVC slowly accelerates over a large spatial scale ($\\sim$~20~km\\,s$^{-1}$ at $\\sim$~500~AU). We suggest that the emission is due to interaction between a slow and unseen wide-angled wind and the surrounding ambient gas."
    },
    "0512/astro-ph0512574_arXiv.txt": {
        "abstract": "In this paper we present and discuss the effects of scattered light echoes (LE) on the luminosity and spectral appearance of Type Ia Supernovae (SNe). After introducing the basic concepts of LE spectral synthesis, by means of LE models and real observations we investigate the deviations from pure SN spectra, light and colour curves, the signatures that witness the presence of a LE and the possible inferences on the extinction law. The effects on the photometric parameters and spectral features are also discussed. In particular, for the case of circumstellar dust, LEs are found to introduce an apparent relation between the post-maximum decline rate and the absolute luminosity which is most likely going to affect the well known Pskowski-Phillips relation. ",
        "introduction": "\\label{sec:intro}  Due to the crucial role played by Type Ia SNe in modern cosmology (see for example the review by \\citealt{leibundgut}), it is essential to understand the systematics which might affect the recent conclusions about the expansion rate of the Universe. The availability of larger telescopes and the renewed interest for these objects as cosmological probes has opened the possibility of studying in great detail their evolution, an investigation which in the past was feasible for the closest events only. In these last few years, in fact, numerous Ia have been followed with an extensive time coverage as late as one year past maximum light, leading to very interesting results (see for example \\citealt{benetti05}).  In two cases, namely those of SN~1991T \\citep{schmidt} and SN~1998bu \\citep{cappellaro,jason}, the spectra at advanced phases were clearly dominated by a phenomenon which had nothing to do with the explosion mechanism but rather with the neighboring environment.  In both cases this very atypical behaviour could be explained with the contamination by SN light scattered into the line of sight by dust present in the surroundings of the explosion. The presence of such a phenomenon, also known as light echo (hereafter LE), was later confirmed by high spatial resolution HST imaging (\\citealt{sparks,garnavich}). Other cases have been discovered for core-collapse SNe like SN~1987A \\citep{crotts89}, SN~1993J \\citep{sugerman02,liu}, SN~1999ev \\citep{maund} and SN~2003gd \\citep{sugerman05,vandyk}. In the most recent years, LEs have been searched also in the case of historical SNe \\citep{martino, krause, rest}.  With the purpose of studying this effect and its possible applications to the analysis of the SN environments and their relation with the SN progenitors, we have started a series of papers on this topic. The first of them (\\citealt{paperI}; hereafter Paper I) was dedicated to the general characteristics of LEs, with particular attention to the effects of self absorption on LE luminosity, spectral and colour appearance. In this and subsequent papers we will instead present and discuss applications to known and test cases.  In principle, LEs can affect all types of SNe. Due to the young stellar population from which they are supposed to be originated, core-collapse SNe are of course favored with respect to thermonuclear explosions. However, both due to their high homogeneity (which makes the analysis much simpler) and their wide use in cosmology, here we will concentrate on Type Ia only. In particular, the present work deals with the effects of LEs on their luminosity and spectra, in order both to help the observers in disentangling intrinsic features from LE contamination and to understand whether LEs can affect the conclusions reached on their nature.  The paper is organized as follows. After giving a short introduction to the basic concepts of LE modeling in Section~\\ref{sec:basic}, in Section~\\ref{sec:speclib} we describe the spectral library used in our calculations. The results are then presented and discussed in Section~\\ref{sec:effects} for two different dust geometries, while Section~\\ref{sec:ext} is devoted to the implications on the extinction law. In Section~\\ref{sec:discuss} we discuss our findings which are briefly summarized in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have analyzed and discussed the effects of LEs on light curves, colours and spectra of Type Ia SNe for two different dust geometries: extended clouds (Case A) and small circumstellar clouds (Case B). The main results can be summarized as follows:  \\begin{enumerate} \\item The effect of Case A LEs is limited to the late phases. The exact time when the LE contribution dominates the global luminosity depends on dust distance and density. \\item There is a transition phase, during which the LEs contribution becomes more and more important and the global spectrum turns from a normal nebular spectrum into a pure LE spectrum. \\item During the transition phase, the overall colour tends to become bluer, spectral features broader and emission line peaks displaced with respect to a pure SN nebular spectrum. \\item During the LE-dominated phase the spectrum retains the peculiarities shown by the SN at maximum, so that the LE detection in an historical SN allows to classify it with a good confidence level. \\item The spectroscopic observation of Case A LEs in the vicinity of SN remnants can help in establishisng a firm and direct relation between the SN type and the remnant. \\item The information on the extinction properties one can derive from LEs (both Case A and B) observations may lead to incorrect conclusions on the dust nature. \\item Case B LEs affect in principle all phases, without necessarily producing notable effects on the spectra. This is particularly true if the dust is confined within $\\sim$10$^{16}$ cm from the SN explosion and $\\tau_d<$1. \\item In the presence of circumstellar dust, the light curves tend to become broader with a smaller $\\Delta m_{15}$. The rising time increases, the late time decay becomes slower and the luminosity difference between maximum light and nebular phase tends to decrease. \\item The SN absolute luminosity deduced from the apparent magnitude, colour excess and a standard extinction law is brighter than the real value. \\item The difference in luminosity is of the same order of the correction foreseen by the Pskowski-Phillips relation. This, coupled to the fact that $\\Delta m_{15}$ is decreased, is a possible source of noise in the luminosity-decline rate relation. \\item As shown by \\cite{wang05}, Case B LEs modify the luminosity and colours in such a way that the apparent selective-to-total extinction ratio $R_V$ is significantly decreased. \\item In this scenario, $R_V$ should show a rapid evolution during the first months after explosion and this, together with a non standard colour evolution, should be a simple diagnostic tool to unveil the underlying LE. \\end{enumerate}"
    },
    "0512/gr-qc0512019_arXiv.txt": {
        "abstract": "Time dilation $\\frac{1}{\\sqrt{1-v^2}}$ and relative velocity $v$ are observationally indistinguishable in the special theory of relativity, a duality that carries over into the general theory under Fermi coordinates along a curve (in coordinate-independent language, in the tangent Minkowski space along the curve). For example, on a clock stationary at radius $r$, a distant observer sees time dilation of $\\frac{1}{\\sqrt{1-v^2}}=\\frac{1}{\\sqrt{1-2M/r}}$ under the Schwarzschild metric and sees the clock receding with a relative velocity of $v=\\sqrt{2M/r}$ under the Painlev{\\'e}-Gullstrand free fall metric. Duality implies that during gravitational collapse, the intensifying time dilation observed at the star's center from a fixed radius $r>0$ is indistinguishable (along a curve) from an increasing relative velocity at which the center recedes as seen from any direction, implying a local inflation. ",
        "introduction": "} This note points out that the observational indistinguishability of time dilation $\\frac{1}{\\sqrt{1-v^2}}$ and relative velocity $v$ in the special theory of relativity is a duality which carries over from the special to the general theory of relativity under Fermi coordinates (along a curve).\\footnote{For more information on Fermi coordinates see Nesterov \\cite{Nesterov:1999ix} and Iliev \\cite{Iliev:1992um}.} ",
        "conclusions": ""
    },
    "0512/astro-ph0512212_arXiv.txt": {
        "abstract": "We study prospects of a method to constrain the inclination of a coalescing compact  binary  by detecting its gravitational waves associated with a three-dimensionally  localized  (direction  and  distance) short-hard gamma-ray burst.  We take advantage of   a synergy of these two observations,  and  our method can be   applied   even with a  single interferometer.  For a nearly face-on binary the inclination angle $I$ can be constrained in the range $1-SNR^{-1} \\le \\cos I\\le 1$ ($SNR$: the signal to noise ratio of gravitational wave detection), provided that the error of the distance estimation is negligible. This method would   help us to study   properties of the short-hard bursts, including   potentially collimated jet-like structures as indicated   by recent observation. ",
        "introduction": "The gamma-ray bursts have been  known to be divided into two classes, the long-soft bursts and the short-hard bursts (SHBs).  While the former are likely to be produced at  explosions of massive stars in star forming galaxies typically at high redshift $z\\gsim 1$ \\cite{Zhang:2003uk}, the nature of SHBs has been a long-standing mystery. However,  recent discoveries of X-ray afterglows of SHBs by Swift and HETE satellites allowed us to localize them accurately and rapidly enough  to specify their host galaxies and finally determine their distances \\cite{Gehrels,Villasenor:2005xj,Hjorth:2005kf,Fox:2005kv,Berger:2005rv}.   One of them, GRB 050724  was found in an elliptical galaxy at  $z=0.257$ with an old stellar population \\cite{Hjorth:2005kf}, and   GRB 050509b is likely to be in a similar galaxy at $z=0.225$ \\cite{Berger:2005rv}. While GRB 050709 was in a star forming galaxy at $z=0.160$,  its light curve excluded a supernova association \\cite{Fox:2005kv}. These results support that coalescing compact binaries (double neutron stars (NS+NSs) or black hole-neutron star (BH+NS) systems) are the promising origins of SHBs \\cite{Eichler:1989ve}, though the estimated typical age of these binaries for SHBs are longer than that of known NS+NSs in our galaxy \\cite{Guetta:2005bb,Nakar:2005bs}.  Coalescing compact binaries are also promising sources of gravitational radiation for LIGO and other ground-based interferometers \\cite{Thorne_K:1987}. A one year scientific run (S5) is ongoing with LIGO that has sensitivity to detect NS+NSs to $\\sim 15$Mpc \\cite{ligo}. Recent theoretical analysis predicts that the probability of a simultaneous detection of gravitational waves by LIGO and a SHB by Swift in one year is $\\sim 30$\\% for BH-NS merger and $\\sim 10$\\% for NS+NS (depending on the lower end of the luminosity function of SHBs) \\cite{Nakar:2005bs}. Therefore, we might soon experience the first detection of  gravitational waves associated with a localized SHB.  The observed afterglows of two SHBs showed steeper power-law decays that indicate SHBs have collimated jet-like structures \\cite{Fox:2005kv,Berger:2005rv}, as found with long-soft bursts.  The estimated beaming factor is $\\sim 0.03$ for GRB 050709 \\cite{Fox:2005kv} and $\\sim 0.01$ for GRB 050724 \\cite{Berger:2005rv}. Here we defined the beaming factor as the fraction of $4\\pi$ steradians into which  jets are emitted. This enabled us to estimate their  total energies $\\sim 3\\times 10^{48}$erg  that is smaller than the long-soft bursts by $\\sim 2$ orders of magnitude. If a  SHB is associated with a coalescing binary, it is likely that the orientation of the jet is aligned with the angular momentum of the binary that would be clearly imprinted on the observed gravitational waveform.  In this paper we propose a method with which the inclination of   a binary will be interestingly constrained  as a synergy of three-dimensional localization  by electro-magnetic waves (EMWs) and observation of gravitational waves,  even using a single interferometer.  Therefore, LIGO could provide us an important geometrical information to understand properties of SHBs.  This paper is organized as follows: in \\S II we describe our method to estimate inclination of binaries with using a single gravitational wave interferometer. Expected error for our method is studied in \\S III. Then, in \\S IV, we extend our study for observation with  multiple detectors, such  as, LIGO-VIRGO network. Brief discussions are presented in \\S V. ",
        "conclusions": "We discussed  prospects of  a  method to constrain the inclination angle of a coalescing compact binary  by detecting its gravitational waves associated with a three-dimensionally  localized SHB. With this method we can get an important geometrical information to understand the properties of SHBs and their afterglows as a function of the viewing angles of the jets (see {\\it e.g.} \\cite{Yamazaki:2004ha} for long-soft bursts). We should comment that there would be  a selection effect toward larger $|V|$ for a simultaneous detection of gravitational waves and SHB. This is because the observed gravitational wave amplitude would be larger (see \\S IV.B) and the SHB would be also luminous with the face-on configuration \\cite{Nakar:2005bs,Kochanek:1993mw}.   By analyzing nearby samples with our method we might establish an efficient empirical criteria to select almost face-on binaries using observed properties of  EMW signals.  Then  we can study cosmological parameters with  face-on binaries at relatively high redshift  where cosmological effects beyond the Hubble-law become important.  The following is the outline of this approach. We first obtain the intrinsic amplitudes $A$ by putting  $R=R_{max}$ for the observed gravitational wave amplitude $(AR)$ for binaries that are likely to be nearly face-on according to the empirical criteria. Then the luminosity distances $r$ are obtained from eq.(10) and the estimated chirp masses. In  this manner  we might observationally study the redshift-distance relation with the redshift information naturally obtained from identified host galaxies with EMW observation. Note that this argument to estimate the distance $r$ is different from the previous method in this paper to constrain the inclination angle by simply converting the observed redshift to the distance $r$ for nearby binaries. Considering the strong correlation between the inclination $V$ and the distance $r$ for parameter estimation \\cite{Cutler:1994ys},  this  could be a powerful approach to investigate the dark energy with  future gravitational wave detectors that can detect compact binaries at $z\\sim 1$ (see also \\cite{Schutz:gp,Dalal:2006qt,Finn:1995ah})."
    },
    "0512/astro-ph0512024_arXiv.txt": {
        "abstract": "{  We have used the FORS-1 camera on the VLT to study the main sequence (MS) of the globular cluster NGC 6218 in the $V$ and $R$ bands. The observations cover an area of $3\\farcm4 \\times 3\\farcm4$ around the cluster centre and probe the stellar population out to the cluster's half-mass radius ($r_{\\rm h} \\simeq 2\\farcm2$). The colour-magnitude diagram (CMD) that we derive in this way reveals a narrow and well defined MS extending down to the $5\\, \\sigma$ detection limit at $V \\simeq 25$, or about $6$ magnitudes below the turn-off, corresponding to stars of $\\sim 0.25$\\,\\Msolar. The luminosity function (LF) obtained with these data shows a marked radial gradient, in that the ratio of lower- and higher-mass stars increases monotonically with radius. The mass function (MF) measured at the half-mass radius, and as such representative of the cluster's global properties, is surprisingly flat. Over the range $0.4 - 0.8$\\,\\Msolar, the number of stars per unit mass follows a power-law distribution of the type $dN/dm \\propto m^{0}$, where, for comparison, Salpeter's IMF would be $dN/dm \\propto m^{-2.35}$. We expect that such a flat MF does not represent the cluster's IMF but is the result of severe tidal stripping of the stars from the cluster due to its interaction with the Galaxy's gravitational field. Our results cannot be reconciled with the predictions of recent theoretical models that imply a relatively insignificant loss of stars from NGC\\,6218 as measured by its expected very long time to disruption. They are more  consistent with the orbital parameters based on the Hipparcos reference system that imply a much higher degree of interaction of this cluster with the Galaxy than assumed by those models. Our results indicate that, if the orbit of a cluster is known, the slope of its MF could be useful in discriminating between the various models of the Galactic potential. ",
        "introduction": "A satisfactory understanding of the properties of the initial mass function (IMF) of globular clusters (GCs) is a major objective of current astrophysical research in that GCs are the closest example of star formation at high redshift (Krauss \\& Chaboyer 2003). The most reliable observations so far available that reach near the bottom of the stellar MS indicate that all halo clusters have a very similar present global MF (PGMF), which peaks at $\\sim 0.35$\\,\\Msolar (Paresce \\& De Marchi 2000; De Marchi, Paresce \\& Portegies Zwart 2005). The internal dynamical relaxation process via stellar encounters is now reasonably well understood (Meylan \\& Heggie 1997) and validated observationally (e.g. De Marchi, Pulone \\& Paresce 2000; Albrow, De Marchi \\& Sahu 2002; Pasquali et al. 2004) and allows us to derive the global properties of the MF from a limited number of measurements within a cluster.  Nevertheless, any hope to infer useful information on the properties of the IMF from the observed present global MF (PGMF) rests on our ability to roll back the effects that the tidal field of the Galaxy has exerted on the stellar population of the clusters (Paresce \\& De Marchi 2000). Gravitational shocking due to repeated  interactions with the bulge and disc of the Galaxy profoundly disrupt the original mass distribution by ejecting low mass stars from the core and by compressing the tidal boundary in phase space at each encounter. These phenomena, integrated over the orbit and time and eventually causing the disruption of the cluster, can substantially alter the shape of the MF thereby completely masking the properties of the IMF (Vesperini \\& Heggie 1997).  Theoretical models describing the interaction of GCs with the Galactic tidal field have become progressively more detailed and, possibly, accurate in the past decade or so. This has been made possible by an in-depth analysis of the mechanisms responsible for cluster disruption (Aguilar, Hut \\& Ostriker 1988; Gnedin \\& Ostriker 1997) and by the availability of more accurate space motion parameters for the clusters (Dauphole et al. 1996; Odenkirchen et al. 1997). Models that make use of GC proper motion information (Dinescu et al. 1999; Baumgardt \\& Makino 2003) should, in principle, provide a more reliable description of the clusters' past dynamical history than those purely based on radial velocity data (Gnedin \\& Ostriker 1997). On the other hand, the models of Gnedin \\& Ostriker (1997) explain rather convincingly, albeit in a statistical sense, why the GCs that we see today are not randomly distributed in parameter space but rather occupy regions of low probability of disruption.  The most useful indicator of the past dynamical history of GCs that models of this type produce is the time to disruption, $T_{\\rm d}$, namely the time over which a cluster would be completely dissolved by tidal forces. Unfortunately, this parameter is not directly observable, thereby making it more difficult to assess the validity of the models. It is, therefore, necessary to relate $T_{\\rm d}$ to other measurable parameters.  A rather obvious indication of a cluster's tidal disruption would seem to be the presence of an extended tidal tail (Leon, Meylan \\& Combes 2000; Odenkirchen et al. 2003). The unequivocal detection of tidal tails, however, is hard to achieve on the basis of photometric information alone (Baumgardt \\& Kroupa 2005), which is often the only available data. A more robust, and potentially more powerful approach consists in looking at the properties of the MF of MS stars and at its variations within the cluster, with respect to that of other well behaved reference GCs.  This has allowed us to identify, for the first time, the clear signature of tidal disruption in the cluster, NGC\\,6712, in the form of a severe  depletion of low-mass stars (De Marchi et al. 1999; Andreuzzi et al. 2001).  While our discovery has proved that GCs are indeed subject to the effects of tidal stripping, the widely different values of $T_{\\rm d}$ predicted by various models for NGC\\,6712 have, at the same time, raised  concerns as to the validity of the present physical understanding of the processes involved. In order to better understand the possible magnitude of the underlying discrepancy, we have studied the properties of the MF of another GC, NGC\\,6218, which, according to all three sets of presently available models (Gnedin \\& Ostriker 1997; Dinescu et al. 1999; Baumgardt \\& Makino 2003) should have experienced an insignificant or very mild interaction with the Galactic tidal field.  We would expect that its PGMF should accurately reflect the IMF and our  aim was to use NGC\\,6218 as a reference for NGC\\,6712 and other clusters.  In this paper we report on the properties of its MF that shows that there may be something terribly wrong with these models or, more likely, with their assumptions.  The structure of the paper is as follows. The observations and their reduction are described in Section\\,2, whereas the results of the photometry are discussed in Section\\,3. Section\\,4 is devoted to the LF and MF of MS stars at various locations in the cluster. The dynamical structure of NGC\\,6218, as derived from these data, is presented in Section\\,5 and the overall implications for the understanding of the interaction between GCs and the potential field of the Galaxy are discussed in Section\\,6. ",
        "conclusions": "What does a flat PGMF tell us about NGC\\,6218? Using a sample of 12 halo GCs for which deep HST observations are available, Paresce \\& De Marchi (2000) showed that their PGMF must directly reflect the properties of the IMF. This conclusion is based on the observation that those clusters have very different properties (total mass, metallicity, concentration and space  motion parameters) but they show, in practice, the same or very similar PGMF. This would be hard to justify if the clusters were initially born with very different IMFs. The underlying common IMF, exemplified by that of NGC\\,6397, is best represented by a power-law distribution that tapers off below $\\sim 0.3$\\,\\Msolar (see Paresce \\& De Marchi 2000 and De Marchi et al. 2005 for details).  The PGMF of NGC\\,6218 clearly does not match that of the 12 objects in the sample of Paresce \\& De Marchi (2000), because it is remarkably flat. However, at least two other cases of flat or even dropping PGMF have been reported, namely those of NGC\\,6712 (De Marchi et al. 1999) and Pal\\,5 (Koch et al. 2004), with NGC\\,6712 revealing an even stronger deficit of low-mass stars than NGC\\,6218. On the other hand, there are good observational reasons to believe that both NGC\\,6712 and Pal\\,5 have suffered severe tidal disruption  that has considerably altered their original distribution of stellar  masses. For instance, Pal\\,5 has well defined tidal tails extending over $10^\\circ$ across the sky (Odenkirchen et al. 2001). But, more generally, the present total mass and space motion parameters of these clusters imply that they have some of the highest destruction rates in the whole Galactic GC system. Gnedin \\& Ostriker (1997) predict a remaining lifetime as low as $\\sim 0.3$\\,Gyr for NGC\\,6712 and $\\sim 1$\\,Gyr for Pal\\,5, whereas Dinescu et al. (1999) give a time to disruption of $\\sim 3.9$\\,Gyr for NGC\\,6712 and just $\\sim 0.1$\\,Gyr for Pal\\,5, respectively. Although not in agreement with one another, these sets of values are far lower than the average time to disruption, which for both authors is of order 12\\,Gyr (and thus comparable with the typical GC age).     Quite surprisingly, however, the total mass and space motion parameters of NGC\\,6218 do not seem to put this cluster in any imminent danger. In fact, the estimated time to disruption for this object varies from $\\sim 23.5$\\,Gyr (Gnedin \\& Ostriker 1997) to $\\sim 33$\\,Gyr (Dinescu et al. 1999). Similarly, collisional N-body simulations of a cluster with the properties of NGC\\,6218 moving in an external tidal field produce a total lifetime for this object of $\\sim 29$\\,Gyr (H. Baumgardt, private comm.; see Baumgardt \\& Makino 2003 for model details). Assuming a GC age of $\\sim 12.5$\\,Gyr (Krauss \\& Chaboyer 2003), this implies $T_{\\rm d} \\simeq 16.5$\\,Gyr for NGC\\,6218. In other words, none of these models implies a dynamically troubled past for NGC\\,6218, leaving in principle open the possibility that it was born with a rather flat IMF, at least in the mass range covered by our observations. While such a hypothesis cannot {\\em a priori} be excluded, this would be the first case known.  On the other hand, models of the dynamical interaction of GCs with the Galactic tidal field are, unfortunately, still subject to large uncertainties. As both Gnedin \\& Ostriker (1997) and Dinescu et al. (1999) point out, different assumptions on the initial conditions (cluster orbits) and  on the Galactic potential can result in rather different destruction rates for the same cluster. While these models are useful to address the replenishment of the halo over time from the disruption of individual clusters, and therefore helpful to understand the global properties of the GC system and its evolution, they may be still too crude to precisely describe the past history of individual clusters.  We, however, believe that the apparent discrepancy between the predicted value of $T_{\\rm d}$ and the shape of the PGMF of NGC\\,6218 is simply due to the wrong assumption as to the cluster orbit. In fact, the space motion parameters used by these authors for NGC\\,6218 are not consistent with the latest determination of its absolute proper motion based on the Hipparcos reference system. The  proper motion measured by Brosche et al. (1991), combined with the radial velocity measurements of Pryor \\& Meylan 1993, allowed Dauphole et al. (1996) to put some constraints on the space motion parameters of this cluster. Their findings were confirmed by an independent analysis of the orbit, based on improved absolute proper motions (Scholz et al. 1996), which indicated that NGC\\,6218 should have a short orbital period ($0.17$\\,Gyr) but also that it never ventures closer than $\\sim 3$\\,Kpc from the Galactic centre, with less than 15\\,\\% of its orbit lying within 1\\,Kpc of the Galactic plane. However, a more recent study of the orbit of NGC\\,6218, based on the new Hipparcos reference system (ESA 1997), has led Odenkirchen et al. (1997) to the conclusion that NGC\\,6218 has a highly irregular motion. In particular, they find that the low value of the cluster's axial angular momentum forces it to pass the Galactic centre at short distance ($R_{\\rm p} = 0.6$\\,kpc) and  to get into strong interaction with the Galactic bulge, since the bulge destruction rate scales with the fourth power of $R_{\\rm p}$ (Dinescu et al. 1999).  With an orbit of this type, the total lifetime of NGC\\,6218 predicted by the models of Baumgardt \\& Makino (2003; H. Baumgardt priv. comm.) would decrease from 29\\,Gyr to 17\\,Gyr, thus implying a time until disruption of only $T_{\\rm d} \\simeq 4.5$\\,Gyr, assuming a typical GC age of $\\sim 12.5$\\,Gyr as above (Krauss \\& Chaboyer 2003). Although not as small as that of NGC\\,6712 or Pal\\,5, this value of $T_{\\rm d}$ places NGC\\,6218 among the clusters at higher risk of disruption and suggests that a considerable fraction of its original stellar population should have been stripped from the cluster.  We can estimate the amount of mass lost by NGC\\,6218 in the hypothesis mentioned above that all GCs were born with a very similar IMF (Paresce \\& De Marchi 2000). As for the latter, we use the tapered power-law proposed by De Marchi et al. (2005) and find that the present total mass due to MS stars is about one fifth of the original $\\sim 6 \\times 10^5$\\,\\Msolar. This value is in full agreement with the revised calculations of H. Baumgardt (priv. comm.) suggesting an initial mass of $7.6 \\times 10^5$\\,\\Msolar. A natural question to pose is when this major mass loss process took place.  If the cluster is in thermo-dynamical equilibrium, one only needs to look at the half-mass relaxation time in order to answer this question. The data are compatible with this hypothesis, since our model imposes the condition of energy equipartition and Figure\\,6 and Table\\,5 show that this is consistent with the radial variation of the MF and the cluster's structural parameters.\\footnote{We are grateful to an anonymous Referee for pointing out that, although plausible, thermo-dynamical equilibrium is not necessarily implied by the data. In fact, the good fit shown in Figure\\,6 and Table\\,5 is a necessary condition for energy equipartition, but it is not a sufficient one. In principle, there could be models that fit the data equally well but in which the equipartition of energy is not complete. Unfortunately, the method that we followed (Meylan 1987; 1988) does not allow us to investigate this hypothesis.} Under this assumption, the half-mass relaxation time suggested by our models is  $t_{\\rm rh}\\simeq 0.5$\\,Gyr. This value is slightly lower than the $t_{\\rm rh} \\simeq 0.7$\\,Gyr given by Djorgovski (1993) and $t_{\\rm rh}  \\simeq 1$\\,Gyr of Gnedin \\& Ostriker (1997), although still compatible with them, because the MF assumed by these authors is very steep at the low mass end, contrary to what we find, thus resulting in a much larger number of objects in the cluster. In any case, it seems possible to exclude that a major mass loss episode happened in the course of the past 1\\,Gyr or so, since in that case the cluster should have not yet reached a condition of equilibrium. The fact that tidal tails are very tenuous around NGC\\,6218 (if at all present; see Lehmann \\& Scholz 1997) gives support to this scenario. Mass loss must, therefore, have happened either long ago or continuously, over the cluster lifetime, at a low rate of about $5 \\times 10^3$\\,\\Msolar per orbit.  The present data do not allow us to explore the past dynamical history of NGC\\,6218 in more detail. Nevertheless, our results show the importance of feeding models of GC distruption with reliable cluster orbits, since the strength and extent of the interaction between the Galaxy and GCs can vary dramatically with the orbit. However, it also depends on the gravitational potential of the Galaxy, and particularly on that of the  bulge, thereby making it crucial for models of this type to rely on a solid observational description of the structure of the Galaxy and of its components (thin/thick disc, bulge, halo), which is presently lacking. This information should become available in the near future with the advent of missions such as SIM and Gaia. But already now, if the cluster orbit is reasonably well understood, observations of the cluster PGMF can set meaningful constraints on the mass distribution in the Galaxy. Indeed, the PGMF is a more reliable indicator of the past dynamical history of a cluster than, for instance, its present location and space motion parameters or even the presence and extent of its tidal tails. The position and velocity of a cluster are instantaneous quantities and can vary largely in time. Tidal tails are made up of unbound stars and, as such, are short lived and can only probe the immediate past of a cluster. The PGMF, on the other hand, reflects the integrated effect of the interaction with the Galaxy throughout the whole life of the cluster and provides an indication of the amount of mass lost to the Galaxy. By measuring the PGMF of a sizeable number of clusters, including those with chaotic or anyhow irregular orbits, it should be possible to set meaningful constraints on the form of the Galactic potential well before the availability of astrometric data from interferometric space observatories."
    },
    "0512/astro-ph0512448_arXiv.txt": {
        "abstract": "Spectral methods are well suited for solving hydrodynamic problems in which the self-gravity of the flow needs to be considered. Because Poisson's equation is linear, the numerical solution for the gravitational potential for each individual mode of the density can be pre-computed, thus reducing substantially the computational cost of the method.  In this second paper, we describe two different approaches to computing the gravitational field of a two-dimensional flow with pseudo-spectral methods.  For situations in which the density profile is independent of the third coordinate (i.e., an infinite cylinder), we use a standard Poisson solver in spectral space.  On the other hand, for situations in which the density profile is a delta function along the third coordinate (i.e., an infinitesimally thin disk), or any other function known a priori, we perform a direct integration of Poisson's equation using a Green's functions approach.  We devise a number of test problems to verify the implementations of these two methods.  Finally, we use our method to study the stability of polytropic, self-gravitating disks. We find that, when the polytropic index $\\Gamma$ is $\\le 4/3$, Toomre's criterion correctly describes the stability of the disk. However, when $\\Gamma > 4/3$ and for large values of the polytropic constant $K$, the numerical solutions are always stable, even when the linear criterion predicts the contrary.  We show that, in the latter case, the minimum wavelength of the unstable modes is larger than the extent of the unstable region and hence the local linear analysis is inapplicable. ",
        "introduction": "\\label{sec:introduction}  In the standard model of accretion disks, turbulent viscosity plays an important role in bringing material inward and transporting angular momentum outward \\citep[see][]{Frank2002}.  At the same time, viscous dissipation converts gravitational potential energy to thermal energy and heats up the disk, which then radiates away this energy as thermal emission.  In most applications of accretion disks around central objects, as in, e.g., X-ray binaries, the self-gravity of the flow is negligible compared to that of the central object.  However, there are disk-like systems, such as active galactic nuclei as well as protostellar and protoplanetary disks, where the effects of self-gravity change not only the properties of angular momentum transport \\citep{Boss1998, Balbus1999, Mejia2005}, but also the energy balance equation \\citep{Bertin1997,Bertin1999,Bertin2001}, which affect the global structure of the disk.  Besides the angular momentum transport problem, self-gravitating disks are also very important in studying star and planet formation. Indeed, gravitational instabilities in protoplanetary disks have been proposed as viable planet formation mechanisms.  Although there has been a large amount of work done on gravitational instabilities in these system \\citep[see, for example,][and references therein]{Pickett2003,Mejia2005}, it still remains an open question whether the fragmentation by gravitational instabilities can produce bound planetary objects, or if reservoirs of small solid cores are needed for the accretion mechanism to lead to rapid planet formation.  In the first paper in this series (Chan, Psaltis, \\& \\\"Ozel 2005), we presented a pseudo-spectral method for solving the equations that describe the evolution of two dimensional, viscous hydrodynamic flows. There, we addressed issues related to the implementation of boundary conditions, spectral filtering, and time-stepping in spectral methods, and verified our algorithm using a suite of hydrodynamic test problems.  In this second paper of the series, we present our implementation of a Poisson solver that allows us to take into account the effects of self-gravity of the flow.  Spectral methods are particularly suitable for incorporating the effects of self-gravity.  Because Poisson's equation is linear, we can pre-compute the numerical solution for the gravitational potential for each individual mode of the density, thus reducing substantially the computational cost of the method.  In fact, Fourier methods have been used extensively in analytical studies of gravitational potentials in systems with periodic boundary conditions \\citep{Binney1987}.  On the other hand, in numerical studies of self-gravitating disks, the strengths of spectral methods have only been partially incorporated.  Hybrid hydrodynamic algorithms with self-gravity have been developed in such a way that modified spectral methods are used only for solving Poissson's equation for the gravitational field, whereas the hydrodynamic parts are still treated with finite difference schemes. For example, \\citet{Boss1992} describe a spherical harmonic decomposition method and a second-order scheme in the radial direction to solve Poisson's equation, whereas \\citet{Myhill1993} use a modified Fourier method to find the gravitational potential of an isolated distribution of sources.  Both algorithms employ explicit second-order finite difference methods to advance the hydrodynamic equations. \\citet{Pickett1998,Pickett2000} describe an implementation of an algorithm that uses Fourier decomposition in the azimuthal direction of cylindrical coordinate to solve Poisson's equation together with a von Neumann \\& Richtmeyer AV scheme for the hydrodynamics.  Some other examples can be found in Grandclement et al. (2001), Broderick \\& Rathore (2004), and Dimmelmeier et al. (2005).  The main advantage of using hybrid methods is that one can employ currently available hydrodynamic algorithms based on finite difference schemes.  However, a hybrid method does not exploit the high order of the spectral algorithm, because the hydrodynamic difference schemes typically have a much lower order compared to that of the Poisson solver.  Contrary to these efforts, our algorithm uses a spectral decomposition method for solving both the hydrodynamic and Poisson's equation, providing a consistent treatment of the whole problem.  To our knowledge, this is the first time that spectral methods have been used in studying astrophysical disks with self gravity.  By construction, there is an ambiguity in the definition of the gravitational field in two-dimensional problems.  We can assume either that the density profile is independent of the third coordinate (i.e., an infinite cylinder) or that it is a delta (or any other predetermined) function along the third coordinate (i.e., an infinitesimally thin disk); the resulting gravitational field on the two-dimensional domain of solution will be different in the two cases.  For example, the gravitational potential of a ``point source'' on the two-dimensional domain of solution will be proportional to $\\log(r)$ in the first case and to $1/r$ in the second, where $r$ is the distance from the source.  In order to consider both geometries, here we describe two different approaches to computing the gravitational field of a two-dimensional hydrodynamic flow with pseudo-spectral methods.  When the flow has the geometry of an infinite cylinder, we use a standard two-dimensional pseudo-spectral Poisson solver, which has been proven to be numerically stable and accurate. When the flow has the geometry of an infinitesimally thin disk, we perform a direct integration of the Green's function for the gravitational potential, following the work of Cohl \\& Tohline (1999).  In the following section, we present our assumptions and equations. In \\S\\ref{sec:numerical_methods}, we discuss the details of our numerical methods, include both the standard Poisson solver and the Green's function integrator.  Next, we present a series of tests in \\S\\ref{sec:tests}, to verify our algorithm.  Finally, we apply our method to a numerical study of Toomre's stability criterion of self-gravitating disks in \\S5. ",
        "conclusions": "Problems that involve self-gravity are usually time-consuming tasks in computational physics.  In standard finite difference methods, Poisson's equation is solved as the steady state of a diffusion equation (using relaxation and over-relaxation methods) and has to be recalculated together with the hydrodynamic equations at every timestep.  Although hybrid algorithms have been developed which use high-order methods to solve Poisson's equation, finite difference schemes are still used for evolving the hydrodynamic equations.  The resulting inconsistency in the order of differencing the hydrodynamic equations and Poisson's equations either significantly reduces the accuracy of the Poisson solver, or requires the extra complication of interpolating the gravitational filed at each grid point.  More importantly, existing two-dimensional hybrid algorithms can only address a limited number of self-gravitating problems because the solutions to the two-dimensional and three-dimensional Poisson's equation are fundamentally different (see \\S\\ref{sec:introduction} and \\S\\ref{sec:equations_assumptions}).  In this paper, we presented two different approaches to using pseudo-spectral methods to solve self-gravity problems.  For the first approach, we described the implementation of a standard pseudo-spectral Poisson solver that solves the two-dimensional Poisson's equations to machine accuracy.  Instead of solving a diffusion equation, spectral methods allowed us to invert the Poisson operator in spectral space, making the algorithm fast and accurate. For the second approach, we investigated a fast gravity integrator for disks-like flows with known, time-independent vertical structures. This algorithm allowed us to study the evolution of flows with finite, but not only infinitesimal, thickness.  This improvement allows two-dimensional algorithms to solve a whole new class of problems.  We demonstrated here the ability of our algorithm to compute properly and efficiently the gravitational potential of flattened flows using different test problems.  Even for the high resolution that corresponds to $129\\times256$ collocation points, our Poisson solvers and integrator use less than 10\\% of the total computational time.  We also explored how to extend Toomre's stability criterion to self-gravitating disks for which the characteristic properties of the flows change over a length scale that is shorter than the minimum unstable wavelength.  Based on our simulations, we find that for Plummer's density model, if the disk is hot and has a polytropic index $\\Gamma > 4/3$, all oscillatory modes in the disk are stable, contrary to the predictions of the analytic calculation."
    },
    "0512/astro-ph0512413_arXiv.txt": {
        "abstract": "We report extensive radio and X-ray observations of SN\\,2003bg whose spectroscopic evolution shows a transition from a broad-lined Type Ic to a hydrogen-rich Type II and later to a typical hydrogen-poor Type Ibc.  We show that the extraordinarily luminous radio emission is well described by a self-absorption dominated synchrotron spectrum while the observed X-ray emission at $t\\approx 30$ days is adequately fit by Inverse Compton scattering of the optical photons off of the synchrotron emitting electrons.  Our radio model implies a sub-relativistic ejecta velocity, $\\overline{v}\\approx 0.24 c$, at $t_0\\approx 10$ days after the explosion which emphasizes that broad optical absorption lines do not imply relativistic ejecta.  We find that the total energy of the radio emitting region evolves as $E\\approx 7.3\\times 10^{48} (t/t_0)^{0.4}$ erg assuming equipartition of energy between relativistic electrons and magnetic fields ($\\epsilon_e=\\epsilon_B=0.1$).  The circumstellar density is well described by a stellar wind profile with modest (factor of $\\sim 2$) episodic density enhancements which produce abrupt achromatic flux variations.  We estimate an average mass loss rate of $\\dot{M}\\approx 3 \\times 10^{-4}~\\rm M_{\\odot}~yr^{-1}$ (assuming a wind velocity of $v_w=10^3~\\rm km~s^{-1}$) for the progenitor, consistent with the observed values for Galactic Wolf-Rayet stars.  Comparison with other events reveals that $\\sim 50\\%$ of radio supernovae show similar short timescale flux variations attributable to circumstellar density irregularities.  Specifically, the radio light-curves of SN\\,2003bg are strikingly similar to those of the Type IIb SN\\,2001ig, suggestive of a common progenitor evolution for these two events.  Based on the relative intensity of the inferred density enhancements, we conclude that the progenitors of SNe 2003bg and 2001ig experienced quasi-periodic mass loss episodes just prior to the SN explosion. Finally, this study emphasizes that abrupt radio light-curve variations cannot be used as a reliable proxy for an engine-driven explosion, including off-axis gamma-ray bursts. ",
        "introduction": "\\label{sec:intro}  Accounting for $\\sim 10\\%$ of the nearby supernova population, Type Ibc supernovae (hereafter SNe Ibc) are identified by their lack of spectroscopic hydrogen and silicon features (see \\citealt{f97} for a review).  Recognized as a rare subclass of core-collapse supernova $\\sim 20$ years ago \\citep{emn+85}, SNe Ibc have recently enjoyed a revitalized interest. Beginning with the discovery of Type Ic SN\\,1998bw in temporal and spatial coincidence with gamma-ray burst (GRB)\\,980425 \\citep{gvv+98,paa+00}, we now know that most long-duration GRBs (e.g. \\citealt{smg+03}) and X-ray flashes (XRFs; \\citealt{skf+05}) are associated with SNe Ibc. However, radio observations of local SNe Ibc indicate that the inverse is {\\em not} true; there is a strict limit of $\\lesssim 10\\%$ on the fraction of SNe Ibc that could be accompanied by a GRB or XRF \\citep{snk+05}.  The lack of hydrogen features in both SNe Ibc and GRB/XRF-associated SNe imply that they represent the explosion of massive stripped-core progenitors \\citep{wl85,wes99}.  The popular models for SNe Ibc progenitors include massive Wolf-Rayet (WR) stars that eject their envelopes through strong dense winds, and close binary systems where the progenitor star is stripped of its hydrogen-rich layer by the companion \\citep{ew88}. Despite dedicated archival searches, the progenitors of SNe Ibc are still poorly constrained by pre-explosion images.  The best photometric constraints are currently associated with SN\\,2002ap and SN\\,2004gt which exclude WR progenitors in the top $\\sim 30\\%$ and $\\sim 50\\%$ of the population, respectively \\citep{svr+02,gfk+05,mss05}.  While intriguing, these archival observations lack the sensitivity to clearly discriminate between the WR and binary progenitor models.  By studying the circumstellar medium around the supernova, however, it is possible to place independent constraints on the progenitor. Observations show that Galactic Wolf-Rayet stars are embedded in wind stratified media \\citep{chu02} while binary systems are associated with disrupted circumstellar media, possibly including an outflow during a common envelope phase \\citep{pjh92}.  Radio observations provide the most direct probe of the circumstellar density structure around supernovae \\citep{wsp+86,c98}. By observing the dynamical interaction of the ejecta with the surrounding medium we are able to map out the mass loss history of the progenitor star.  The identification of irregular density profiles may therefore distinguish between Wolf-Rayet and binary progenitor systems.  Here we present extensive radio observations of SN\\,2003bg, discovered as part of our ongoing radio survey of local Type Ibc supernovae.  The peculiar SN\\,2003bg spectroscopically evolved from a broad-lined Type Ic to a hydrogen-rich Type II and later to a typical Type Ibc (Hamuy {\\it et al.}, in prep), thereby bridging the hydrogen-rich and poor divisions of the core-collapse classification system. Our densely-sampled radio light-curves show that the extraordinarily luminous radio emission for a SN Ibc is characterized by episodic short-timescale variations. We show that these variations are well described by abrupt density enhancements in the circumstellar medium. Comparison with other radio supernovae shows that $\\sim 50\\%$ of all well-studied events similarly show evidence for abrupt light-curve variations. We review the mass loss evolution observed (and inferred) for the progenitors of these core-collapse SNe.  Using the observed radio properties for the peculiar SN\\,2003bg, we place constraints on the nature of its progenitor system.  The organization of this paper is as follows: observations from the Very Large Array (VLA) and the {\\it Chandra} X-ray Observatory (CXO) are described in \\S\\ref{sec:obs}.  Preliminary estimates of the energy, velocity and density of the radio emitting region are presented in \\S\\ref{sec:ep}.  Modeling of the radio light-curves are presented in \\S\\ref{sec:model} and \\S\\ref{sec:ssa} while a discussion of the radio polarization mechanism follows as \\S\\ref{sec:pol_disc}.  Our modeling of the X-ray emission is discussed in \\S\\ref{sec:xray_model}. In \\S\\ref{sec:bumps} we present a compilation of radio supernovae with abrupt light-curve variations and review their circumstellar irregularities.  Finally, in \\S\\ref{sec:prog} we discuss the possible causes of the density enhancements surrounding SN\\,2003bg and the implications for the progenitor system. ",
        "conclusions": "\\label{sec:conc} We report extensive radio and X-ray observations of SN\\,2003bg.  The spectroscopic evolution of this unusual supernova shows a transition from a broad-lined Type Ic to a hydrogen-rich Type II and later to a typical hydrogen-poor Type Ibc.  This evolution strengthens the connection between Type II and Type Ibc (including broad-lined) events.  We show that the extraordinarily luminous radio emission for a SN Ibc is well described by a self-absorption dominated synchrotron spectrum while the observed X-ray emission at $t\\approx 30$ days is adequately fit by inverse Compton scattering of the optical photons off of the synchrotron emitting electrons.  Our radio model implies a sub-relativistic ejecta velocity, $\\overline{v}\\approx 0.24 c$, and a size of $r\\approx 6.2\\times 10^{15}$ cm at $t_0\\approx 10$ days.  This analysis emphasizes that broad optical absorption lines do not imply relativistic ejecta.  We find that the total energy of the radio emitting region evolves as $E\\approx 7.3\\times 10^{48} (t/t_0)^{0.4}$ erg assuming equipartition of energy between relativistic electrons and magnetic fields ($\\epsilon_e=\\epsilon_B=0.1$).  The circumstellar density is well described by stellar wind profile, $\\propto r^{-2}$, with modest (factor of $\\sim 2$) episodic enhancements which produce abrupt achromatic flux variations.  We show that free-free absorption does not contribute significantly to the radio spectrum and estimate an average mass loss rate of $\\dot{M}\\approx 3 \\times 10^{-4}~\\rm M_{\\odot}~yr^{-1}$, consistent with observed values for local Wolf-Rayet stars.  Comparison with other events reveals that $\\sim 50\\%$ of radio supernovae show similar short timescale light-curve variations which are attributable to circumstellar density irregularities.  This compilation emphasizes that abrupt radio light-curve variations cannot be used as a reliable proxy for an engine-driven explosion, including off-axis gamma-ray bursts.  Finally, the radio light-curves and spectroscopic evolution for SN\\,2003bg are strikingly similar to those of SN\\,2001ig, suggestive of a common progenitor evolution for these two events.  The overall similarity of SNe 2003bg and 2001ig to radio SNe Ibc is suggestive of a compact Wolf-Rayet progenitor model.  Based on the relative intensity of the inferred density enhancements, we conclude that the progenitors of hydrogen-poor SNe 2003bg and 2001ig experienced quasi-periodic mass loss episodes just prior to the SN explosion."
    },
    "0512/astro-ph0512275.txt": {
        "abstract": "High energy neutrino emission from GRBs is discussed. In this paper, by using the simulation kit GEANT4, we calculate proton cooling efficiency including pion-multiplicity and proton-inelasticity in photomeson production. First, we estimate the maximum energy of accelerated protons in GRBs. Using the obtained results, neutrino flux from one burst and a diffuse neutrino background  are evaluated quantitatively. We also take account of cooling processes of pion and muon, which are crucial for resulting neutrino spectra. We confirm the validity of analytic approximate treatments on GRB fiducial parameter sets, but also find that the effects of multiplicity and high-inelasticity can be important on both proton cooling and resulting spectra in some cases. Finally, assuming that the GRB rate traces the star formation rate, we obtain a diffuse neutrino background spectrum from GRBs for specific parameter sets. We introduce the nonthermal baryon-loading factor, rather than assume that GRBs are main sources of UHECRs. We find that the obtained neutrino background can be comparable with the prediction of Waxman \\& Bahcall, although our ground in estimation is different from theirs. In this paper, we study on various parameters since there are many parameters in the model. The detection of high energy neutrinos from GRBs will be one of the strong evidences that protons are accelerated to very high energy in GRBs. Furthermore, the observations of a neutrino background has a possibility not only to test the internal shock model of GRBs but also to give us information about parameters in the model and whether GRBs are sources of UHECRs or not. ",
        "introduction": "Introduction} Gamma-ray bursts (GRBs) are the most powerful explosions in the universe. The observed isotropic energy can be estimated to be larger than ${10}^{52}  \\, \\mr{ergs}$ \\cite{Blo1,Fra1}. The high luminosity and the rapid time variability lead to the compactness problem. This problem and the hardness of observed photon spectra imply that $\\gamma$-ray emission should be results of dissipation of kinetic energy of relativistic expanding shells. In the standard model of GRBs, such a dissipation is caused by internal shocks - internal collisions among the shells (see reviews e.g., \\cite{Pir1,Zha1,Pir2}). Internal shock scenario requires a strong magnetic field, typically $10^{4} - 10^{7} \\, \\mr{G}$, which can accelerate electrons to high energy enough to explain observed $\\gamma$-ray spectra by synchrotron radiation. Usually, Fermi acceleration mechanism is assumed to be working not only for electrons but also for protons, which can be accelerated up to a high energy within the fiducial GRB parameters. Physical conditions allow protons accelerated to greater than ${10}^{20} \\, \\mr{eV}$ and energetics can explain the observed flux of ultra-high-energy cosmic rays (UHECRs), assuming that similar energy goes into acceleration of electrons and protons in the shell \\cite{Wax1,Vie1,Wax2,Bah1}. Such protons accelerated to the ultra-high energy cannot avoid interacting with GRB photons. This photomeson production process can generate very high energy neutrinos and gamma rays \\cite{Wax3}. Whether GRBs are sources of UHECRs or not, internal shock models predict the flux of very high energy neutrinos at the Earth \\cite{Der1,Gue1,Asa1}. Many authors have investigated neutrino emission from GRBs. Ice \\v{C}herenkov detectors such as AMANDA at the South Pole have already been constructed and are taking data \\cite{And2,Ahr1,Har1}. Now, the future $1\\, {\\mr{km}}^3$ detector, IceCube is being constructed \\cite{Alv1,Ahr2,Ahr3}. In the Mediterranean Sea, ANTARES and NESTOR are under construction \\cite{Kat1}. If the prediction is correct, these detectors may detect these neutrinos correlated with GRBs in the near future.\\\\ In this paper, we execute the Monte Carlo simulation kit GEANT4 \\cite{Ago1} and simulate the photomeson production that causes proton energy loss. As a result, we can get meson spectrum and resulting neutrino spectrum from GRBs quantitatively. In our calculation, we take into account pion-multiplicity and proton-inelasticity, which are often neglected or approximated analytically in previous works, although they may be important in some cases \\cite{Muc1,Muc2}. We also take into account the synchrotron loss of mesons and protons. These cooling processes play a crucial role for the resulting neutrino spectrum. Our models and method of calculation are explained in Sec. \\ref{sec:level2}. One of our purposes is to seek physical conditions allowing proton to be accelerated up to ultra-high energy. Similar calculations are carried by Asano \\cite{Asa1}, in which the possibility that a nucleon creates pions multiple times in the dynamical time scale is included but multiplicity is neglected. Using obtained results, we also investigate efficiency of neutrino emission for various parameter sets. Since there are many model uncertainties in GRBs, such a parameter survey is meaningful. Our final goal is to calculate a diffuse neutrino background from GRBs for specific parameter sets. A unified model for UHECRs from GRBs is very attractive \\cite{Wax1,Vie1,Wax3}, although it requires GRBs being strongly baryon-loaded \\cite{Wic1}. Even if observed UHECRs are not produced mainly by GRBs, high energy neutrino emission from GRBs can be expected. Hence, we leave the nonthermal baryon-loading factor as a parameter. In this paper, assuming GRBs trace star formation rate, we calculate a neutrino background from GRBs and compare our results with the flux obtained by Waxman \\& Bahcall \\cite{Wax2}. The design characters of neutrino detectors are being determined in part by the best available theoretical models, so numerical investigation of high energy neutrino fluxes should be important. Finally, we will consider the implications of neutrino observations and discuss a possibility that neutrino observation gives some information on the nonthermal baryon-loading factor and the inner engine, if the internal shock model is true. Our numerical results are shown in Sec. \\ref{sec:level3}. Our summary and discussions are described in Sec. \\ref{sec:level4}. ",
        "conclusions": ""
    },
    "0512/astro-ph0512555_arXiv.txt": {
        "abstract": "We analyze several recently detected gamma-ray bursts (GRBs) with late X-ray flares in the context of late internal shock and late external shock models. We find that the X-ray flares in GRB 050421 and GRB 050502B originate from late internal shocks, while the main X-ray flares in GRB 050406 and GRB 050607 may arise from late external shocks. Under the assumption that the central engine has two periods of activities, we get four basic types of X-ray light curves. The classification of these types depends on which period of activities produces the prompt gamma-ray emission (Type 1 and Type 2: the earlier period; Type 3 and Type 4: the late period), and on whether the late ejecta catching up with the early ejecta happens earlier than the deceleration of the early ejecta (Type 1 and Type 3) or not (Type 2 and Type 4). We find that the X-ray flare caused by a late external shock is a special case of Type 1. Our analysis reveals that the X-ray light curves of GRBs 050406, 050421, and 050607 can be classified as Type 1, while the X-ray light curve of GRB 050502B is classified as Type 2. However, the X-ray light curve of GRB 050406 is also likely to be Type 2. We also predict a long-lag short-lived X-ray flare caused by the inner external shock, which forms when a low baryon-loading long-lag late ejecta decelerates in the non-relativistic tail of an outer external shock driven by an early ejecta. ",
        "introduction": "In the pre-{\\it Swift} era, only a few early optical afterglows of gamma-ray bursts (GRBs) were detected, while the observations of X-ray afterglows usually started at $\\sim10^{4}$ seconds after the prompt trigger. Since early afterglows contain important information about GRB central engines, understanding the early afterglows is therefore one of the most interesting scientific goals of the NASA's {\\it Swift} satellite (Gehrels et al. 2004). After about one year of operations, a significant fraction of well localized GRBs have not been detected with early optical emissions down to moderate limiting magnitudes, although the UV/Optical Telescope (UVOT) on board {\\it Swift} slewed to the error circle of position quickly after the burst (Roming et al. 2005). On the other hand, observations by the X-ray Telescope (XRT) on board {\\it Swift} have revealed several new features of X-ray emissions. First, steep declines in some X-ray light curves at the transition from prompt phase to afterglow phase have been discovered, which are interpreted as tail emissions of prompt GRBs (Tagliaferri et al. 2005; Zhang et al. 2005; Nousek et al. 2005). Secondly, a portion of early X-ray afterglows have shallow-than-normal temporal decays before they enter ``normal'' decaying phase (Nousek et al. 2005; Zhang et al. 2005). Thirdly, late X-ray flares have been observed in several GRBs (Burrows et al. 2005a; see also Burrows et al. 2005b for a review). Piro et al. (2005) also reported that X-ray flares were discovered in a few GRBs (e.g., GRBs 011121 and 011211) by the Italian-Dutch {\\it Beppo}SAX satellite. GRBs are usually divided into two main classes by their durations and spectral hardness ratios, i.e., long GRBs have durations larger than 2 seconds and softer spectra while short GRBs have durations shorter than 2 seconds and harder spectra (Kouveliotou et al. 1993). Up to now, both long GRBs, including X-ray rich bursts and a high redshift burst (GRB 050904), and short GRBs have been found with late X-ray flares (Galli \\& Piro 2005; Watson et al. 2005; Fox et al. 2005; Barthelmy et al. 2005).  One leading explanation for the early shallow decaying X-ray afterglows is that the forward external shock has continuous energy injection from the long active central engine (Zhang et al. 2005; Nousek et al. 2005; Dai \\& Lu 1998a; Zhang \\& \\meszaros 2002; Dai 2004; Wei, Yan \\& Fan 2005). Another leading explanation is that the Lorentz factor of the GRB ejecta has a distribution shaped by the central engine so that the behind slower material catches up with the ahead faster material when the latter is decelerated in the circum-burst medium, acting as a continuous energy injection (Zhang et al. 2005; Nousek et al. 2005; Rees \\& \\meszaros 1998; Sari \\& \\meszaros 2000; Zhang \\& \\meszaros 2002; Granot \\& Kumar 2005). In this explanation the central engine does not need to be active for a long time (Granot \\& Kumar 2005). It is also possible that the early shallow decaying X-ray afterglows are caused by late continuous energy releases if the GRB ejecta is initially dominated by Poynting flux (Zhang \\& Kobayashi 2004).  As for X-ray flares, both late external shock model and late internal shock model have been proposed (Piro et al. 2005; Burrows et al. 2005; Fan \\& Wei 2005; Zhang et al. 2005). Piro et al. (2005) and Galli \\& Piro (2005) argued for the late external shock model because they found there is no obvious spectral evolution in some X-ray flares, while Burrows et al. (2005b) preferred the late internal shock model based on several other X-ray flares which showed opposite features. These two conclusions are only based on their qualitative and empirical analysis. In this paper we quantitatively compare the X-ray flares with the predictions of the above two models. In Section 2 we describe the intrinsic X-ray light curves from the internal and external shocks. Then we analyze four GRBs with X-ray flares in Section 3 by comparing with the ``delayed'' (relative to the GRB trigger) intrinsic light curves of the late internal shock and late external shock. In Section 4 we establish four basic types of X-ray light curves, assuming that the central engine has two periods of activities. We also classify the four GRBs by these types. Our discussion and conclusions about X-ray flares are presented in Section 5. ",
        "conclusions": "In this paper we have quantitatively analyzed late X-ray flares in the frameworks of the late internal shock model and late external shock model. As we have shown above, both the late internal shock model and late external shock model require late time activities of central engines. Some of previous works suggested the late internal shock origin for X-ray flares (Burrows et al. 2005a; Zhang et al. 2005; Fan \\& Wei 2005), while others suggested the late external shock origin (Piro et al. 2005; Galli \\& Piro 2005). In fact, these two kinds of late shocks can coexist within a certain gamma-ray burst. GRB 050607 may be such a possible candidate.  Here one caveat should be made for the late external shock model. An X-ray flare originating from the late external shock comes to being only if the prompt gamma-ray burst is produced by internal shocks within the early time ejecta, and the late time ejecta carries more energy than the early one and its initial Lorentz factor $\\Gamma_{l0}$ must be larger than the initial Lorentz factor of the early ejecta $\\Gamma_{e0}$ by a factor of $(t_{{\\rm{dec}},e}/t_{\\rm{lag}})^{1/2}>1$, where $t_{\\rm{lag}}$ is the time separation between the two ejecta by the central engine and $t_{{\\rm{dec}},e}$ is the deceleration time of the early ejecta in the circum-burst medium. The merger of the two ejecta is mainly dominated by the late ejecta with a much shorter deceleration time relative to $t_{{\\rm{dec}},e}$. Therefore the beginning of the afterglow from the merger happens much later than the prompt GRB and the peak of afterglow behaves as an X-ray flare. Note this description of the late internal shock model is different from that in Piro et al. (2005) and Galli \\& Piro (2005). They attributed the late external shock to arise from a thick shell which is ejected by the long active central engine. Despite different descriptions, this model requires the $t_{0}$ effect in the afterglow emission. Given $t_{0}\\sim t_{\\rm{lag}}$, the smaller the deceleration time of the merger $t_{{\\rm{dec}},m}$, the more remarkable the afterglow as an X-ray flare. Since $t_{{\\rm{dec}},m}\\propto E_{m}^{1/(3-k)}\\Gamma_{m0}^{-2(4-k)/(3-k)}$, a large energy or especially a large Lorentz factor will easily cause an X-ray flare. This result has also been obtained by Galli \\& Piro (2005).  Theoretically, there are four basic types of X-ray light curves if the central engine has two periods of activities before it entirely ceases. According to our analysis, the observed X-ray light curves of GRBs with X-ray flares all belong to {\\it Type 1} or {\\it Type 2}. This means the prompt gamma-ray emissions of these GRBs result from the early ejecta, while the late X-ray flares are produced by the internal shock within the late ejecta, by the internal shock between the early and late ejecta, or by the late external shock. The difference between {\\it Type 1} and {\\it Type 2} is that for the former the collision between the early and late ejecta happens before the deceleration of the early ejecta, while for the latter the situation is inverse. One may wonder where are {\\it Type 3} and {\\it Type 4} X-ray light curves? In these two types, the prompt gamma-ray burst is produced from the late ejecta. There would be no internal shocks within the early ejecta, or these internal shock emission is too weak to be detected. It is also possible that the internal shock emission within the early ejecta is very weak and regarded as the precursor of the main burst. According to our classification, {\\it Type 3} X-ray light curves correspond to the case that the collision between the early and late ejecta takes place before the deceleration of the early ejecta. The emission from the collision has a small time lag relative to the prompt gamma-ray emission and can be regarded as the last pulse of the prompt GRB. The afterglow emission in {\\it Type 3} hardly suffers the $t_{0}$ effect. Therefore it is hard to distinguish {\\it Type 3} X-ray light curves from the light curves in which the central engine has only one active period. However, {\\it Type 4} X-ray light curves may have already been detected. Some of recently discovered X-ray afterglows whose initial decay is extraordinarily slow may originate from this type.  In this paper we also investigate a new kind of collision between the early and late ejecta, the {\\it inner external shock}. Such a shock will be developed when the late ejecta is decelerated by the non-relativistic tail of the outer external shock, which is driven by the early ejecta. To develop the inner external shock requires the time lag $t_{\\rm{lag}}$ between these two ejecta is very long, typically $\\geq1$ day, and the baryon loading of the late ejecta is very low, or equivalently, its initial Lorentz factor must be larger than a critical value. Since the required $t_{\\rm{lag}}$ is not very extreme compared to recent discoveries and the low baryon loading is also plausible, the inner external shock maybe exist. Because the time scale of emission from this shock in the observer's frame is equal to its deceleration time and therefore is much shorter than $t_{\\rm{lag}}$, the emission behaves as a short spike. Detecting such a spike in the typical afterglow time scale ($t\\sim t_{\\rm{lag}}\\sim1$ day) is quite difficult. We therefore just make a prediction of the inner external shock emission in this paper, and have not considered such kind of emission in the above four basic X-ray light curves.  For simplicity we have only considered the central engine having two periods of activities. In reality the central engine may have more than two periods of activities. The multiple X-ray flares detected in the high redshift gamma ray burst GRB 050904 indicate that the central engine of this burst has been active for many periods (Watson et al. 2005; Cusumano et al. 2005). Cusumano et al. (2005) reported that the hardness ratio for $t>100$ s is nearly a constant, $H(14-73{\\rm{ keV}})/S(1.4-14{\\rm{ keV}})\\sim0.2$, which corresponds to $\\beta\\sim-0.9$. The X-ray light curves from multiple ejections by the central engine are complicated and difficult to be classified. As for GRB 050904, Zou, Xu \\& Dai (2005) found that ordered late internal shocks formed by collisions between late ejecta and the earliest ejecta can reproduce the temporal evolution of peak luminosity of X-ray flares and the entire X-ray spectral evolution.  At last we would like to point out other ways besides those developed in this paper (see Section 3) to distinguish an X-ray flare caused by a late external shock from other X-ray flares in the same GRB by late internal shock(s). As we know in most pulses of GRBs it is found that higher energy photons arrive to the observer earlier than lower energy photons (e.g., Chen et al. 2005). This is the spectral-lag phenomenon. Fenimore et al. (1995) found that the width $W$ of peaks in prompt GRBs tends to be narrower at higher photon energy $E$, i.e., $W\\propto E^{-0.4}$. This is called the spectral-width relationship. If these two phenomena have been detected in a particular X-ray flare, then the X-ray flare can be proved to originate from a late internal shock. Statistics help us to understand X-ray flares. In prompt GRBs, several relationships have been found, such as the isotropic gamma-ray released energy $E_{\\rm{iso}}$ and spectral peak photon energy $E_{p}$ relationship (Amati et al. 2002), the peak luminosity $L_{p}$ and spectral peak photon energy $E_{p}$ relationship (Yonetoku et al. 2004), the peak luminosity $L_{p}$ and spectral time lag $\\tau$ relationship (Norris, Marani \\& Bonnell 2000). Moreover, Liang, Dai \\& Wu (2004) found there exists a relationship between the flux and $E_{p}$ within a GRB with time resolved spectra (see also Borgonovo \\& Ryde 2001). Therefore, investigating the above relations in X-ray flares is urgent since it can provide direct evidence that whether late X-ray flares and prompt GRBs have the same mechanism.  {"
    },
    "0512/astro-ph0512249_arXiv.txt": {
        "abstract": "We present the results of two {\\it XMM-Newton} observations of Jupiter carried out in 2003 for 100 and 250 ks (or 3 and 7 planet rotations) respectively. X-ray images from the EPIC CCD cameras show prominent emission from the auroral regions in the 0.2$-$2.0 keV band: the spectra are well modelled by a combination of emission lines, including most prominently those of highly ionised oxygen (OVII and OVIII). In addition, and for the first time, {\\it XMM-Newton} reveals the presence in both aurorae of a higher energy component (3$-$7 keV) which is well described by an electron bremsstrahlung spectrum. This component is found to be variable in flux and spectral shape during the Nov. 2003 observation, which corresponded to an extended period of intense solar activity. Emission from the equatorial regions of the Jupiter's disk is also observed, with a spectrum consistent with that of solar X-rays scattered in the planet's upper atmosphere. Jupiter's X-rays are spectrally resolved with the RGS which clearly separates the prominent OVII contribution of the aurorae from the OVIII, FeXVII and MgXI lines, originating in the low-latitude disk regions of the planet. ",
        "introduction": "Jupiter was first detected as an X-ray source with the {\\it Einstein} observatory \\citep{Met:83}. By analogy with the Earth's aurorae, the emission was expected to be produced via bremsstrahlung by energetic electrons precipitating from the magnetosphere. However, the observed spectrum is softer (0.2$-$3 keV) and the observed fluxes larger than predicted from this mechanism. Model calculations by \\citet{Sing:92} confirmed that the expected bremsstrahlung flux is lower by 1 to 2 orders of magnitude compared with the observed $<$2 keV X-ray flux. The alternative process is K shell line emission from ions, mostly of oxygen, which charge exchange, are left in an excited state and then decay back to the ground state (see Bhardwaj and Gladstone, 2000, for a review of early work on planetary auroral emissions). The ions were thought to originate in Jupiter's inner magnetosphere, where an abundance of sulphur and oxygen, associated with Io and its plasma torus, is expected \\citep{Met:83}.  The first {\\it ROSAT} soft X-ray (0.1$-$2.0 keV) observations produced a spectrum much more consistent with recombination line emission than with bremsstrahlung (Waite et al., 1994; Cravens et al., 1995). Subsequent {\\it ROSAT} observations also revealed low-latitude `disk' emission from Jupiter \\citep{Waite:97}, and this too was attributed to charge exchange. However, the X-rays were brightest at the planet's limb corresponding to the position of the subsolar point relative to the sub-Earth point, suggesting that a solar-driven mechanism may be at work \\citep{gla:98}. Scattering of solar X-rays, both elastic (by atmospheric neutrals) and fluorescent (of carbon K-shell X-rays off methane molecules below the Jovian homopause), was put forward as a way to explain the disk emission \\citep{mau:00}.  With the advent of the {\\it Chandra} observatory we gained the clearest view yet of Jupiter's X-ray emission, but more questions arose as well: HRC-I observations in Dec. 2000 and Feb. 2003 clearly resolve two bright, high-latitude sources associated with the aurorae, as well as low-latitude emission from the planet's disk (Gladstone et al., 2002; Elsner et al., 2005). However, the Northern X-ray hot spot is found to be magnetically mapped to distances in excess of 30 Jovian radii, rather than to the inner magnetosphere and the Io plasma torus. Since in the outer magnetosphere ion fluxes are insufficient to explain the observed X-ray emission, another ion source (solar wind?) and/or acceleration mechanism are required. Strong 45 min quasi-periodic X-ray oscillations were also discovered using {\\it Chandra} data in the North auroral spot in Dec. 2000, without any correlated periodicity being observed in {\\it Cassini} upstream solar wind data, or in {\\it Galileo} and {\\it Cassini} energetic particle and plasma wave measurements \\citep{gla:02}. {\\it Chandra} ACIS-S observations \\citep{els:05} show that the auroral X-ray spectrum is made up of oxygen line emission consistent with mostly fully stripped ions. Line emission at lower energies could be from sulphur and/or carbon. The high charge states and the observed fluxes imply that the ions must have undergone acceleration, independently from their origin, magnetospheric or solar wind. Rather than periodic oscillations, chaotic variability of the auroral X-ray emission was observed, with power peaks in the 20$-$70 min range. A promising mechanism which could explain this change in character of the variability, from organised to chaotic, is pulsed reconnection at the day-side magnetopause, as suggested by \\citet{bun:04}. ",
        "conclusions": "{\\it XMM-Newton} observations of Jupiter on two epochs in Apr. and Nov. 2003 convincingly demonstrate that auroral and low-latitude disk X-ray emissions are different in spectral shape and origin. The Jovian auroral soft X-rays ($<$ 2 keV) are most likely due to charge exchange by energetic ions from the outer magnetosphere, or solar wind, or both. For the first time a higher energy component in the auroral spectra has been identified, and has been found to be variable over timescales of days: its spectral shape is consistent with that predicted from bremsstrahlung of energetic electrons precipitating from the magnetosphere. The variability observed in its flux and spectrum is likely to be linked to changes in the energy distribution of the electrons producing it and may be related to the particular period of intense solar activity reported in Oct. - Nov. 2003 by a number of spacecraft mesurements."
    },
    "0512/astro-ph0512280_arXiv.txt": {
        "abstract": "We present an analysis of fine-structure transitions of \\ion{Fe}{2} and \\ion{Si}{2} detected in a high-resolution optical spectrum of the afterglow of \\grb\\ ($z=1.54948$).  The fine-structure absorption features arising from \\ion{Fe}{2}* to \\ion{Fe}{2}****, as well as \\ion{Si}{2}*, are confined to a narrow velocity structure extending over $\\pm 30$ km s$^{-1}$, which we interpret as the burst local environment, most likely a star forming region.  We investigate two scenarios for the excitation of the fine-structure levels by collisions with electrons and radiative pumping by an infra-red or ultra-violet radiation field produced by intense star formation in the GRB environment, or by the GRB afterglow itself.  We find that the conditions required for collisional excitation of \\ion{Fe}{2} fine-structure states cannot be easily reconciled with the relatively weak \\ion{Si}{2}* absorption.  Radiative pumping by either IR or UV emission requires $>10^3$ massive hot OB stars within a compact star-forming region a few pc in size, and in the case of IR pumping a large dust content.  On the other hand, it is possible that the GRB itself provides the source of IR and/or UV radiation, in which case we estimate that the excitation takes place at a distance of $\\sim 10-20$ pc from the burst.  Detailed radiative transfer calculations are required in order to verify that excitation of the low-ionization fine-structure states is possible given the intense UV flux from the burst.  Still, it is clear that GRB absorption spectroscopy can provide direct information on the mode and conditions of star formation at high redshift. ",
        "introduction": "\\label{sec:intro}  Understanding the physical conditions in star forming regions at high redshift plays a crucial role in our ability to trace the evolution of star formation and the associated production of metals. Unfortunately, it is exceedingly difficult to probe these regions using background quasars because of their compact size, and possible dust extinction.  Absorption studies of bright star forming galaxies (e.g., the lensed Lyman break galaxy MS 1512$-$cB58; \\citealt{prs+02}) provide greater promise because they probe the regions where star formation is taking place, but even in these rare cases information is only available on the {\\it integrated} properties and individual regions cannot be probed.  Gamma-ray bursts (GRBs), on the other hand, are known to be the end product of massive stars, and as a result their bright afterglows can be used to trace the interstellar medium of their host galaxies.  More importantly, since GRB progenitors are short-lived, they are likely to be buried within star-forming regions undergoing active star formation \\citep{bkd02}.  As a result, GRBs may provide an ideal direct probe of the physical conditions in high redshift star forming regions.  Over the last several years, and particularly with the advent of the {\\it Swift} satellite, several GRB absorption spectra have been obtained revealing a large fraction of objects with unusually large neutral hydrogen column densities even for damped Ly$\\alpha$ (DLA) systems, $N\\gtrsim 10^{22}$ cm$^{-2}$ (e.g., \\citealt{vel+04,bpc+05}).  In addition, it has been noted in several cases that the column densities of non-refractory elements (e.g., Zn) are higher than in QSO-DLAs, with a depletion pattern suggestive of a large dust content \\citep{sf04}.  These properties suggest that at least some GRB-DLAs are physically different from QSO-DLAs, and that they may represent individual star-forming regions with at least a modest dust obscuration.  Given the possible association with star-forming regions, several authors have also investigated the impact of the GRB prompt and afterglow emission on the local environment, particularly in the context of dust destruction, photo-ionization, and dissociation and excitation of $H_2$ molecules \\citep{dh02}.  The basic conclusion of these studies is that bright afterglows are capable of significantly modifying their local environment on a scale of about $10^{19}$ cm. The detection of such effects can therefore provide a unique diagnostic of the immediate environment of the burst, and in addition may address the issue of relatively low dust extinction in GRB optical afterglows.  In this {\\it Letter} we present a high-resolution absorption spectrum of \\grb, which reveals strong absorption features from \\ion{Fe}{2} and \\ion{Si}{2} fine-structure states at a redshift $z=1.54948$.  This is the first example of an extra-galactic sight line that reveals the full range of \\ion{Fe}{2} ground-level fine-structure transitions, suggesting that their excitation is intimately related to the burst environment or the burst itself.  We investigate both possibilities and their implications in the following sections. ",
        "conclusions": "\\label{sec:conc}  The high-resolution spectrum of \\grb\\ reveals a wide range of absorption features from fine-structure states of \\ion{Fe}{2} and \\ion{Si}{2} at $z=1.54948$.  We associate the absorption with the local environment of the burst based on the fact that the excitation of these fine-structure states requires large densities or intense IR/UV radiation fields, which are not common in typical interstellar environments.  This conclusion is also supported by the kinematic coincidence with the burst systemic redshift.  A comparison of the conditions required to excite the \\ion{Fe}{2} and \\ion{Si}{2} fine-structure levels indicates that collisional excitation is not likely to be the dominant mechanism, unless the two ions are segregated in different regions of the absorber.  A more likely mechanuism is direct IR pumping or indirect UV pumping.  In this scenario we find that if the radiation field is due to star formation activity in the burst environment, then this requires a large concentration of several thousand OB stars in a compact region ($\\sim {\\rm few}$ pc), perhaps reminiscent of the super star cluster in the dwarf galaxy NGC\\,5253.  If the UV and/or IR radiation are supplied by the GRB, on the other hand, then the distance to the absorber is likely to be about $10-20$ pc.  This is consistent with calculations by \\citet{dh02}, which indicate that within $\\sim 3$ pc, the burst will ionize its environment, leading to a very low column of low-ionization states such as \\ion{Fe}{2} and \\ion{Si}{2}; outside of this region it is likely that \\ion{Fe}{2} and \\ion{Si}{2} survive and may be excited to fine-structure levels instead.  Clearly, detailed radiative transfer calculations, along with refined calculations of \\ion{Fe}{2} fine-structure excitation, are required in order to distinguish between the different scenarios.  Still, regardless of the exact details it is clear that GRB afterglows allow us to directly probe the conditions within individual star forming regions, and may therefore provide direct information on the mode of star formation across a wide redshift range.  The absorption spectrum of \\grb\\ seems to indicate that compact, dusty, and dense star forming regions may be prevalent, at least in the context of GRB progenitor formation."
    },
    "0512/hep-th0512118_arXiv.txt": {
        "abstract": "Dark energy dynamics of the universe can be achieved by equivalent mathematical descriptions taking into account generalized fluid equations of state in General Relativity, scalar-tensor theories or modified $F(R)$ gravity in Einstein or Jordan frames. The corresponding technique transforming equation of state description to scalar-tensor or modified gravity is explicitly presented. We show that such equivalent pictures can be discriminated  by matching solutions with data capable of selecting the true physical frame. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/hep-th0512268_arXiv.txt": {
        "abstract": "The most entropic fluid can be related to a \\emph{dense} gas of black holes that we use to study the beginning of the universe. We encounter difficulties to compatibilize an adiabatic expansion with the growing area for the coalescence of black holes. This problem may be circumvented for a quantum black hole fluid, whose classical counterpart can be described by a percolating process at the critical point. This classical regime might be related to the energy content of the current universe. ",
        "introduction": "To understand the initial conditions of the universe we need a description of the gravitational field when the curvature is the same order or smaller than the Planck length. The most important candidate for a consistent description of quantum gravity is the string theory where the fundamental particles are no longer point-like but have a linear extension. There is a small number of consistent ways to quantize strings, all related by dualities, signaling an underlying structure that has been termed M theory. Despite this strong restriction in the consistent way to quantize the gravitational field, the true vacuum in which the theory lives is not predicted. Instead, a great amount of vacua are possible all compatible with the premises of the theory. The situation is such, that arguments outside the theory, mainly anthropic, have been used to select a vacuum \\cite{weinberg}. Our purpose is to address the initial conditions using entropic arguments instead.  One of the basis underlying the theory of quantum gravity is the holographic principle\\cite{thooft}. This states that a quantum description of the gravitational phenomena can be mapped into a dual theory on the boundary. In the string scenario such dual description is exemplified by the  celebrated AdS/CFT correspondence \\cite{maldacena}. If a theory can be mapped onto its boundary, the operative number of degrees of freedom grows with the area of such boundary, in contrast to the case of ordinary quantum fields which scales with the volume of the system. A way to understand this drastic reduction on the number of degrees of freedom is by the formation of black holes (BH's). Due to the fact that the information carriers have energy, when we try to accumulate an amount of information in a smaller volume the subsystem can collapse forming a BH and lose the corresponding information. An amount of entropy is generated in this process.  In the semiclassical regime, when the curvature is larger than the Planck length, the holographic principle translates into different types of entropy bounds all sharing a property: the limitation on the number of degrees of freedom by the area which circumventes the system; an improvement of the situation was given by Bousso's covariant entropy bound \\cite{bousso} .  A BH is the most entropic system for a given amount of energy \\cite{thooft}; if information is not lost in the subsequent evaporation, the BH is the most efficient recipient of information. Such information would be encoded in its event horizon and ejected away in subtle, higher non local correlation of its atmosphere. A gas of BH's will be the more entropic fluid and such universe will be able to attain the maximum complexity \\cite{banks}. ",
        "conclusions": ""
    },
    "0512/astro-ph0512005_arXiv.txt": {
        "abstract": "{The high mass X-ray binary pulsar 4U 1538-52 was observed between July 31 and August 7, 2003. Using these observations, we determined new orbital epochs for both circular and elliptical orbit models. The orbital epochs for both orbit solutions agreed with each other and yielded an orbital period derivative $\\dot{P} / P = (0.4 \\pm 1.8 ) \\times 10^{-6}$ yr$^{-1}$. This value is consistent with the earlier measurement of $\\dot{P} / P = (2.9 \\pm 2.1 ) \\times 10^{-6}$ yr$^{-1}$ at the $1 \\sigma$ level and gives only an upper limit to the orbital period decay. Our determination of the pulse frequency  showed that the source spun up at an average rate of $2.76 \\times 10^{-14}$ Hz sec$^{-1}$ between 1991 and 2003. ",
        "introduction": " ",
        "conclusions": "\\begin{figure}[h] \\begin{center} \\psfig{file=pgplot.ps,height=10cm,width=12cm,angle=-90} \\small{Fig. 4 -- Pulse frequency history of 4U 1538-52. The rightmost point corresponds to most recent RXTE observation of ID 80016.} \\end{center} \\end{figure}  Before CGRO observations, 4U 1538-52 had been found to have a long-term spin down trend. A linear fit to pre-CGRO pulse frequency history gives $\\dot{\\nu} /\\nu \\sim -8 \\times 10^{-12} s^{-1}$ and a linear fit to CGRO and our RXTE result yields $\\dot {\\nu} /\\nu \\sim 1.45 \\times 10^{-11} s^{-1}$. Rubin et al (1997) constructed the power spectrum of pulse frequency derivative fluctuations. Their analysis showed that the pulse frequency derivative fluctuations can be explained on timescales from 16 to 1600 days with an average white noise strength of (7.6 $\\pm$ 1.6)$\\times 10^{-21}$ (Hz s$^{-1}$)$^{2}$ Hz$^{-1}$.  A random walk in pulse frequency (or white noise in pulse frequency derivative) can be explained as a sequence of {\\bf{steps in pulse frequency with an RMS value of $<(\\delta \\nu ^{2})>$}} which occur at a constant rate R. Then the RMS variation of the pulse frequency scales with elapsed time $\\tau$ as $<(\\Delta \\nu) ^{2}> = R <(\\delta \\nu) ^{2}> \\tau$ (Hz), where $S= R <\\delta \\nu ^{2}>$ is defined as noise strength. Then, RMS scaling for the pulse frequency derivatives can be obtained as $<(\\Delta \\dot \\nu) ^{2}>^{1/2} = (S/\\tau )^{1/2} {\\rm{Hz.s}}^{-1}$. As seen from Table 1, in our fits, upper limits on intrinsic pulse frequency derivatives are 7-10 times higher than the long-term spin up rates. If white noise in the pulse frequency derivative can be interpolated to a few days, {\\bf{then the upper limit on the change of}} frequency derivative obtained from a $\\sim 1$ week observation should typically have a magnitude that can be estimated from $<(\\Delta \\dot \\nu) ^{2}_{week}>= <(\\Delta \\dot \\nu) ^{2}_{1600 days} \\times 15^2 $ {\\bf{This value is 15/7 - 15/10 times higher than the measured upper limit values. Therefore the measured upper limits on the intrinsic pulse frequency derivatives for 1 week are consistent with the values from the extrapolation of the power spectrum within a factor of a few.}}  Previous marginal measurement of change in the orbital period, was (-2.9 $\\pm$ 2.1)$\\times 10^{-6}$ yr$^{-1}$ (Clark 2000), and our new value for the orbital period change, $\\dot{P}/P = (0.4 \\pm 1.8 ) \\times 10^{-6}$yr$^{-1}$, are consistent with zero. These two measurements are consistent with each other in $1 \\sigma $ level.   In most of the X-ray binaries with accretion powered pulsars, the evolution of the orbital period seems to be too slow to be detectable. Yet there are still some such systems in which this evolution was measured and $\\dot{P}/P$ were reported. These systems include Cen X-3 with (-1.8 $\\pm$ 0.1)$\\times$10$^{-6}$ yr$^{-1}$ (Kelley et al. 1983; Nagase et al. 1992), Her X-1 with (-1.32 $\\pm$ 0.16)$\\times$10$^{-8}$ yr$^{-1}$ (Deeter et al. 1991), SMC X-1 with (-3.36 $\\pm$ 0.02) $\\times$10$^{-6}$ yr$^{-1}$ (Levine et al. 1993), Cyg X-3 with (1.17 $\\pm$ 0.44) $\\times$10$^{-6}$ yr$^{-1}$ (Kitamoto et al. 1995), 4U 1700-37 with (3.3 $\\pm$ 0.6) $\\times$10$^{-6}$ yr$^{-1}$ (Rubin et al. 1996), and LMC X-4 with (-9.8 $\\pm$ 0.7)$\\times$10$^{-7}$ yr$^{-1}$ (Levine et al. 2000). Change in the orbital period of Cyg X-3 was associated with the mass loss rate from the Wolf-Rayet companion star. For 4U 1700-37, the major cause of orbital period change was also thought to be mass loss from the companion star. For Her X-1, mass loss and mass transfer from the companion were proposed to be the reasons of the change in the orbital period of the system.  On the other hand, for the high mass X-ray binary systems Cen X-3, LMC X-4 and SMC X-1, the major cause of change in the orbital period is likely to be tidal interactions (Kelley et al. 1983; Levine et al. 2000; Levine et al. 1993). For these three systems, orbital period decreases (i.e. derivative of the orbital period is negative). Our new measurement of orbital period change ($\\dot{P}/P$) gives the value of about $-10^{-6}$ yr$^{-1}$ which is similar to the observed values of SMC X-1 and Cen X-3. Further observations can give further information about the orbital period change of this source.  {\\bf"
    },
    "0512/astro-ph0512233_arXiv.txt": {
        "abstract": "We measure the clustering of DEEP2 galaxies at $z=1$ as a function of luminosity on scales $0.1$ \\mpch \\ to $20$ \\mpch.  Drawing from a parent catalog of 25,000 galaxies at $0.7<z<1.3$ in the full DEEP2 survey, we create volume-limited samples having upper luminosity limits between $M_B=-19$ and $M_B=-20.5$, roughly $0.2-1 L^*$ at $z=1$.  We find that brighter galaxies are more strongly clustered than fainter galaxies and that the slope of the correlation function does not depend on luminosity for $L<L^*$.  The brightest galaxies, with $L>L^*$, have a steeper slope.  The clustering scale-length, \\rr, varies from $3.69 \\pm0.14$ for the faintest sample to $4.43 \\pm0.14$ for the brightest sample.  The relative bias of galaxies as a function of $L/L^*$ is steeper than the relation found locally for SDSS galaxies \\citep{Zehavi05} over the luminosity range that we sample. The absolute bias of galaxies at $z\\sim1$ is scale-dependent on scales $r_p<1$ \\mpch, and rises most significantly on small scales for the brightest samples. For a concordance cosmology, the large-scale bias varies from $1.26 \\pm0.04$ to $1.54 \\pm0.05$ as a function of luminosity and implies that DEEP2 galaxies reside in dark matter halos with a minimum mass of $\\sim1-3 \\times10^{12} h^{-1} {\\rm M}_\\sun$. ",
        "introduction": "The clustering of galaxies has long been used as a fundamental measure of the large-scale structure of the universe.  Clustering measures place strong constraints on galaxy formation and evolution models and provide estimates of the average or minimum parent dark matter halo mass of a given galaxy population, allowing placement in a cosmological context \\citep[e.g.,][]{Mo96, Sheth99}.  Locally, large surveys such as the 2-Degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS) have measured the two-point correlation function of galaxies precisely enough to allow detailed halo occupation distribution (HOD) modeling \\citep[e.g.,][]{Yan03b, Phleps05,Yang05HOD,Zehavi05}, as well as cosmological parameter constraints \\citep[e.g.,][]{Peacock01,Abazajian05}.  In this paper we focus on the observed clustering of galaxies as a function of luminosity. Both 2dF and SDSS have found that at $z\\sim0.1$ brighter galaxies cluster more strongly, with the relative bias on large scales increasing linearly with luminosity \\citep[e.g.,][]{Alimi88,Hamilton88,Benoist96, Willmer98,Norberg01,Norberg02,Zehavi04}.  A simple interpretation is that brighter galaxies reside in more massive dark matter halos, which are more clustered in the standard theories of hierarchical structure formation.  The full galaxy samples for these surveys are generally unbiased with respect to the dark matter density field on large scales \\citep[e.g.,][]{Peacock01,Verde01,Croton04}.  Locally, the two-point correlation function is relatively well-fit by a power law, \\xir=$(r/r_0)^{-\\gamma}$, with $r_0\\sim5$ \\mpch \\ and $\\gamma\\sim1.8$ \\citep[e.g.,][]{Norberg02,Zehavi05}.  Small deviations from a power law have recently been detected on scales $r<10$ \\mpch \\ and can be naturally be explained in the HOD framework as the transition between pairs of galaxies within the same halo (the `one-halo' term) on small scales and galaxies in different halos (the `two-halo' term) on larger scales \\citep{Zehavi04}.  The largest deviations are seen for the brightest galaxy samples, with $L>L^*$. These galaxies are generally found in groups and clusters, which have larger `one-halo' terms due to the many galaxies residing in a single parent dark matter halo and therefore have a steeper correlation slope on small scales.  Lyman-break galaxies (LBGs) at $z\\sim3-4$ also display a luminosity-dependence in their clustering properties \\citep{Allen05,Ouchi05,Lee05}.  The large-scale bias of LBGs ranges from $\\sim3-8$, depending on the luminosity of the sample. There is also a strong scale-dependence to the bias on small scales, where the correlation function rises sharply below $r\\sim0.5$ \\mpch \\citep{Ouchi05, Lee05}.  This rise has been interpreted as evidence for multiple LBGs residing in single dark matter halos at $z\\sim4$.  In this paper, we present results on the luminosity-dependent clustering of galaxies at $z\\sim1$, using the nearly-completed DEEP2 Galaxy Redshift Survey.  The data sample used here is 10 times larger than our initial results presented in \\cite{Coil03xisp} and represents the most robust measurements of \\xir \\ at $z>0.7$ to date.  In future papers we will address clustering properties as a function of other galaxy properties such as color, stellar mass and redshift; here we focus solely on the luminosity-dependence.  We choose to use integral, rather than differential, luminosity bins to facilitate comparison with theoretical models \\citep[e.g.,][]{Kravtsov04,Tinker05}. Luminosity-threshold samples are preferred for HOD modeling, as theoretical predictions for HODs have been studied more extensively for mass (luminosity) thresholds and afford fewer parameters to fit than differential bins \\citep[for details see e.g.][]{Zehavi05}. In this case the HOD has a relatively simple form, e.g., the mean occupation function is a step function (for central galaxies) plus a power law (for satellites).  The outline of the paper is as follows: \\S 2 briefly describes the DEEP2 survey and data sample.  In \\S 3 we outline the methods used in this paper, while \\S 4 presents our results on the luminosity-dependence of galaxy clustering at $z\\sim1$.  In \\S 5 we discuss the galaxy bias and the relative biases between our luminosity samples and conclude in \\S 6. ",
        "conclusions": "Using volume-limited subsamples of the 25,000 galaxies with spectroscopic redshifts between $0.7\\leq z \\leq1.4$ from the nearly-completed DEEP2 Galaxy Redshift Survey, we have measured the clustering of galaxies as a function of luminosity at $z\\sim1$.  We find that the clustering scale length, \\rr, increases linearly with luminosity over the range sampled here and that the bias of clustering relative to $L^*$ is strong to what is found at $z\\sim0$.  The brightest DEEP2 sample, with $L>L*$, has a significantly steeper correlation slope ($\\gamma=1.82 \\pm0.03$) on all scales than the fainter luminosity samples, which all have similar slopes of $\\gamma\\sim1.73$.  This is not due to differences in the mix of galaxies in each sample; the red galaxy fraction depends only weakly on luminosity in the DEEP2 data.  \\cite{Conroy05} attempt to model the clustering results presented here with a simple prescription in which luminosities are assigned to dark matter halos and sub-halos in an N-body simulation by matching the observed DEEP2 luminosity function to the sub-halo/halo maximum circular velocity, used as a proxy for mass.  Their model fits our data extremely well, implying that brighter galaxies reside in more massive dark matter halos and that the luminosity-dependent clustering is dominated by the physical distribution of halos on large scales and sub-halos on small scales, as a function of mass.  It is remarkable that the results presented in this paper can be explained using a model in which the luminosity of a galaxy depends solely on the mass of the dark matter halo in which it resides.  The mean galaxy bias in the DEEP2 data on scales $r_p=1-10$ \\mpch \\ for a concordance cosmology ranges from $1.26 \\pm0.04$ to $1.54 \\pm0.05$ as a function of luminosity for $L/L^*=0.7-1.3$, and there is a scale-dependence to the bias for $r_p<1$ \\mpch. The brightest samples show the strongest rise in the correlation function on small scales, which is likely due to these galaxies preferentially residing in groups. This upturn on small scales has also been seen both locally for SDSS galaxies with $L>L^*$ \\citep{Zehavi05} and at high redshift for LBG galaxies at $z\\sim3-4$ \\citep{Lee05,Ouchi05}.  The strength of the rise in the DEEP2 sample is smaller than what is found at $z\\sim3-4$, as expected from simulations \\citep{Kravtsov04}, and the large-scale bias is intermediate between the bias of LBGs and local galaxies of similar $L/L^*$ in 2dF and SDSS."
    },
    "0512/astro-ph0512143_arXiv.txt": {
        "abstract": "We have identified two new galaxies with gas counter-rotation (NGC~1596 and NGC~3203) and have confirmed similar behaviour in another one (NGC~128), this using results from separate studies of the ionized-gas and stellar kinematics of a well-defined sample of $30$ edge-on disc galaxies. Gas counter-rotators thus represent $10\\pm5\\%$ of our sample, but the fraction climbs to $21\\pm11\\%$ when only lenticular (S0) galaxies are considered and to $27\\pm13\\%$ for S0s with detected ionized-gas only. Those fractions are consistent with but slightly higher than previous studies. A compilation from well-defined studies of S0s in the literature yield fractions of $15\\pm4\\%$ and $23\\pm5\\%$, respectively. Although mainly based on circumstantial evidence, we argue that the counter-rotating gas originates primarily from minor mergers and tidally-induced transfer of material from nearby objects. Assuming isotropic accretion, twice those fractions of objects must have undergone similar processes, underlining the importance of (minor) accretion for galaxy evolution. Applications of gas counter-rotators to barred galaxy dynamics are also discussed. ",
        "introduction": "\\label{sec:intro} It has been already almost $20$ years since the phenomenon of counter-rotation in disc galaxies was discovered \\citep{g87}, but both the exact incidence and the origin of counter-rotating gas and stars remain to be clarified. Most statistical studies indicate that roughly $20$--$25\\%$ of all lenticular (S0) galaxies with detected ionized-gas (usually observed through optical emission lines) contain a non-negligible fraction of gas-stars counter-rotation \\citep*[e.g.][]{bbz92,kfm96,kf01,pcvb04}, typically in the central regions. The fraction is much lower in later type systems \\citep[e.g.][]{kf01,pcvb04}, presumably because co-rotating and counter-rotating gas can not coexist at a given radius. Galaxies with counter-rotating H$\\,${\\small I} or CO gas are also known (e.g.\\ NGC~4826, \\citealt*{bwk92}; NGC~3626, \\citealt*{gsb98}), but to our knowledge no sound statistics exists. The fraction of S0s with stars-stars counter-rotation is also much lower (at most $10\\%$ but perhaps lower; \\citealt{kfm96,pcvb04}), perhaps due to the increased difficulty of detecting small numbers of counter-rotating stars. Generally, however, the number of objects on which well-defined studies are based is still rather small, and it is still the case that most known counter-rotators were discovered fortuitously in targeted studies (e.g.\\ NGC~4546, \\citealt{g87}; NGC~4138, \\citealt*{jbh96}).  Our goal in this paper is thus to improve the (ionized-gas) counter-rotation statistics in S0 galaxies and summarize the current situation. Since counter-rotating material is a strong argument in favour of hierarchical galaxy formation scenarios, this is an important goal. In the interest of conciseness, we do not discuss here the issue of counter-rotation in elliptical galaxies, although it is probably related (see \\citealt{r94} and \\citealt{s98} for general reviews of galaxies with misaligned angular momenta).  In Section~\\ref{sec:obs}, we describe gas and stellar kinematics observations of a large sample of edge-on disc galaxies. The structure, kinematics, and environment of $2$ newly discovered ionized-gas counter-rotators and $1$ known object are discussed in detail in Section~\\ref{sec:galaxies}. In Section~\\ref{sec:discussion}, we quantify the incidence of counter-rotators using all available observations and discuss their possible origin through continuous gas infall and/or discrete gas and satellites accretion. Applications of gas counter-rotators to barred galaxy dynamics are also discussed. We conclude briefly in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Using our previous studies of the ionized-gas and stellar kinematics of a relatively large sample of $30$ edge-on disc galaxies with (mostly) boxy and peanut-shaped (B/PS) bulges, we have identified two new galaxies (NGC~1596 and NGC~3203) where the ionized gas is counter-rotating with respect to the bulk of the stars. We have also confirmed similar kinematics in one additional object (NGC~128). Counter-rotating gas is thus present in $10\\pm5\\%$ of our entire sample. However, if only lenticular (S0) galaxies are considered, the fraction climbs to $21\\pm11\\%$. This fraction further climbs to $27\\pm13\\%$ if only S0s with extended ionized-gas are considered. As discussed at length in the text, those fractions are consistent with but slightly higher than the few existing systematic studies available in the literature. Merging those studies, fractions of $15\\pm4\\%$ and $23\\pm5\\%$ are obtained for, respectively, all S0s and S0s with ionized-gas only.  Based on the presence of probable companions near the galaxies discussed here, we argued that minor mergers and tidally-induced transfer of material from nearby objects are primarily responsible for the counter-rotating gas. If accretion on to the objects is isotropic, similar processes must have been at work in roughly twice the fractions of objects discussed. This strongly argues for a non-negligible role of (minor) accretion in galaxy formation and evolution.  The presence of counter-rotating gas in barred galaxies offers a number of largely unexplored ways to probe the structure and dynamics of those objects. In edge-on galaxies in particular, where counter-rotation is easiest to detect, it offers an independent way of testing the barred nature of B/PS bulges."
    },
    "0512/astro-ph0512375_arXiv.txt": {
        "abstract": "We explore the growth of super-massive black holes and host galaxy bulges in the galaxy population using the Millennium Run $\\Lambda$CDM simulation coupled with a model of galaxy formation.  We find that, if galaxy mergers are the primary drivers for both bulge and black hole growth, then in the simplest picture one should expect the \\bb\\ relation to evolve with redshift, with a larger black hole mass associated with a given bulge mass at earlier times relative to the present day.  This result is independent of an evolving cold gas fraction in the galaxy population.  The evolution arises from the disruption of galactic disks during mergers that make a larger fractional mass contribution to bulges at low redshift than at earlier epochs.  There is no comparable growth mode for the black hole population.  Thus, this effect produces evolution in the \\bb\\ relation that is driven by bulge mass growth and not by black holes. ",
        "introduction": "\\label{intro}  Super-massive black hole masses are strongly correlated with their host bulge stellar mass, the so-called \\bb\\ relation \\citep{Magorrian1998, Marconi2003, Haring2004}.  This is at least true in the local universe, but also expected to extend out to higher redshifts.  This correlation suggests a common mechanism linking the growth of these two galactic components, with evidence proposing galaxy mergers as the most likely candidate. If true, and given that the global galaxy merger rate in a $\\Lambda$CDM universe evolves strongly with time, one may ask if we should expect to see the \\bb\\ relation also evolve.  Support for the idea that bulges and black holes grow through mergers arises primarily from the success of numerical simulations and galaxy formation models in reproducing many observed galaxy scaling relations.  Such works illustrate that much of the bulge mass of a galaxy can be accounted for by the disruption of disk stars from the merger progenitors, and merger triggered starbursts in the cold gas disk \\citep{Barnes1992, Mihos1994, Mihos1996, Cox2004}.  As a growth mechanism for black holes, merger induced perturbations of the gas close to the central massive object can drive gas inward, fueling what is observed to be a `quasar' period in a galaxy's history \\citep[see e.g.][]{Kauffmann2000, DiMatteo2005}.  In this simple picture the amount of cold gas present in a merging system plays a large part in how rapidly the black hole and bulge can grow.  If the growth dependence for each is a simple constant scaling with gas mass, as is commonly assumed in many models of galaxy formation, then their mass ratio will, on average, be approximately independent of any evolution in the global cold gas fraction.  This is because both bulges and black holes then co-evolve at a similar pace (drawing their new mass from the same gas reservoir).  Furthermore, during a merger bulges will add to bulges, and black holes may coalesce.  Thus, from this alone, one should expect little evolution in the \\bb\\ relation.  In this paper we explore an additional growth channel through which bulges gain mass that black holes do not have.  This is the disruption of merged satellite \\emph{disks}, and in the event of a major merger, the disruption of the central galaxy \\emph{disk}.  The stellar mass in such disks will have previously never contributed to the \\bb\\ relation.  If the bulge growth rate from such disrupted disks is not constant with time, then evolution in the \\bb\\ relation can occur.  We investigate this behavior using the Millennium Run $\\Lambda$CDM simulation \\citep{Springel2005} and a model for galaxy formation \\citep{Croton2006}.  This model follows the growth of galaxies (including their individual disk, bulge and black hole components) from high redshift to the present day, and provides a solid framework within which to undertake our analysis.  The results we find, however, are not be unique to our particular implementation of the galaxy formation physics but arise from the simple assumptions described above regarding black hole and bulge growth in galaxies.  Our aim in using this particular model is to illustrate what one may expect to see if these underlying growth mechanisms turn out to be true.  This paper is organized as follows.  In Section~\\ref{method} we briefly describe the Millennium Run $\\Lambda$CDM dark matter simulation and our model of galaxy formation, including the simple implementation of bulge and black hole growth.  In Section \\ref{results} we will use this model to investigate how black hole and bulges co-evolve together in the galaxy population from high redshift to the present.  We finish in Section~\\ref{discussion} with a discussion of the \\bb\\ relation in light of these results. ",
        "conclusions": "\\label{discussion}   \\begin{figure} \\plotone{./figures/figure2.ps} \\caption{The static model described in Section~\\ref{static}. (top) The evolution of the relative growth rates for black holes and bulges, $\\dot{m}_{\\rm BH}$/$\\dot{m}_{\\rm bulge}$, for all galaxies hosting black holes with bulge masses $m_{\\rm bulge}\\!>\\!10^9 M_{\\odot}$.  As discussed in the text, black holes grow from merger triggered cold gas accretion.  Bulges, on the other hand, grow from both merger induced starbursts and disrupted disks, so we plot these two growth channels independently (dashed and solid lines respectively).  Disrupted disks are found to contribute a larger faction to the bulge at low redshift relative to high.  (bottom) The accelerated contribution of disrupted disks drive an evolution in the \\bb\\ relation, shown here using the $\\chi^2$ model fits from Fig.~\\ref{fig1} across the entire redshift range.  The dotted bounding lines show the 1$\\sigma$ scatter of galaxies along the relation.} \\label{fig2} \\end{figure}  The rapid increase of bulge growth at late times in our static model is a consequence of two well understood effects.  The first is the steady rise in the star formation rate density of the universe from high redshift to approximately $z\\!=\\!1\\!-\\!2$.  If one accepts, as a general rule, the conventional wisdom that the bulk of this star formation occurs in stellar disks, then the outcome is a strongly increasing growth of disk mass across the galaxy population with time. As disks grow the second effect then becomes increasingly important. This effect stems from the hierarchical nature of a CDM universe, where mergers become more frequent as the universe ages, assembling structure from the bottom up.  As we discussed in Section~\\ref{model}, mergers also transform disks into bulges.  Thus, at late times, a larger fraction of the total stellar mass in the universe becomes locked up in the spheroid component of the population relative to earlier epochs.  This results in the accelerated bulge growth seen in Fig.~\\ref{fig2}, and which drives the evolution in the \\bb\\ relation shown in Fig.~\\ref{fig1}.  Observationally it is difficult to measure black hole and bulge masses. In the local universe \\cite{Magorrian1998} estimate $m_{\\rm BH}\\!\\sim 0.006\\,m_{\\rm bulge}$ from a sample 32 galaxies, while both \\cite{Marconi2003} and \\cite{Haring2004} independently find $m_{\\rm BH}\\!\\sim 0.002\\,m_{\\rm bulge}$ from improved measurements of $\\sim\\!30$ galaxies.  Although the statistics are still poor and the uncertainty large, locally at least all observations appear to be converging to a consistent result.  At higher redshifts the picture is much less clear.  For example, \\cite{Shields2003} claim little evolution in the relation can be inferred out to $z\\!\\sim 3$\\footnote{However see their most recent work \\citep{Shields2005}}, and \\cite{Adelberger2005} measure the quasar--galaxy cross-correlation function and find consistency with the local \\bb\\ ratio from a sample of $79$ $z\\!\\sim\\!2.5$ quasars.  On the other hand, \\cite{Treu2004} see variations in the $m_{\\rm BH}$--$\\sigma$ relation at $z\\!=\\!0.37$, while \\cite{Mclure2005} measure some evolution in the \\bb\\ relation using the 3CRR sample of radio galaxies.  Similarly, \\cite{Rix2001} use gravitational lensing to find that quasar host galaxies at $z\\!\\sim\\!2$ are much fainter than their low redshift counterparts containing quasars of similar luminosity, and \\cite{Walter2004} find a significant deviation from the local \\bb\\ relation for a $z\\!=\\!6.4$ quasar host galaxy.  Future observations will need to clarify the exact nature of both the high redshift black hole and host galaxy populations.  Recent theoretical work to understand the cosmological assembly of stars and super-massive black holes have led to interesting results. \\cite{Wyithe2003} present a model for super-massive black hole growth that successfully matches many local and high redshift AGN related observations. Their work results in a $m_{\\rm BH}$--$\\sigma$ relation constant with redshift while predicting the \\bb\\ relation evolves as $m_{\\rm BH}/m_{\\rm bulge} \\propto \\xi(z)^{1/2}(1+z)^{3/2} \\sim (1+z)^{1.15}$, where the final approximation is valid when $z\\!<\\!2$, and $\\xi(z)$ depends only on the cosmological parameters and is a weak function of redshift.  An evolution of this kind would be consistent with the scaling assumed in our dynamical model (Section~\\ref{dynamic}).   \\begin{figure} \\plotone{./figures/figure3.ps} \\caption{As for Fig.~\\ref{fig2}, however now showing the result for the dynamic model, where an evolution in the black hole feeding rate has been assumed (Section~\\ref{dynamic}). Evolution in the \\bb\\ relation is now much stronger than that seen previously.} \\label{fig3} \\end{figure}   Similarly, \\cite{Merloni2004b} constrain phenomenologically the joint evolution of super massive black holes and their host spheroids by fitting simultaneously the total stellar mass and star formation rate densities as a function of redshift, as well as the hard X-ray selected quasar luminosity function.  With the latter they assume that black holes grow exclusively through accretion.  Assuming a present day disc to spheroid ratio of $0.5$ \\citep{Tasca2005}, their work favors a model in which the \\bb\\ relation evolves as $\\sim\\!(1+z)^{1/2}$. This is a weaker effect than found by \\citeauthor{Wyithe2003}, however demonstrates both the range of evolution that may be expected, and most importantly, that such non-zero evolution can arise naturally from simple studies of black hole and bulge growth.  As discussed in Section~\\ref{intro}, in a $\\Lambda$CDM universe the effect described in this paper will be present in any model of black hole and bulge assembly driven by mergers.  Indeed, this has already been seen in the semi-analytic model of \\cite{Cattaneo2005} who find similar \\bb\\ evolution to that found here (compare their Fig.~6 with our Fig.~\\ref{fig1}).  Unfortunately they do not discuss the origin of this behavior, but instead choose to focus on the disruption of galactic discs in relation to the scatter and slope of the relation. \\citeauthor{Cattaneo2005} grow bulges both as we do \\emph{and} from disk instabilities, which interestingly produces a bi-modal \\bb\\ distribution at high redshift.  For simplicity we have removed bulge growth through disk instabilities \\cite[as originally used in][]{Croton2006}, although when included we also see such bi-modality.  This bi-modal prediction of the high redshift \\bb\\ relation provides a novel test of the mechanisms through which bulge growth may occur.  Theoretical arguments and numerical work have demonstrated that galaxy mergers are capable of simultaneously triggering growth in both bulges and black holes in a way so as to jointly reproduce many of their properties currently observed in the local universe.  If mergers are the primary drivers of black hole and bulge growth in the galaxy population, then we have shown one should expect to see an evolution in the \\bb\\ relation which arises from an increasing contribution of disrupted disks to bulges as the universe ages.  In this picture, \\emph{evolution in the growth of bulges drives an evolution in the \\bb\\ relation}, distinct from the growth rate of black holes.  At the very least, even if the physics governing bulge and black hole growth turns out to be much more complex and cannot be expressed in a simplified manner (as is currently assumed by most models of galaxy formation), in a $\\Lambda$CDM universe this effect should still be present and must be included in any interpretation of the \\bb\\ relation measured at different redshifts.  We await future high redshift observations, e.g. the Galaxy Evolution from Morphological Studies (GEMS) project \\citep{Rix2004}, to clarify the situation further."
    },
    "0512/astro-ph0512361_arXiv.txt": {
        "abstract": "{We show that aperiodic and quasiperiodic variability of bright LMXBs -- atoll and Z- sources,  on $\\sim$ sec -- msec time scales  is caused primarily by variations of the luminosity of the boundary layer. The emission of the accretion disk is  less variable on these time scales and its power density spectrum follows $P_{\\rm disk}(f)\\propto f^{-1}$  law, contributing to observed flux variation at low frequencies and low energies only. The kHz QPOs have the same origin as variability at lower frequencies, i.e. independent of the nature of the \"clock\", the actual luminosity modulation takes place on the neutron star surface. The boundary layer spectrum remains nearly constant in the course of the luminosity variations  and is represented to certain accuracy by the  Fourier frequency resolved spectrum. In the investigated range of $\\dot{M}\\sim (0.1-1) \\dot{M}_{\\rm Edd}$ it depends weakly on the global mass accretion rate and in the limit $\\dot{M}\\sim \\dot{M}_{\\rm Edd}$ is close to Wien spectrum with $kT\\sim 2.4$ keV. Its independence on the global value of $\\dot{M}$ lends support to the  theoretical suggestion by \\citet{inogamov99} that the  boundary layer is radiation pressure supported. \\\\ Based on the knowledge of the boundary layer spectrum we attempt to relate the motion along the Z-track to changes of physically meaningful parameters. Our results suggest that the contribution of the boundary layer to the observed emission decreases along the Z-track from conventional $\\sim 50\\%$ on the horizontal branch to a rather small number on the normal branch. This decrease can be caused, for example, by obscuration of the boundary layer by the geometrically thickened accretion  disk at $\\dot{M}\\sim\\dot{M}_{\\rm Edd}$. Alternatively, this can indicate significant change of the structure of the accretion flow at $\\dot{M}\\sim\\dot{M}_{\\rm Edd}$ and disappearance of the boundary layer as a distinct region of the significant energy release associated with the neutron star surface. ",
        "introduction": "\\label{sec:intro}  Accreting neutron stars in low mass X-ray binaries (LMXB) are among the most luminous compact X-ray sources in the Milky Way. A number of them have luminosities exceeding $\\sim {\\rm few}\\times 10^{38}$ erg/s and presumably  accrete matter at the level close to the critical Eddington accretion rate. In the bright state these sources have rather soft X-ray spectra, indicating that their X-ray emission is predominantly formed in the optically thick media. Similar to black holes,  at lower luminosities, $\\log(L_{\\rm X})\\la 36.5-37$, neutron stars undergo a transition to the hard spectral state \\citep[e.g.][]{barret01}. The energy spectra in this state point at the low optical depth in the emission region, indicating a significant change of the geometry of the accretion flow.  In the soft spectral state, the commonly accepted picture of accretion at not too extreme values of  accretion rate has two main ingredients -- the  accretion disk (AD) and the boundary layer (BL). While in the disk matter rotates with nearly Keplerian velocities, in the boundary layer it decelerates down to the spin frequency of the neutron star and settles onto its surface. For the typical neutron star spin frequency ($\\la500-700$Hz) comparable amounts of energy are released in  these two regions \\citep{ss86,sibg00}. This picture is based on rather obvious qualitative expectations as well as more sophisticated theoretical considerations and numerical modeling \\citep{ss86,kluzniak,inogamov99,sibg00}. It has been receiving, however, little direct observational confirmation.  Due to similarity of the spectra of the accretion disk and boundary layer the total spectrum has a smooth curved shape, which is difficult to decompose into separate spectral components \\citep{mitsuda84,white88,disalvo01,done02}. This made  application of physically motivated spectral models to the description of observed spectra of luminous neutron stars difficult, in spite of very significant increase in the sensitivity of X-ray instruments. Not surprisingly, the best fit parameters derived from the data of different instruments and, correspondingly, the inferred values of the physically meaningful quantities  are often in contradiction to each other.  This ambiguity can be resolved if the spectral information is analysed together with  timing data. Early results of \\citet{mitsuda84} and \\citet{mitsuda86} suggested that the boundary layer and accretion disk may have different patterns of spectral variability. Based on the TENMA data, they studied the difference between the spectra averaged at different intensity levels -- that restricted the range of accessible time scales to $\\ga 10^3$ sec. \\cite{gilfanov03} and \\cite{mikej05} have exploited the technique of Fourier frequency resolved spectroscopy \\citep{freq_res99} to study spectral variability of luminous LMXBs in a broad range of time scales, including kHz QPO.  Their findings are reviewed and discussed below.       \\begin{figure} \\includegraphics[width=0.5\\textwidth]{freqres_gx340.ps} \\caption{Average and frequency resolved spectra of GX340+0 on the horizontal branch of the color-color diagram. The solid lines show the Comptonization  spectrum with parameters similar to those given in the Section \\ref{sec:bl_spectrum}. \\label{fig:freqres_gx340}} \\end{figure}   \\begin{figure} \\includegraphics[width=0.5\\textwidth]{lags_gx340.ps} \\caption{Phase lags in GX340+0 on the horizontal branch as function of energy ({\\em upper panel}) and Fourier frequency ({\\em lower panel}). The energy dependent phase lags were computed in the 1--32 Hz frequency range, the frequency dependent lags are between 3--6.5 keV and 6.5--13 keV energy bands. The phase is normalized to 0--1 interval. \\label{fig:lags_gx340}} \\end{figure} ",
        "conclusions": "\\begin{enumerate}  \\item The X-ray variability in luminous LMXBs on the short timescales, $f\\ga 1$ Hz, is caused by variations of the luminosity of the boundary layer. The accretion disk emission is significantly less variable at these frequencies. The BL spectrum remains nearly constant in the course of luminosity variations and its shape equals the frequency resolved spectrum, i.e. can be directly derived from the timing data (Fig.\\ref{fig:disk_bl}).  \\item In the investigated range of the mass accretion rate $\\dot{M}\\sim (0.1-1)\\dot{M}_{\\rm Edd}$, the boundary layer spectrum  depends weakly on $\\dot{M}$.  Its shape is remarkably similar in atoll and Z-sources (Fig.~\\ref{freq_spectra}), despite an order of magnitude difference  in the mass accretion rate. Data indicates that in the limit of high $\\dot{M}\\sim\\dot{M}_{\\rm Edd}$, the boundary layer spectrum can be described by Wien spectrum with $kT\\approx 2.4$ keV (Fig.~\\ref{fig:bl_alongz}). At lower values of $\\dot{M}$ the spectra are better described by model of saturated Comptonization with electron temperature of $\\sim 2-4$ keV and Comptonization parameter $y\\sim 1$. Weak dependence of the BL spectrum  on the global value of $\\dot{M}$ lends support to the  theoretical suggestion by \\citet{inogamov99} that the  boundary layer is radiation pressure supported.  \\item The  kHz QPOs appear to have the same origin as aperiodic and quasiperiodic variability at lower frequencies. The msec flux modulations originate on the surface of the neutron star although the kHz ``clock''  might reside in the disk or be determined by the  disk -- neutron star interaction.   \\item We  attempt to relate the motion of Z-sources along the Z-track to changes in the values of the physically meaningful parameters. Our results  suggest that the contribution of the boundary layer  component to the observed emission decreases along the Z-track from the conventional value of $\\sim 50\\%$ on the horizontal branch to a rather small number at the end of the normal branch (Fig.\\ref{blcontr},\\ref{z_model}). The main difference of our approach from previous attempts is in the a priori knowledge of the shape of the boundary layer spectrum. This allowed us to avoid ambiguity of the spectral decomposition into boundary layer and disk components.   \\end{enumerate}"
    },
    "0512/astro-ph0512157_arXiv.txt": {
        "abstract": "The scientific data collected during slews of the \\xmm\\ satellite are used to construct a slew survey catalogue. This comprises of the order of 4000 sources detected in the EPIC-pn $0.2-12$\\,keV band with exposures of less than 15\\,s and a sky coverage of about 6300 square degrees (source density $\\sim 0.65$ per square degree). Below 2 keV the sensitivity limit is comparable to the ROSAT PSPC All-Sky Survey and the \\xmm\\ slew survey offers long-term variablity studies. Above 2 keV the survey will be a factor of 10 more sensitive than all previous all-sky X-ray surveys. The slew survey is almost complementary to the serendipitous survey compiled from pointed \\xmm\\ observations. It is aimed to release the first source catalogue by the end of 2005. Later slew observations and detections will continuously be added. This paper discusses the \\xmm\\ slew survey also in a historical context. ",
        "introduction": "The development of new space instrumentation for X-ray astronomical applications aims towards higher collecting areas, higher spatial resolution, and higher spectral resolution. This is related to smaller and smaller fields of view. Observations like {\\em Deep Surveys} (e.g.\\ in the directions of the Lockman Hole, the Hubble Deep Field North, etc.) -- with exposures of the order of $10^6$\\,s until they reach the confusion limit -- can help to study the faint end of luminosity functions and thus to analyse the most abundant sources in the Universe.  {\\em All-Sky Surveys}, on the other hand, with shallow exposures but a large sky coverage, are the proper database to study rare objects, with a small surface number density, and also the bright end of luminosity functions. As an example, the ROSAT All-Sky Survey (RASS) with its {\\em Bright Source Catalogue} with 18811 sources in the $0.1-2.4$\\,keV band (Voges et al.\\ 1999a,b) exceeded any previous large-area X-ray survey in terms of sensitivity and number of new sources.  \\begin{figure}[!thb] \\centerline{\\psfig{file=s05_01.ps,width=0.46\\textwidth,angle=270.0,clip=}}% \\vspace*{-0.5mm} \\caption[]{Sky distribution of 465 EPIC-pn slews (FF, eFF, LW modes) in ecliptic coordinates. Note that slews are performed close to great circles due to solar angle constraints.} \\label{fig:slewpaths} \\end{figure}   {\\em Slew Surveys} play an intermediate role between specially designed all-sky programs and dedicated pointed observations. The {\\em Einstein} IPC slew survey \\citep{elvis1992} covered half of the sky in the $0.5-3.5$\\,keV band with a sensitivity of about $3\\times 10^{-12}$erg\\,cm$^{-2}$\\,s$^{-1}$ (0.1 IPC cts/\\,s$^{-1}$) and the resulting catalogue contained 819 sources. 15\\% of those had no counterpart in the (slightly softer) RASS \\citep{schachter1993}. All-sky survey extensions into (and beyond) this harder energy range have been proposed like ABRIXAS as a path-finder for \\xmm\\ \\citep{truemper1998} and ROSITA \\citep{predehl2003}, but the first failed and the latter is not yet approved.  \\xmm\\ with its superior collecting area would also be an ideal mission for serendipitous science as already in short exposures during slews enough photons could be detected for a classification of the sources (X-ray colours, extents, etc.). 5 years before the launch the potential of such a slew survey was outlined \\citep{lumb1995}. Pre-launch predictions and feasibility studies, however, were based on assumptions on the slew rate of smaller than $20^\\circ$ per hour \\citep{jones1998, lumb1998, joneslumb1998}; the actual slew rate of $90^\\circ$ per hour reduces the typical number of photons per source but increases also the sky coverage (faster slew gives more possible observations and thus more slews). ",
        "conclusions": "It has been shown that the \\xmm\\ EPIC-pn data collected during slews represent an important scientific database. The catalogue currently under construction will provide a complement to catalogues compiled from pointed observations (1XMM, 2XMM). Moreover, not only point sources but also extended sources had been detected in the slew survey \\citep{lazaro2005}. While supernova remnants are most likely already detected in the RASS, harder sources like clusters of galaxies may be new extended objects originating from the slew survey.  In further versions of the slew survey catalogue it is planned to recover part of the slews that were disregarded due to high background by selecting periods of low background (using good time intervals similar to pointed observations). The slew sky coverage will increase and therefore serendipitous overlaps with pointed observations and with other slew paths will increase as well. This will greatly enhance the possibility of time variability studies."
    },
    "0512/astro-ph0512227_arXiv.txt": {
        "abstract": "We compute a series of three-dimensional general relativistic magnetohydrodynamic simulations of accretion flows in the Kerr metric to investigate the properties of the unbound outflows that result. The overall strength of these outflows increases sharply with increasing black hole rotation rate, but a number of generic features are found in all cases.  The mass in the outflow is concentrated in a hollow cone whose opening angle is largely determined by the effective potential for matter orbiting with angular momentum comparable to that of the innermost stable circular orbit.  The dominant force accelerating the matter outward comes from the pressure of the accretion disk's corona.  The principal element that shapes the outflow is therefore the centrifugal barrier preventing accreting matter from coming close to the rotation axis. Inside the centrifugal barrier, the cone contains very little matter and is dominated by electromagnetic fields.  The magnetic fieldlines inside the cone rotate at a rate tied closely to the rotation of the black hole, even when the black hole spins in a sense opposite to the rotation of the accretion flow.  These fields carry an outward-going Poynting flux whose immediate energy source is the rotating spacetime of the Kerr black hole. When the spin parameter $a/M$ of the black hole exceeds $\\simeq 0.9$, the energy carried to infinity by these outflows can be comparable to the nominal radiative efficiency predicted in the Novikov-Thorne model. Similarly, the expelled angular momentum can be comparable to that accreted by the black hole.  Both the inner electromagnetic part and the outer matter part can contribute in significant fashion to the energy and angular momentum of the outflow. ",
        "introduction": "Whether from supermassive objects in galactic nuclei or stellar mass objects in Galactic binaries, outflows---often relativistic---are frequently seen from black holes.  However, there is little that is known with confidence about their nature and dynamics.  The essential ingredients are, of course, well-known:  magnetic field, accretion, and rotation.  But basic questions such as the forces that drive the jets, the mechanisms that regulate their content, and the constraints that collimate them remain open.  Magneto-centrifugal effects have received the greatest attention (Blandford \\& Znajek 1977; Blandford \\& Payne 1981; Shibata \\& Uchida 1985; Punsly \\& Coroniti 1990), but there are many possible variations on this idea.  In most such studies, magnetic field boundary conditions are guessed and their consequences derived without much consideration of how the posited field structures would arise in the context of accretion flows.  We approach these questions from the opposite point of view: we simulate numerically how accretion dynamics self-consistently create magnetic field and explore the outflows that result.  Employing the general relativistic three dimensional MHD code described in De Villiers \\& Hawley (2003),  we have previously investigated the general character of accretion onto black holes in a series of simulations for which an overview is presented in De Villiers, Hawley \\& Krolik (2003; hereafter Paper~I).  The magnetic field structures that develop in these accretion flows were described in Hirose et al. (2004; hereafter Paper~II).  We discussed the resulting outflows in De Villiers et al. (2005; hereafter Paper~III), and the dynamics of the inner disk in Krolik, Hawley, \\& Hirose (2005; hereafter Paper~IV). In this paper we return to a more in-depth consideration of the type of unbound outflows that were discussed in Paper~III.  We will present a more quantitative and dynamically-oriented analysis of the outflows, and draw on several new simulations designed to elucidate their properties. We will focus our attention on three issues in particular: the underlying principles governing the material component in the outflow; the dynamics and spin-dependence of the outflow's electromagnetic segment; and the relationship between the electromagnetic contribution and the class of ideas placed under the heading of the Blandford-Znajek mechanism. ",
        "conclusions": "In this paper we have examined the unbound axial outflows (jets) that develop within GRMHD accretion simulations.  What is noteworthy in these models is that the jets develop as a natural result of the accretion process, rather than as a direct consequence of conditions imposed on the simulation.  The jets have two major components:  a matter-dominated outflow that moves at a modest velocity ($v/c \\sim 0.3$) along the centrifugal barrier surrounding an evacuated axial funnel, and a highly-relativistic Poynting flux dominated jet within the funnel. The funnel wall jet is accelerated and collimated by magnetic and gas pressure forces in the inner torus and the surrounding corona. The Poynting flux jet results from the formation of a large scale radial magnetic field within the funnel.  This field is spun by the rotating spacetime of the black hole, hence the energy for this jet component comes not directly from accretion but from the black hole's rotation, although the field mediating this energy extraction is a product of accretion. The energies in both components of the jet are a function of black hole spin; greater spins yield greater energy.  The total jet power can be a significant fraction of the accretion energy estimated by the product of the mass accretion rate and the specific binding energy of the last stable orbit.  We next consider the broader implications of these results within the context of several of the standard theoretical categories for jets, namely jet acceleration, power, collimation, and the Blandford-Znajek mechanism.  \\subsection{Accelerating forces}  It has been expected for many years (since Blandford \\& Znajek 1977), that within the axial funnel the spacetime rotation enforced by spinning black holes can drive a significant electromagnetic outflow.  Our results are clearly in line with that expectation.  The outflow-generation process within the accretion flow itself that has for many years received the greatest attention has been magnetocentrifugal acceleration, such as envisioned in the model of Blandford \\& Payne (1981), but this mechanism does not appear to be important here.  On the other hand, the main driving force we identify for the funnel wall jet---a high-pressure corona squeezing material against an inner centrifugal wall---has received very little attention in the past. In consequence of this compression, the pressure at the base of the matter outflow is kept high enough to create an outward radial pressure gradient that provides the momentum for the outflow.  It is also worth commenting that these forces act over a substantial radial range, from a few gravitational radii out to many tens.  \\subsection{Outflow strength}  To measure the relative power in the unbound outflow, we have compared the jet energy flux with the mass accretion rate into the hole over the last 70\\% of the simulation.  The simulation has no radiative losses, so we cannot say what the radiative efficiency of the accretion disk would be.  The Novikov-Thorne efficiency provides a convenient standard of comparison.  This nominal efficiency corresponds to the binding energy of the innermost stable circular orbit, although the actual radiative efficiency could well be different because of, for example, magnetic torques in the plunging region (Krolik 1999; Gammie 1999), or heat advection into the hole.  Judged by this measure, the energy carried off in the outflow is considerable.  In two cases ($a/M = -0.9$ and 0.99), the electromagnetic power at $r=100M$ is, by itself, comparable to the Novikov-Thorne prediction of radiated power.  In the other high spin cases, it is about a quarter of the Novikov-Thorne number.  One should remember, however, that energy can be (and often is) exchanged between the fields and the matter, so these measures of the electromagnetic power in the outflow are quantitatively accurate {\\it only} at the fiducial location where they are measured.  Because the rate of energy loss in the matter outflow is usually a few times as great as the electromagnetic contribution, the total is typically comparable to the classical prediction of radiated power.  Another way to gain a feel for the magnitude of the electromagnetic power in the outflow is via dimensional analysis.  The Poynting flux at the inner boundary is $\\propto {\\cal B}^r {\\cal B}^\\phi r_{\\rm in} \\omega_{\\rm in}$ (cf. eqn.~\\ref{eqn:bzpoynting}); but what determines the magnetic field strength?  Indeed, this has been the key question examined by such works as Ghosh \\& Abramowicz (1997) and Livio, Ogilvie, \\& Pringle (1999). These authors argued that the strength of the field responsible for a Blandford-Znajek process must be set by the fields in the accretion flow, and this may limit the total jet power.  At radii in the disk well beyond the marginally stable region, angular momentum conservation ensures that $\\langle {\\cal B}^r {\\cal B}^\\phi \\rangle \\propto \\dot M \\Omega/h$ if magnetic torques are the primary agent of accretion, while the competition between orbital shear and other stresses tends to make $\\langle ||b||^2 \\rangle \\propto \\langle {\\cal B}^r {\\cal B}^\\phi \\rangle$.  In these simulations, in which $h/r$ is common to all, but the black hole spin varies, if we wish to analyze the spin-dependence of the Poynting flux near the event horizon, it makes sense to examine the ratio of the magnetic field intensity to the mass accretion rate at that location. We find that the time- and surface area-averaged ratio $||b||^2/\\dot M$ grows rapidly with increasing black hole spin, very nearly $\\propto (1 - |a/M|)^{-1.1}$ (Fig.~\\ref{fig:bsqeff}).  Some of this growth is attributable to mere compression, as the event horizon shrinks as the spin rate rises.  However, the dependence is much stronger than anything simple geometry could explain.  It appears that the strong frame-dragging associated with extreme Kerr spacetimes leads to substantial field amplification as well.  The rapid growth in electromagnetic power output with faster black hole rotation is predominantly due to this effect.  \\clearpage  \\begin{figure} \\centerline{\\psfig{file=./f12.ps,angle=90,width=5.0in}} \\caption{The ratio $||b||^2/\\dot M$ averaged over time, with the magnetic field intensity evaluated at the inner boundary. \\label{fig:bsqeff}} \\end{figure}  \\clearpage  \\subsection{Jet shape}  One of the key issues in jet formation and propagation is the degree to which the jets are collimated.  Of course, jets are observed on far larger scales than we are modeling in these simulations, but the degree of collimation seen here is nevertheless of some interest. Generally speaking, whatever collimation there is occurs in the inner regions of the flow, where the inner disk and hot corona provide confinement.  Relatively close to the black hole ($r \\lesssim 10M$), the boundary between the magnetically-dominated Poynting flux jet and the funnel wall outflow follows the shape of the centrifugal barrier. At larger distances, however, the matter flows almost straight outward, so that the funnel is nearly conical, while the contours of constant effective potential gradually move toward smaller polar angle (Fig.~\\ref{fig:momflux}).  The precise shape and collimation of the jet is somewhat uncertain because the outer boundary of the funnel-wall, matter-dominated funnel wall outflow is somewhat indistinct.  We see a smooth transition as a function of polar angle between mildly relativistic outflowing unbound matter and slightly slower, but bound, coronal matter.  On the other hand, the boundary between the low-density magnetically-dominated funnel interior and the higher-density funnel-wall outflow is sharp and clear. It is possible that magnetic forces might provide additional collimation on larger scales, but this remains to be seen in future simulations.   \\subsection{Comparison to other models and the ``Blandford-Znajek\" mechanism}  In the astrophysical literature, the general idea that energy can be conveyed from rotating black holes to the outside world via magnetic connections is often dubbed the ``Blandford-Znajek mechanism\" because they were the first to discuss a specific version of this process.  In our simulations, when all is said and done, the funnel interior is dominated by a magnetic field that would be purely radial but for the rotation created by its connection to the black hole's ergosphere.   Thus, any toroidal field in the funnel, and consequently any Poynting flux, is due entirely to general relativistic dynamics peculiar to spinning black holes.  In other words, the outward Poynting flux must be due to a process that could fit under the general rubric of ``Blandford-Znajek\".  On the other hand, we also find specific differences from the classical Blandford-Znajek model, specifically the particular idealized picture presented in their original paper (Blandford \\& Znajek 1977).  For example, although the funnel region is magnetically-dominated, it is {\\it not} in general in a state of force-free equilibrium.  Indeed, the very large fluctuations that continually occur in the outflow show that it is never in any state of equilibrium, force-free or otherwise.  In addition, as originally shown in Phinney (1983), the collapse of all wave speeds to $c$ in the exact force-free limit means that to understand the dynamics properly requires retention of a finite inertia.  Moreover, whereas in the classical Blandford-Znajek model there was zero accretion (the only function of nearby orbiting matter is to support the currents sustaining the magnetic field), in our picture the electromagnetic luminosity in the outflow takes place entirely as a byproduct of accretion.  In the classical model the magnetic field was a relic of plasma inflow that took place in the distant past; because there was no continuing accretion, the energy output could be attributed unambiguously to the black hole's rotation.  In our simulation, the magnetic field is a direct result of on-going accretion.  While the proximate energy source for the jet is the black hole's rotation, accretion replenishes both the black hole's mass and angular momentum, and whether the rotational energy of the black hole actually decreases depends on details such as the radiative efficiency and, possibly, the accretion of photons (Krolik 2001; cf. Thorne 1974).  In our $a/M = 0.99$ simulation, the net angular momentum accreted is so small that, for reasonable values of the radiative efficiency, the system is very near the edge of an actual decrease in rotational kinetic energy.  Even if the black hole's rotational energy does not decrease, there can nevertheless be a significant increase in the efficiency of accretion.  We find that for rapidly rotating black holes the total power in the Poynting flux jet can be comparable to the nominal accretion power.  A number of other numerical studies of the generic Blandford-Znajek mechanism have been carried out recently.  One type of simulation has been devoted to examining the idealized problem of a rotating black hole embedded in thin plasma containing a specified external magnetic field (Koide 2003; Semenov et al. 2004; Komissarov 2005).  These simulations focus on points of basic principle, such as the extent to which accretion of either matter or electromagnetic fields with negative energy-at-infinity plays a role.  Komissarov (2005) particularly notes that plasma effects within the ergosphere were significant early on in his simulation while the initial plasma is accreting into the hole, but that the long-term steady state mechanism resembles closely the classic Blandford-Znajek model.  The simulations of McKinney \\& Gammie (2004) were more similar in their focus to ours, as they too were principally concerned with the connection between continuing accretion and the electromagnetic power output of black holes.  For this reason, their (axi-symmetric) simulations employed initial conditions resembling ours (isolated initial gas torus containing field loops).  They found good agreement with several aspects of the classical Blandford-Znajek model for slowly rotating holes, but growing differences as the spin of the hole increases and the field becomes less monopolar.  That this should be so is not too surprising, as the specific predictions of the Blandford-Znajek model were developed on the basis of a perturbation expansion in $a/M$.  In regard to the global energy and angular momentum budgets, McKinney \\& Gammmie's results qualitatively resemble ours.  They also saw that as the hole loses energy and angular momentum electromagnetically, its total mass and angular momentum are replenished by accretion.  Likewise, they, too, found that the overall efficiency (energy released to infinity per rest-mass accreted) can be significantly enhanced by the energy carried off in the jet.  McKinney (2005) expanded upon these results and proposed an analytic fitting function for the electromagnetic efficiency of the jet as a function of black hole rotation rate, suggesting that this efficiency scales $\\propto \\Omega_H^5$, where $\\Omega_H$ is the rotation rate of the black hole.  If the Poynting flux were due simply to rotation of a magnetic dipole whose magnitude is independent of black hole rotation but $\\propto \\dot M^{1/2}$, the efficiency would, of course, increase $\\propto \\Omega^4$.  The slightly faster dependence McKinney (2005) found suggests either that there is a mixture of a higher multipole component or that the magnitude of the magnetic dipole increases slowly with increasing black hole spin.   The choice of this functional form was motivated by the hope that electromagnetic radiation in this case resembled classical dipole radiation, in which the luminosity is $\\propto B^2 \\Omega^4$.   Whether this applies in the present context depends, of course, on whether dipole field geometry is a reasonable approximation and on whether the mean field strength at the horizon produced by a fixed accretion rate is independent of  spin.  Unfortunately, for a variety of technical reasons it is difficult to compare quantitatively the results of different numerical simulations. Without the inclusion of additional physics in the basic equations, there is no absolute density scale in the gas.  Although efficiencies are well-defined (ratios of energies to masses within the simulation), one cannot define an absolute accretion rate.  Even the comparison of efficiencies from one type of simulation to another poses problems. For example, two dimensional simulations tend to over-emphasize axisymmetric MRI modes.  As a result, they tend to have large accretion rates early on, followed by a rapid decline as the magnetic field and accompanying turbulence die out, as expected from the anti-dynamo theorem.  Thus, when computing the efficiency in a two dimensional simulation simulation, one must be careful about choosing an appropriate interval over which to average the accretion rate.  A second problem is that, as discussed previously, energy can be exchanged between the electromagnetic and matter components of the outflow, so the purely electromagnetic power is a function of radius. To cross-compare different simulations, it is therefore important to use the same fiducial distance for evaluating the electromagnetic power.   Lastly, jet properties may depend on the vertical thickness of the accretion disk, a parameter determined by the initial temperature of the disk and the treatment of the energy equation in the simulation.  The gross accretion rate in general scales $\\propto h^2$, but we do not know how the magnetic field in the vicinity of the black hole scales with disk temperature.  Further simulational work will be required to determine such effects, although because the strength of the axial funnel field is determined by the accretion flow, it is a reasonable expectation that the jet power might scale with the accretion rate (Ghosh \\& Abramowicz 1997).  Keeping these {\\it caveats} in mind, we can compare our electromagnetic efficiency in the jet with the scaling relation proposed by McKinney (2005) by modifying his relation to use the rotation rate of our inner grid radius instead of the rotation rate of the horizon.  The quantitative match between our data and McKinney's formula is not particularly good, although qualitatively we, too, see a strong increase in efficiency as black hole spin increases.  Our results are better described by the simple formula \\begin{equation} \\eta_{em} \\simeq 0.002/(1 - |a/M|) , \\end{equation} a form that describes both the high and lower spin models.  Of course, without any clear theory to predict the strength of the magnetic field on the event horizon relative to the mass accretion rate as a function of black hole rotation rate, one cannot expect to find a general formula for total jet power.  Finally, we would like to close this paper by pointing out that, although the Blandford-Znajek mechanism was originally thought of as a scheme for powering jets, outflows are not the only place where electromagnetic power from the black hole can go.  In this paper we have concentrated on jet properties.  However, preliminary analysis of the accretion flow in these simulations indicates that the Poynting flux and electromagnetic angular momentum flux from the black hole into the disk can be comparable to those within the outflow.  These fluxes and their consequences will be examined in subsequent work."
    },
    "0512/astro-ph0512588_arXiv.txt": {
        "abstract": "We have defined a complete sample of 228 southern radio sources at 408 MHz with integrated flux densities $S_{\\rm 408} > 4.0$\\,Jy, Galactic latitude $|b|>10^{\\circ}$ and declination $-85^{\\circ} < \\delta < -30^{\\circ}$.  The main finding survey used was the Molonglo Reference Catalogue.  We describe in detail how the Molonglo Southern 4 Jy sample (MS4) was assembled and its completeness assessed.  Sources in the sample were imaged at 843\\,MHz with the Molonglo Observatory Synthesis Telescope to obtain positions accurate to about $1''$, as well as flux densities and angular sizes; follow-up radio and optical observations are presented in Paper II.  Radio spectra for the MS4 have been compiled from the literature and used to estimate flux densities at 178\\,MHz. The strong-source subset of MS4, with $S_{\\rm 178} > 10.9$\\,Jy (SMS4), provides a southern sample closely equivalent to the well-studied northern 3CRR sample.  Comparison of SMS4 with 3CRR shows a reassuring similarity in source density and median flux density between the two samples. ",
        "introduction": "Well defined, complete surveys are essential starting points for learning about the global characteristics of astronomical sources.  At radio frequencies, the easiest way to select a large, well defined sample is to choose all objects brighter than a given flux density at a given observing frequency.  Such a sample has the advantage of containing the objects with the highest signal-to-noise ratio.  Because any one finding survey is likely to be affected either by confusion or by over-resolution for some of the sources, it is often necessary to use more than one survey to define a flux density-limited sample.  Even then the sample may not be complete, but may, for example, be missing sources (or portions of sources) of low surface brightness or very large angular size.  One of the best studied radio samples in the past few decades has been the northern hemisphere Third Cambridge Catalogue (3C; \\citealp{3c}) and its revised versions the 3CR \\citep{3cr} and 3CRR \\citep{lrl}. Being selected at low frequency (159 and later 178\\,MHz), it contains a large proportion of extended, steep-spectrum sources.  Because of the high flux-density limit ($S_{\\rm 178}$) of around 11\\,Jy, the sample contains a large proportion of very powerful sources.  The most complete version, the 3CRR, is the one discussed in this paper.  As so many studies of high-power radio sources have been based on this one sample, it is essential to have a comparison sample to test whether the 3CRR is truly representative.  This paper is concerned with the definition of such a comparison sample, selected at low frequency to have properties similar to 3CRR.  Increasing the sky coverage is more important at high flux densities where the total number of sources is limited.  This sample was selected at 408\\,MHz to contain the $\\sim 200$ brightest extragalactic objects south of $\\delta = -30^{\\circ}$.  The sample is defined in Section~\\ref{sec.sampdef}.  Follow-up observations of the whole sample with the Molonglo Observatory Synthesis Telescope at 843\\,MHz are described in Section~\\ref{sec.mostobs}.  These observations have provided more accurate positions and angular sizes, as well as flux densities in a gap in the radio spectrum.  The definition of a strong subsample at 178\\,MHz, for comparison with 3CRR, is described in Section~\\ref{sec.est178}.  In \\papern\\/ \\citep{paper2} we present follow-up imaging at 5\\,GHz of the more compact sources with the Australia Telescope Compact Array, together with optical identifications using the UK Schmidt sky survey and CCD images from the Anglo-Australian Telescope. ",
        "conclusions": "A summary of the 843\\,MHz data for the MS4 sample, along with flux density measurements from the literature at other frequencies, is given in Table~\\ref{tab5}.  Although different measurements in the literature have used (nominally) different flux-density scales, no correction has been made; the figures in the table are quoted {\\it directly\\/} from the papers cited.  The columns of Table~\\ref{tab5} are as follows:  \\begin{enumerate}  \\item Source name, taken from the MRC, unless the source has more than one MRC entry, in which case the name is from the original Parkes catalogue \\citep{rad.B1,rad.P1} or Parkes 2700\\,MHz catalogue (\\citealp{2700.last} and references therein).  \\item \\label{tab5.ra} Right ascension in J2000 coordinates at 843\\,MHz.  \\item \\label{tab5.dec} Declination in J2000 coordinates at 843\\,MHz. The reference for the radio position in columns~\\ref{tab5.ra} and \\ref{tab5.dec} is the same as for the 843\\,MHz flux density in column~\\ref{tab5.843}.  \\item \\label{178col} Estimated flux density at 178\\,MHz, as described in Section~\\ref{sec.est178}, and the order of the polynomial used to calculate the estimate.  \\item Flux density at 80\\,MHz; the observing frequency is 85.5\\,MHz, if taken from MSH, or 80\\,MHz, if measured with the Culgoora Circular Array.  The latter flux densities are from the compilation of \\citet{rad.S12}, with revised calibration.  \\item \\label{tab5.s408} Flux density at 408\\,MHz, taken from the MRC, except for some extended sources.  In general, flux densities were taken from published data in the following order of preference: SM75, MC4, \\citet{rad.W6}, the Parkes catalogues, \\citet{rad.E1}.  Flux densities for a few sources were estimated using the Molonglo Transit Catalogue or from the radio spectrum.  \\item \\label{tab5.843} Flux density at 843\\,MHz, measured with MOST. The values of \\citet{rad.C1} were used for MOST calibrators, as were those of \\citet{rad.S16} for six giant radio sources.  For the remaining sources the flux density was taken either from JM92 or from the present study, depending on which image had the higher dynamic range.  \\item Flux density at 1400\\,MHz.  Data are mostly from Parkes (1410\\,MHz: \\citealt{rad.B1,rad.P1,rad.W2}), from Owens Valley (1425\\,MHz: \\citealt{rad.F1,rad.F2}), or from the NRAO VLA Sky Survey catalog (NVSS; 1400\\,MHz: \\citealt{rad.C6}).  For a few sources data were taken from the Instituto Argentino de Radioastronomia (IAR) 30\\,m telescope (1410\\,MHz: \\citealt{rad.Q1,rad.Q2,rad.Q3}), from the Fleurs Synthesis Telescope, or the Australia Telescope Compact Array. The flux-density scale for the Parkes catalogues \\citep{rad.B1,rad.P1} and Owens Valley data was that of \\citet{flux408.ckl}, whereas the data of \\citet{rad.W2} and the IAR 30\\,m data were on the scale of \\citet{wills.scale}.  At 1410\\,MHz the scale of \\citet{wills.scale} is 1.08 times the scale of \\citet{flux408.ckl} \\citep{baars77}.  NVSS used the flux density scale of \\citet{baars77}.  \\item Flux density at 2700\\,MHz, measured with the Parkes telescope, except for one source which was observed at 2640\\,MHz with Owens Valley \\citep{rad.R1}.  2650\\,MHz flux densities are from the original Parkes catalogue or from \\citet{rad.W2}, while 2700\\,MHz values are from the Parkes 2700\\,MHz survey or from \\citet{rad.W2}.  The Parkes 2700\\,MHz survey used a scale in which the peak flux density of Hydra~A was 23.5\\,Jy; \\citet{rad.W2} used the flux-density scale of \\citet{wills.scale}.  \\item Flux density at 5000\\,MHz, measured with the Parkes telescope. Values at 4850\\,MHz are from the Parkes-MIT-NRAO survey (PMN: \\citealt{rad.G1,rad.W7}), those at 5000\\,MHz are from \\citet{rad.W2}, and those at 5009\\,MHz are from \\citet{rad.W2}, or the Parkes 2700\\,MHz survey, or \\citet{rad.S1}, or \\citet{rad.S2}.  The PMN survey used the flux density scale of \\citet{baars77}; \\citet{rad.W2} used the scale of \\citet{wills.scale}; the Parkes 2700\\,MHz surveys used a 5 GHz scale in which the peak flux density of Hydra~A was $S_{\\rm 5009} = 13.05$\\,Jy, as did \\citet{rad.S1} and \\citet{rad.S2}.  \\item Radio spectral index $\\alpha$ measured between 408 and 2700\\,MHz, defined in the sense $S_{\\nu} \\propto \\nu^{\\alpha}$. For sources without measurements at both those frequencies, the spectral index was measured either between 408 and 5000\\,MHz or between 843 and 5000\\,MHz.  \\item \\label{tab5.las} Largest angular size (LAS) at 843\\,MHz.  The reference is the same as for the 843\\,MHz flux density, except for a few sources, noted in Section~\\ref{sec.com843}.  We have imaged all sources with ${\\rm LAS}<35''$ with the ATCA at 5\\,GHz (\\papern\\/), with the exception of MRC~B0521$-$365, MRC~B0743$-$673, MRC~B1740$-$517, MRC~B1814$-$519, \\\\ MRC~B2153$-$699, and MRC~B2259$-$375.  \\item Position angle of extension at 843\\,MHz, defined in degrees east of north (modulo 90).  The reference is the same as for the 843\\,MHz flux density, except for a few sources which are noted in Section~\\ref{sec.com843}.  \\end{enumerate}  \\subsection{Comments on Individual Sources} \\label{sec.com843}  \\cind{MRC~B0208$-$512} Detected with EGRET \\citep{egret.cat2}.  \\cind{MRC~B0214$-$480} The Parkes 408\\,MHz and MSH 85.5\\,MHz flux densities are probably affected by blending with MRC~B0211$-$479. The position, angular size, and position angle in Table~\\ref{tab5} are from JM92, but $S_{843}$ was measured from MOST images from the SUMSS survey \\citep{sumss2003}.  \\cind{MRC~B0252$-$712} The spectrum turns over at low frequency. The flux density of $S_{\\rm 85.5} = 7$\\,Jy is noted as uncertain in MSH, but corresponds well with the slight spectral curvature evident at higher frequencies.  \\cind{PKS~B0319$-$45} Giant radio galaxy.  Because of the large angular size, the flux densities are uncertain.  The MSH 85.5\\,MHz measurement was affected by sidelobes of Fornax~A.  \\cind{MRC~B0320$-$373} Fornax~A.  \\cind{MRC~B0336$-$355} Blended at low resolution with the foreground source MRC~B0336$-$356.  At 2650 and 5009\\,MHz, the peak flux densities \\citep{rad.W2} were used, as they are expected to be less affected by blending.  \\cind{MRC~B0411$-$346A} The weaker source MRC~B0411$-$346B is unrelated.  \\cind{MRC~B0453$-$301, MRC~B0456$-$301} MSH\\,04$-$3{\\it 14} is a blend of these two sources.  \\cind{MRC~B0506$-$612} The 8.87 GHz flux density of 2.63\\,Jy \\citep{rad.S5} is substantially higher than expected from the radio spectrum, probably indicating high-frequency variability.  \\cind{MRC~B0511$-$484} Unequal double, with two neighbouring weaker sources, presumed to be unrelated.  The SM75 integrated value of $S_{\\rm 408} = 8.8$\\,Jy is affected by blending.  The MRC fitted value of 6.84\\,Jy was preferred but is still uncertain.  The Parkes flux densities at 2.7 and 5 GHz are also affected by blending, so the radio spectrum is not well determined.  \\cind{MRC~B0518$-$458} Pictor~A.  \\cind{MRC~B0521$-$365} Well studied blazar, detected as a gamma-ray source by EGRET \\citep{egret.cat2}.  \\cind{PKS~B0707$-$35} The flux density $S_{\\rm 5000} = 0.84$\\,Jy \\citep{rad.W5} appears anomalously high.  \\cind{MRC~B0743$-$673} The 1410 MHz flux density \\citep{rad.P1} is higher than expected from the radio spectrum, and probably affected by blending with MRC\\,B0742$-$674.  \\cind{MRC~B1017$-$421, MRC~B1017$-$426} MSH\\,10$-$4{\\it 4} is a blend of these two sources.  \\cind{MRC~B1136$-$320} In the Texas catalog \\citep{rad.D2} the double separation is quoted as 58$''$ in P.A. $-15^{\\circ}$, consistent with the values in Table~\\ref{tab5}.  \\cind{MRC~B1143$-$316} The Texas catalog \\citep{rad.D2} lists this source as a double with separation $34''$ in P.A. $86^{\\circ}$.  The MRC value of $S_{\\rm 408} = 5.77$\\,Jy is preferred to SM75's anomalously low value of 3.8\\,Jy.  \\cind{MRC~B1302$-$491} Nearby edge-on spiral galaxy NGC\\,4945, identified by \\citet{rad.M1}.  \\cind{PKS~B1318$-$434} Complex double source, with bent edge-darkened lobes, lying behind the southern lobe of Centaurus~A \\citep{rad.C4,map.cena.haynes}.  \\cind{MRC~B1322$-$427} The well studied low-luminosity radio galaxy Centaurus~A (see review by \\citealp{cena.review}).  The flux density at 843 MHz was determined from the radio spectrum.  \\cind{PKS~B1400$-$33} The only source in the sample which is not in the MRC.  An unusual extended source of low surface brightness.  There is a nearby compact source associated with the E0 galaxy NGC\\,5419 (see \\citealp{rad.E1}).  \\citet{rad.G4} suggest that the extended component may be a relic radio source associated with the poor cluster Abell S753 around NGC\\,5419.  The radio properties have been studied in detail by \\citet{rad.S17}; we have used their flux densities at 843 and 1398\\,MHz.  \\cind{MRC~B1425$-$479} Because of the angular size of $4'.5$, the values of $S_{\\rm 408}$ and $S_{\\rm 5000}$ in Table~\\ref{tab5} may be slightly underestimated.  \\cind{MRC~B1445$-$468} The published values of $S_{\\rm 80}$ --- 2\\,Jy \\citep{rad.S12} and 13\\,Jy \\citep{rad.S9} --- are grossly discrepant. We prefer the latter value, after multiplying it by 1.1 to correct the flux-density scale \\citep{rad.S11}.  \\cind{MRC~B1549$-$790} Flat-spectrum source.  The flux density at 1410\\,MHz \\citep{rad.P1} is affected by blending with MRC~B1547$-$795.  \\cind{MRC~B1814$-$637, MRC~B1817$-$640} MSH\\,18$-$6{\\it 1} is a blend of these two sources.  \\cind{MRC~B1917$-$546} Ultra-steep spectrum source.  The 1410 MHz flux density may be affected by blending with two weak neighbouring sources (Hunstead 1972).  \\cind{MRC~B1934$-$638} Archetypal Gigahertz-Peaked-Spectrum (GPS) source \\citep{disc.1934}.  \\cind{MRC~B1940$-$406} Blending with three weaker sources (JM92) makes the radio spectrum uncertain.  \\cind{MRC~B2006$-$566} Very extended, diffuse source, associated with the cluster Abell\\,3667 \\citep{aco}.  The flux density of $S_{\\rm 843} = (5.5 \\pm 0.5)$\\,Jy \\citep{rad.R2} was obtained from a MOST full-synthesis image after subtracting the contribution from background sources within the source envelope (0.6\\,Jy); the discrepancy with JM92 is probably due to the higher noise level in their image.  \\cind{MRC~B2052$-$474} % Detected with EGRET \\citep{egret.cat2}.  \\cind{MRC~B2122$-$555} The MRC position appears to be in error, lying $30''$ east of the MOST position.  The MRC flux density of 4.05\\,Jy is also lower than expected from the radio spectrum, suggesting that the observation may have been affected by a large ionospheric wedge (Hunstead 1972).  Both Parkes positions \\citep{rad.W1,rad.G1} agree with the MOST centroid.  \\cind{MRC~B2152$-$699} Some flux densities are affected by blending with MRC~B2153$-$699.  The 1415\\,MHz measurement \\citep{rad.C5} was made with the Fleurs Synthesis Telescope.  At 2650\\,MHz, the peak flux density \\citep{rad.W2} was used, to reduce any contribution from MRC~B2153$-$699.  \\cind{MRC~B2153$-$699} Double, extended along a position angle similar to that of MRC~B2152$-$699.  Flux densities are affected by blending with MRC~B2152$-$699: both the MRC value of $20.9 \\pm 1.4$\\,Jy and the SM75 value of $6.0 \\pm 0.4$\\,Jy are regarded as unreliable.  Therefore, the flux densities at 843 and 1415\\,MHz, together with the 468 MHz value of $10.7 \\pm 0.9$\\,Jy \\citep{rad.E2}, have been used to estimate the flux density at 408 MHz.    The 178\\,MHz flux densities estimated in Section~\\ref{sec.est178} have been used to define a strong source subsample of MS4 which we call SMS4.  It has been chosen to have the same flux-density cutoff, 10.9\\,Jy, as the northern 3CRR sample, and contains 137 sources, compared with 172 in 3CRR.  Comparison of SMS4 with 3CRR (Table~\\ref{tab6.medians}) shows the southern sample to have a slightly higher source density, but only at the $2.8\\sigma$ level of significance.  The difference should be treated with caution and may simply reflect biases in the way each sample was compiled. We have identified three possible causes of such a bias:  \\begin{itemize} \\item[(i)] As the spectra of many radio sources turn over at low radio frequency, extrapolation from high frequencies is more likely to overestimate than underestimate $S_{\\rm 178}$.  The survey with the Mauritius radio telescope \\citep{maurit} at 151.5\\,MHz will be valuable for checking flux densities of sources north of $\\delta = -70^{\\circ}$, and testing for such a bias in the SMS4 sample.  \\item[(ii)] Because of the steep slope of the radio source counts, a small systematic difference in flux-density scale can strongly bias the source density.  The Mauritius values will be useful for checking such effects in SMS4.  \\item[(iii)] The 3CRR may be missing sources of low surface brightness.  The angular size distributions are compared in Figure~\\ref{fig4.las}.  The median angular size of SMS4 is 32$^{+6}_{-5}$ arcsec, compared with 35.5$^{+8.7}_{-7.5}$ arcsec for 3CRR.  Although the medians are similar, the distributions appear different: the SMS4 has proportionally more sources with angular size $>300''$, and proportionally fewer sources with angular size between 100$''$ and 300$''$.  However, because of the small numbers of sources involved, it is not possible to draw firm conclusions about differences in the angular size distributions.  \\end{itemize}  The median flux density of SMS4 ($S_{\\rm 178} = 15.7$\\,Jy), is consistent with the median value of 15.6\\,Jy for 3CRR.  This is to be expected if the radio source counts are similar.  Histograms of spectral indices for the MS4, SMS4, and 3CRR samples are plotted in Figure~\\ref{fig5.alpha}.  The distributions for the MS4 sample show a longer tail towards flatter spectral indices than do those of the SMS4 or 3CRR.  This is to be expected given the higher selection frequency of the MS4, as flat-spectrum sources will be included which would be missed with a higher flux-density cutoff and lower selection frequency.  The median spectral indices, $\\alpha=-0.91$ for SMS4 and $-0.81$ for 3CRR (Table~\\ref{tab6.medians}), are not consistent, but can be explained simply by the different frequency ranges used to measure $\\alpha$ for each sample: 408--2700\\,MHz for SMS4 and 178--750\\,MHz for 3CRR.  Given that the spectra of many radio sources turn over at low frequency, and that radio spectra often steepen at high frequency, the flatter median spectral index of the 3CRR sources can be explained by the lower frequency range over which $\\alpha$ was measured.  A comparison over more similar frequency ranges will be possible when 151.5\\,MHz data from Mauritius become available."
    },
    "0512/astro-ph0512541_arXiv.txt": {
        "abstract": "We present ground-based mid-infrared observations of Class~0 protostars in LDN 1448. Of the five known protostars in this cloud, we detected two, L1448N:A and L1448C, at 12.5, 17.9, 20.8, and 24.5~\\microns, and a third, L1448 IRS~2, at 24.5~\\microns.  We present high-resolution images of the detected sources, and photometry or upper limits for all five Class~0 sources in this cloud. With these data, we are able to augment existing spectral energy distributions (SEDs) for all five objects and place them on an evolutionary status diagram. ",
        "introduction": "\\label{sec:intro}  The majority of young stellar objects (YSOs) are classified using a scheme based on characteristics of their SEDs introduced by \\citet{lad84} \\citep[see also][]{lad87,lad91}.  They proposed three classes: I, II, and III; distinguished from one another by their near-infrared to mid-infrared spectral indices, which were presumed to be indicators of the evolutionary status of the object in question. Class I sources were initially thought to be representative of the earliest stage of pre-main-sequence evolution, {\\em i.e.} protostars.  The spectral index, $\\alpha$, is defined as  \\begin{equation} \\alpha = -\\frac{\\mbox{d}\\,\\log{\\nu F_{\\nu}}}{\\mbox{d}\\,\\log{\\nu}}; \\end{equation}  \\noindent for purposes of YSO classification, the region of the SED usually used for this calculation is from 2.2 $\\mu$m to 10 $\\mu$m. The spectral index is the negative of the slope of the SED, when graphed in the frequency domain \\citep{lad91}.  The most current version of this slope classification scheme is presented in \\citet{gre94}. For a Class I source, the SED tends to be broader than a single-temperature blackbody and $\\alpha > 0.3$, due to the fact that the peak of the SED of such a source is shifted towards the far-infrared because of the very low effective temperature of the YSO. These sources are usually deeply embedded in their nascent dust envelopes. For Class II and Class III YSOs, on the other hand, $ -0.3 > \\alpha \\ge -1.6$ and $\\alpha < -1.6$, respectively. The SED of a Class II will be much broader than that of a Class III, due to the presence of circumstellar dust with a fairly wide temperature distribution. \\citet{gre94} identify a fourth class of ``flat spectrum'' sources, for which $0.3 > \\alpha \\ge -0.3$, that have near-infrared spectra which are strongly veiled by continuum emission from hot circumstellar dust.  Observations of candidate protostars, in particular VLA 1623 in $\\rho$ Ophiuchi A, an extremely cold ($T_{e\\!f\\!\\!f} \\leq$ 20 K) object that does not appear at {\\em IRAS} wavelengths, led to the introduction of a new class of young stellar objects (YSOs) by \\citet{an93}.  They contended that Class I sources could not be the youngest examples of pre-main-sequence objects ({\\em i.e.}, not true protostars), and promptly created a new category for such objects, designated Class 0.  The search for these protostars has gained impetus in the last several years due to the development of highly sensitive submillimeter continuum detectors.  Class~0 protostars are the rarest of YSOs. Pre-main-sequence stars typically spend $\\ll 10^5$ years in this phase, on their way towards the onset of hydrogen fusion and stellar birth.  The current set of defining characteristics of a Class 0 protostar are as follows \\citep{an00}:  \\begin{itemize} \\item Detection of a radio continuum source, presence of a molecular outflow, or some other evidence of a central protostellar source; \\item Centrally peaked but extended submillimeter continuum emission indicating an envelope as opposed to a mere disk; \\item A ratio of L$_{submm}$/L$_{bol}$ $>$ 0.5\\% (where L$_{submm}$ is the luminosity radiated by the object at wavelengths longward of 350 $\\mu$m and L$_{bol}$ is the total luminosity of the object). \\end{itemize}  This last criterion indicates a circumstellar mass greater than the mass of the central protostellar core, as discussed in \\citet{an93}. There is no mention of the spectral index between 2 and 10 $\\mu$m, since these sources have generally not been detected shortward of 10 $\\mu$m using ground-based instruments. This will change, no doubt, as detector technology evolves.  Although reports of newly discovered Class~0 protostars have increased dramatically in the last few years, confirmed sources in this category still number well under a hundred. \\citep{an00,fro05}. The spectral energy distribution (SED) of a Class~0 object, with a typical effective temperature of $\\leq$40~K, peaks at around 100~\\microns; until fairly recently, essentially the only tool with which one may study the Wien side of the SEDs of protostars has been the \\emph{IRAS} database.  The majority of known Class~0 sources were detected by \\emph{IRAS} at 25~\\microns, if not at 12~\\microns; however, due to the nature of the \\emph{IRAS} data, the effective beam size at mid-infrared wavelengths is never less than $\\sim$30--40\\arcsec, even when using the high-resolution data-processing program YORIC \\citep{aum90,hur96,bar98,oli99}. Figure \\ref{fig:hires} provides an example of coadded 60-\\micron\\ \\emph{IRAS} data before (contours) and after HIRES processing (greyscale) via YORIC. The poor spatial resolution of the \\emph{IRAS} data makes associating point sources in the catalog with data from other wavelengths difficult and, in cases of close binary or multiple systems, does not allow the contributions of the individual components to the total flux to be distinguished. HIRES point-source modeling, which makes use of a little-known feature of the YORIC program, has been used successfully to estimate individual \\emph{IRAS} fluxes from sources as close as 20\\arcsec\\ apart, but this still does not compare with the sub-arcsecond resolution obtainable with millimeter-wave interferometry \\citep{bar98}. Consequently, the mid-infrared spectral energy distributions of Class 0 sources have tended to be poorly constrained. Data from the latest NASA Great Observatory, the \\emph{Spitzer Space Telescope} are already contributing significantly to our understanding of the mid-IR characteristics of young stellar objects, but \\emph{Spitzer} lacks the resolution to augment the information on the small-scale structure of their circumstellar environments which may be obtained from submillimeter and millimeter wavelength observations \\citep{ch00, lo00}.  Some of the most important questions in star formation studies today center around the multiplicity of these young systems, and how the binary properties evolve with time \\citep{lo00,rei00}. The majority of nearby young T Tauri stars are known to be binary systems, yet only a very few Class~0 protobinary systems have been resolved. It has been postulated that the giant Herbig-Haro flows driven by many young stellar objects may be initiated by the dynamical decay of unstable multiple systems \\citep{rei00}. Understanding star formation in the context of multiple systems has important implications for understanding the IMF and its physical meaning.  In order to address the issue of multiplicity and obtain details about the immediate circumstellar environments of Class~0 sources, it is necessary to study these objects at as many wavelengths as possible with high angular resolution.  A particularly rich cluster of Class~0 sources is located in the compact globule L1448, a molecular cloud which is part of the Perseus molecular complex. L1448 lies at a distance of $\\sim$300~pc; its mass, estimated from ammonia data, is $\\sim$50~M$_{\\sun}$ \\citep{ba86a}. Three embedded infrared sources were detected by \\emph{IRAS} in this cloud; since then, it has been surveyed at several wavelengths and mapped in a number of molecular lines \\citep{ba86b,ang89,bal97,ei00,wo00a}. Listed from west to east, the \\emph{IRAS} Point Source Catalog designations for the three sources are IRAS 03220+3035, IRAS 03222+3034, and IRAS 03225+3034. In the literature, these objects are more commonly known as L1448 IRS~1, L1448 IRS~2, and L1448 IRS~3, respectively (see Figure \\ref{fig:hires}).  Only the western-most of these objects (IRS~1) has been detected at near-infrared and shorter wavelengths \\citep{co79,ne84}; this source has been identified as a Class I YSO and hence was not included in our survey \\citep{ei00,bar00}. The other two \\emph{IRAS} point sources have been found to contain multiple Class~0 objects. IRS~2 was recently reported by \\citet{wo00a} to be a candidate protobinary system, based on the two distinct molecular outflows which emerge from that position, although the binary components were unresolved at that time. Subsequently \\citet{vo02} announced that both components of the IRS~2 binary had in fact been detected at BIMA. IRS~3 consists of L1448N:A,B (a 7\\arcsec\\ protobinary), L1448NW, and L1448C (note: L1448C is frequently referred to as L1448-mm), all of which drive molecular outflows as well \\citep{ter97,bar98,ei00,wo00a}. The locations of all of these sources, from previously published high-resolution submillimeter or millimeter-wave observations, are indicated by crosses in Figure \\ref{fig:hires} \\citep[and references therein]{oli99,lo00,an00}.  We observed all of the known L1448 Class~0 protostars in the mid-infrared, and present herein images of the detected sources plus multi-wavelength photometry and/or upper limits for all five Class~0 sources. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Spectral Energy Distributions - Individual Sources}  \\subsubsection{L1448C}  This Class~0 source, which drives a powerful, highly-collimated outflow, is a favorite target of many observers \\citep[see][and references therein]{bar98}. When first discovered, it was believed to be one of the youngest Class~0 sources found, with a kinematic age (based on the outflow parameters) of $\\sim$3500 years \\citep{bac90}. This figure was revised upwards to $\\sim$32,000 years by \\citet{bar98}, who found that the outflow was far more extensive than previously thought \\citep[see also][]{wo00a}.  The SED for L1448C is presented in Figure \\ref{fig:sedc}, with a two-greybody fit to the data shown by the solid line. The dashed line indicates a single-temperature greybody fit to the data.  \\subsubsection{L1448N:A,B}\\label{sec:nab}  In order to analyze the SEDs for L1448N:A and N:B, and use them to derive the parameters listed in Table 2, it is necessary to make certain assumptions regarding the individual fluxes of these two objects at far-infrared wavelengths. There are no high-angular-resolution data that can separate the 7\\arcsec\\ protobinary components at or near the peak of the dust emission ($\\sim$100~\\microns), and this causes problems when attempting to fit modified blackbodies to the spectral energy distributions of these objects.  Due to the excellent correspondence between the 12 and 25~\\micron\\ {\\it IRAS} HIRES point-source model fluxes from \\citet{bar98} for L1448N (which represent the sums of the fluxes for both components of the protobinary) and our MIRLIN fluxes (see Figure \\ref{fig:sednab}), it seems clear that the mid-infrared fluxes are due chiefly to L1448N:A, with very little, if any, contribution from N:B or NW.  Again, these results agree with those presented in \\citet{cia03}. In the millimeter, the situation is quite different: N:B produces most of the flux seen at 2.7 mm \\citep{ter97,lo00}. Using HIRES modeling, \\citet{bar98} showed it is possible to separate the {\\it IRAS} fluxes produced by L1448NW from the total flux emitted by the L1448N triple system. The L1448N system produces a total of 89~Jy at 100~\\microns, while the HIRES-modeled 100~\\micron\\ flux from NW is $\\sim$23~Jy. Therefore the {\\it maximum} possible flux from the protobinary at 100~\\microns\\ must be on the order of 66~Jy. Since some of the emission detected by \\emph{IRAS} at that wavelength is presumably due to shock-heated dust caused by the impact of the L1448C outflow on the dense core containing the L1448N protobinary, the true 100~\\micron\\ flux from the binary is probably somewhat less than 66~Jy \\citep{bar98,cur99}.  The flux ratio is approximately 5:1 in favor of N:B in the millimeter regime \\citep{lo00}, therefore it seems reasonable to explore a range of possibilities for the 100~\\micron\\ fluxes.  We have done so, in the following fashion: (1) Assume that the fluxes at 100~\\microns\\ have the same ratio as the 2.7 mm fluxes, so that N:A produces 11~Jy and N:B emits the remaining 55~Jy; (2) assume the inverse scenario (\\emph{i.e.}, N:B has a flux of 11~Jy, N:A produces the other 55~Jy).  These two possibilities define the endpoints of the range that we consider in this paper. A third instructive case to consider is the intermediate scenario which involves dividing the 66~Jy into equal portions; this assumes fluxes of $\\sim$33~Jy at 100~\\micron\\ from each of the protobinary components.  The assumptions inherent in this range of possibilities were used to generate a series of different SEDs for N:A and N:B; each of the SEDs for N:A were then fit with two-greybody models, while those for N:B were done with single-temperature greybodies. Sanity checks were performed for each case by using the various fits to calculate ranges of L$_{bol}$ for the two sources, and thus to ascertain whether or not the sum of their luminosities approximates the total luminosity expected for the protobinary. The luminosity of the protobinary was derived by taking the total fluxes of the L1448N system reported in the literature and subtracting the flux data for L1448NW at all wavelengths, then plotting the resultant SED and fitting with the sum of two modified blackbodies; this yields a maximum expected L$_{bol}$ of $\\sim$14.1~L$_{\\sun}$ for the protobinary system. The derivation of the L1448NW luminosity, using a single greybody fit to the data, is discussed in \\S\\ref{sec:nw}.  From the analysis described above, we find the plausible range of L$_{bol}$ for L1448N:A is 3.1--6.6~L$_{\\sun}$, and for N:B is 2.8--7.4~L$_{\\sun}$.  Since N:B dominates in the millimeter, and N:A in the mid-infrared, and the two objects have bolometric luminosities of the same order of magnitude, their SEDs must cross each other somewhere in the far-infrared, probably in the general vicinity of 100~\\microns. The scenario of approximately equal 100 \\micron\\ fluxes thus seems the most illustrative; it yields bolometric luminosities of $\\sim$4.8~L$_{\\sun}$ for N:A, and 5.5~L$_{\\sun}$ for N:B.  The sum is indeed less than the maximum expected luminosity.  Figure \\ref{fig:sednab} is a plot of the literature and MIRLIN flux data for each protobinary component along with a two-greybody fit for L1448N:A (dashed line) and a single-temperature greybody fit for N:B (dotted line) using the ``equal 100 \\micron\\ fluxes'' scenario. The previously published data points and derived fluxes for the entire L1448N protobinary system are plotted, and the sum of the fits to the individual SEDs (i.e., dashed + dotted) is shown as a solid line. The values of the physical parameters derived from the assumption of approximately equal 100 \\micron\\ fluxes (case 3) are reported in Table 2.  \\subsubsection{L1448NW}\\label{sec:nw}  L1448NW was not detected by MIRLIN at any wavelength, as may be seen in Figure \\ref{fig:imn}. This was the expected result, due to the extensive HIRES point-source modeling work reported in \\citet{bar98}, which established \\emph{IRAS} upper limits for this low-luminosity source of 0.015 and 0.05~Jy at 12 and 25~\\microns, respectively. Our MIRLIN upper limits do not improve on these results (see Table 1).  We have incorporated the submillimeter and millimeter-wave flux data from \\citet{ch00} and \\citet{lo00} in the plot shown in Figure \\ref{fig:sednw}; these data are entirely consistent with the FIR \\emph{IRAS} fluxes derived from HIRES modeling \\citep{bar98}. The results from a single-temperature greybody fit to the data at wavelengths longer than 60~\\microns\\ are used to calculate the parameters reported in Table 2, which allow us to place L1448NW on the evolutionary diagram in Figure \\ref{fig:evo}.  \\subsubsection{L1448 IRS~2}\\label{sec:irs2}  L1448 IRS~2 was identified as a Class~0 source by \\citet{oli99}, and first proposed as a candidate protobinary in \\citet{wo00a}, based on CO observations which mapped two distinct molecular outflows driven by IRS~2. \\citet{vo02} claimed detection of both binary components using BIMA, although the spatial separation was not reported. With the high spatial resolution of MIRLIN at the IRTF ($<$2\\arcsec\\ at all wavelengths), we had hoped to be able to detect and resolve both components of the system, but only a single source was found, at the 4$\\sigma$ level in the 24.5~\\micron\\ filter (see Figure \\ref{fig:irs2}). This suggests that either the binary components are too close together for MIRLIN to resolve (see the first paragraph of \\S\\ref{sec:res} for MIRLIN resolutions), or that the fluxes from one of the components are below the detection thresholds at all MIRLIN wavebands.  Figure \\ref{fig:sedi2} presents the SED for IRS~2, with a two-greybody fit to the data shown by the solid line. The dashed line indicates a single-temperature greybody fit to the data.   \\subsection{Evolutionary Status of L1448 Sources}  \\subsubsection{M$_{env}$ vs. L$_{bol}$}  Figure \\ref{fig:evo} is a plot of M$_{env}$ vs. L$_{bol}$ for all of the sources in this paper (filled squares), along with all other Class 0 (open diamonds) and Class I YSOs (filled triangles) with L$_{bol} \\> 1.0$ from Table 1 in \\citet{bon96} plotted for comparison.  This plot may be used as a diagnostic tool for determining the relative ages of embedded sources \\citep{an94,sar96}. The diagram shows infall envelope mass (along the vertical axis) vs. bolometric luminosity (horizontal axis). L$_{bol}$ is directly related to the mass of the central protostellar core at very early stages of evolution, assuming all luminosity is generated by gravitational infall, according to:  \\begin{equation} L_{bol}\\ =\\ {G{\\dot M} M_* \\over R_*}, \\end{equation}  \\noindent where a protostellar radius of $\\sim$3~R$_{\\sun}$ is assumed \\citep{sta80b}.  Our results depicted in Figure \\ref{fig:evo} for L1448C and L1448NW are consistent (within errors) with those presented in \\citet{bar98} for those two sources. The new results in the diagram presented here are (1) L1448 IRS~2 is plotted on the diagram to allow comparison with the other sources from this cloud, and (2) we have plotted the L1448N protobinary components {\\it separately} on this graph, using the assumptions discussed in \\S\\ref{sec:nab}.  All of the target objects discussed in this paper lie within the Class~0 region of the evolutionary diagram in Figure \\ref{fig:evo}, although L1448N:A appears to be close to the transition zone between Class~0 and Class I, indicating N:A may be the most evolved source in this cloud after the Class I source L1448 IRS 1.  By the same criteria, L1448N:B is very likely the youngest source that we attempted to image in this cloud.  These classifications are reflected in the values of the parameter L$_{submm}$/L$_{bol}$, listed in Table 2, which range from 1.7--7.3\\%: the value of 1.7\\% corresponds to L1448N:A, and 7.3\\% to N:B (for the case 3 scenario of equal 100 \\micron\\ fluxes).  We note that recently \\citet{cia03} have argued that L1448N:A is probably a Class I object. However, our smallest value of L$_{submm}$/L$_{bol}$ for N:A, derived for case 2 (assuming 100 \\micron\\ flux for N:A $\\sim 55$ Jy, for N:B $\\sim 11$ Jy), was 1.5\\%, which lies well above the lower-limit value of 0.5\\% established for Class 0 objects by \\citet{an00}.  We point out two additional factors which bolster the argument for the Class 0 status of N:A: (1) Its molecular outflow is extremely well-collimated, similar to outflows associated with other Class 0 sources \\citep{wo00a}. (2) We find that it is not detected in the near-infrared by 2MASS, although the Class I source L1448 IRS 1 is strongly detected in all 2MASS bands (JHK$_s$). Although fluxes are presented in all bands for a ``source'' detected by 2MASS within a few arcseconds of L1448N:A, the J-band flux is presented as an upper limit, and the H-band and K-band fluxes suffer photometric confusion. The 2MASS Quicklook images reveal a bow-shaped knot north of L1448N:A in all three bands, and an additional jet-like feature south of L1448N:A in the K-band image. Inspection of these features reveals that they correspond to H$_2$ outflow features imaged by \\citet{ei00}. The jet-like feature in the K-band image of L1448N:A is part of the blueshifted outflow lobe associated with L1448C, which is located $\\sim 1.5^{\\prime}$ to the southeast of the L1448N sources. This outflow was also identified in HIRES-processed {\\it IRAS} images and high-velocity CO emission \\citep{bar98,wo00a} The bow-shaped knot that appears just north of L1448N:A in the J,H, \\& K$_s$ 2MASS images is very prominent in the H$_2$ image. There is no trace of a near-infrared point source counterpart to L1448N:A in the 2MASS data.  Additional near-infrared data for this source became available during a late revision of this paper: a FLAMINGOS survey of L1448 down to $m_{K_s} \\sim$ 17 mag \\citep{tsu05}.  Again, no near-IR counterpart to L1448N:A was found during this survey.  However, these same investigators also used {\\em Chandra} to observe L1448; they detected a weak X-ray source in the general vicinity of N:A.  No X-ray emission was detected from any of the other Class 0 sources under discussion in this paper.  \\citet{tsu05} discuss the fact that, given the non-detection in the near-infrared bands, L1448N:A must have an extremely steep near-to-mid-infrared spectral index ($\\alpha>3.2$), and that nearly all known Class I protostars have $\\alpha \\leq 2$ (the single exception being the peculiar young stellar object, WL 22, with $\\alpha = 3$).  We therefore contend that all currently available data for L1448N:A indicate that it is a bonafide (albeit borderline) Class 0 source; the weak X-ray detection does not rule this out, and all other data support this conclusion.  \\subsubsection{Effects of the L1448C Outflow on the L1448N/NW Core}\\label{sec:out}  Using ammonia data, \\citet{cur99} report that the high-velocity L1448C outflow has had a significant affect on the northern core, creating a deep ``dent'' in its southern edge, directly south of the protobinary. \\citet{bar98} suggest that the impact of the outflow from L1448C on the dense core fragmented it and actually induced the formation of the L1448N/NW system.  At first glance, this scenario of outflow-induced fragmentation appears to be at odds with our findings above that N:A is the most evolved Class 0 source in this cloud, while N:B is the youngest and most deeply embedded.  Of course, if these two sources do not form a gravitationally-bound protobinary system, it is unnecessary for them to be coeval. But for the purposes of this discussion, we will assume with \\citet{ter97}, \\citet{bar98}, \\citet{ei00}, and \\citet{wo00a}, that N:A and N:B are the bound components of a protobinary. Is it possible for presumably coeval sources, which appear to be physically similar, to travel different evolutionary paths?  One recently-proposed explanation for such a dichotemy in evolutionary status of protobinary components is found in the disintegration of unstable triple or higher-order multiple systems; such events may cause disk truncation for some of the components, subsequent episodic outflow activity linked to giant Herbig-Haro flows, and accelerated evolution of one or more components due to the rapid dissipation of circumstellar envelopes \\citep{rei00, rei01}. It is true that most of the young sources in L1448 have been identified as driving sources of parsec-scale Herbig-Haro flows \\citep{bal97,bar98,oli99,ei00,wo00a}. However, at $\\sim$2000~AU separation, with a rotational period on the order of 60,000 years \\citep{ter97}, the L1448N:A,B system does not seem to fit into this particular theoretical paradigm, since a remnant binary from such a former higher-multiplicity system is predicted to have a much smaller separation. The L1448N:A,B components appear to be too young to have had time to interact with each other or with L1448NW in any significant way.  Furthermore, according to the \\citet{rei00} theory, binary components that are remnants of unstable multiples having undergone ejection events might be expected to have more pronounced differences between their apparent evolutionary stages; \\emph{e.g.}, one component may transition abruptly to Class II status, while the other remains in Class 0 or Class I.  All of the L1448 sources under discussion in this paper still fall within the same evolutionary category, although N:A appears to be close to transitioning to Class I.  We propose an alternative explanation for the (admittedly minor) difference between the protobinary components: this could be nothing more than a chance side effect of the powerful L1448C outflow. Lacking a good understanding of the three-dimensional relative geometry of the sources in this cloud, we are unable to verify our hypothesis, but one possible scenario is that the outflow from L1448C has partially stripped the envelope surrounding N:A, while leaving N:B deeply embedded. This would allow N:A and N:B to be coeval sources that formed due to the impact of the L1448C outflow, while explaining the different \\emph{apparent} ages of these objects."
    },
    "0512/astro-ph0512631_arXiv.txt": {
        "abstract": "I consider constraints from observations on a cutoff scale in clustering due to free streaming of the dark matter in a warm dark matter cosmological model with a cosmological constant.  The limits are derived in the framework of a sterile neutrino warm dark matter universe, but can be applied to gravitinos and other models with small scale suppression in the linear matter power spectrum.  With freedom in all cosmological parameters including the free streaming scale of the sterile neutrino dark matter, limits are derived using observations of the fluctuations in the cosmic microwave background, the 3D clustering of galaxies and 1D clustering of gas in the Lyman-alpha (Ly$\\alpha$) forest in the Sloan Digital Sky Survey (SDSS), as well as the Ly$\\alpha$ forest in high-resolution spectroscopic observations.  In the most conservative case, using only the SDSS main-galaxy 3D power-spectrum shape, the limit is $m_s > 0.11\\rm\\ keV$; including the SDSS Ly$\\alpha$ forest, this limit improves to $m_s > 1.7\\rm\\ keV$.  More stringent constraints may be placed from the inferred matter power spectrum from high-resolution Ly$\\alpha$ forest observations, which has significant systematic uncertainties; in this case, the limit improves to $m_s > 3.0\\rm\\ keV$ (all at $95\\%$ CL). ",
        "introduction": "The scale to which the ansatz of a (nearly) scale-invariant power spectrum describes the cosmological primordial density fluctuations is unquantified at the smallest scales.  The standard paradigm of an unmodified, nearly scale-invariant power spectrum arising from the amplification of quantum vacuum fluctuations of an inflaton field agrees well with the inferred matter power spectrum in the linear to mildly nonlinear regime of cosmological structure formation \\cite{Spergel:2003cb,Tegmark:2003ud,Seljak:2004xh,Viel:2004np,Abazajian:2005dt}. Models of the nonlinear clustering of galaxies are also consistent with a nearly scale-invariant perturbation spectrum \\cite{Abazajian:2004tn}.    There are a number of physical properties of the dark matter that may alter its perturbation spectrum at small scales.  The hot big bang primordial plasma in the early universe may partially or completely thermodynamically couple to the dark matter.  Any primordial velocity distribution that is imparted to the dark matter will allow dark matter to escape perturbations with sufficiently small gravitational potentials and damp structure formation at such scales.  The most cited cold dark matter (CDM) candidate, the lightest supersymmetric particle (LSP), has a small but non-zero velocity dispersion, which damps structures below Earth mass scales \\cite{Hofmann:2001bi,Green:2005fa}.  In warm dark matter (WDM) models, a light mass particle such as a sterile neutrino~\\cite{Dodelson:1993je} or gravitino~\\cite{Blumenthal:1982mv} with a large velocity dispersion can suppress linear structure formation up to galactic scales.  The dark matter power spectrum may also be suppressed on small scales due to the production of the dark matter LSP through decay of charged progenitor causing a suppression of structure below the horizon scale at decay~\\cite{Sigurdson:2003vy,Kaplinghat:2005sy} or inflation with broken scale invariance~\\cite{Kamionkowski:1999vp}.  Models with a suppression of small scale power have drawn attention due to their potential alleviation of several unresolved problems in galaxy and small scale structure formation.  These include, first, the reduction of satellite galaxy halos \\cite{Kauffmann:1993gv,Klypin:1999uc,Moore:1999wf,Willman:2004xc}, second, the reduction of galaxies in voids~\\cite{Peebles:2001nv,Bode:2000gq}, third, the low concentrations of dark matter in galaxies~\\cite{Dalcanton:2000hn,vandenBosch:2000rz,Zentner:2002xt,Abazajian:2005kz}, fourth, the angular momentum problem of galaxy formation~\\cite{Dolgov:2001nq}, and fifth, the formation of disk-dominated galaxies \\cite{Governato:2002cv,Kormendy2005}.  A significant reduction of power on sufficiently large scales will be in conflict with observations of clustering at small scales.  It was noted by Narayanan et al.~\\cite{Narayanan:2000tp} that observations of power on the smallest linear scales in the Lyman-alpha (Ly$\\alpha$) forest may provide particularly stringent constraints. In this paper, I examine the upper limits on the scale of such a suppression of small scale power from measures of matter clustering at small scales yet in the linear regime, where systematic effects in modeling can be minimized.  I frame the constraints in terms of a sterile neutrino WDM candidate, motivated by new data and an accurate calculation of the transfer function for sterile neutrino dark matter based on an accurate calculation of the nonthermal production of sterile neutrino dark matter~\\cite{abazajian05}.  Limits on gravitino WDM and general departures of matter power spectra at small scales also can be inferred from these results.  The small-scale clustering data employed here are the Sloan Digital Sky Survey (SDSS) main galaxy 3D power spectrum \\cite{Tegmark:2003uf}, the inferred linear matter power spectrum from observations of the flux power spectrum of Ly$\\alpha$ absorption in the SDSS quasar catalogue~\\cite{McDonald:2004eu,McDonald:2004xn}, and the inferred linear matter power spectrum from high-resolution observations of the Ly$\\alpha$ forest of Croft et al.~\\cite{Croft:2000hs} as reanalyzed and augmented by Viel et al.~\\cite{Viel:2004bf} (hereafter VHS).  In order to alleviate degeneracies with other cosmological parameters, particularly the optical depth to the cosmic microwave background (CMB) and primordial scalar perturbation amplitude, as well as tighten constraints within allowable possibilities, I include observations of the fluctuations of the CMB from the first year Wilkinson Microwave Anisotropy Probe (WMAP)~\\cite{hinshaw03,kogut03}, the Cosmic Background Imager (CBI)~\\cite{readhead04}, Boomerang~\\cite{Jones:2005yb}, the Arcminute Cosmology Bolometer Array (ACBAR)~\\cite{Kuo:2002ua}, and the Very Small Array (VSA)~\\cite{dickinson04}.  Viel et al.~\\cite{Viel:2005qj} have done a similar statistical analysis to that presented here, using WMAP and the inferred linear matter power spectrum of VHS to place constraints on a gravitino WDM candidate, which is applied to sterile neutrino dark matter via a central-value relationship.  The work presented here is motivated by using a significantly more accurate transfer function for sterile neutrino dark matter of Ref.~\\cite{abazajian05}, as well as the results of inclusion of many more independent cosmological structure data.  It should be noted that observations of nonlinear structures may be a very powerful handle of primordial power at extremely small scales. The observations of anomalous flux ratios in gravitational lens systems can be an indication of substructure having a significant fraction of the mass of the lensing galaxies, which would not be present in certain WDM models~\\cite{dalal2002,kochanek2004}.  However, the exact level of the suppression of substructure that can be tolerated by these lensing observations is not clear, particularly given that some small mass halos may be formed by fragmentation in WDM models~\\cite{Knebe:2003hs}; moreover, there may be a significant enhancement of the observed anomaly by line-of-sight isolated halos~\\cite{Chen:2003uu}.  Another highly nonlinear process asserted to place strong constraints on the presence of small scale power is high-redshift (high-$z$) star formation inferred to be required to cause reionization early enough to produce the anomalously high $TE$ cross-correlation at low multipoles seen by WMAP~\\cite{Barkana2001,Yoshida:2003rm}.  CDM models themselves generally have difficulty in producing such high-$z$ reionization even when including the invocation of a very early generation of high-mass star formation~\\cite{Rozas:2005pt}, and it is not clear that the resolution of the problem of an anomalously high $TE$ cross-correlation at large scales observed by WMAP will be solved by the presence of small mass halos at high-$z$. ",
        "conclusions": "\\label{conclusions}  Potential problems in galaxy and small-scale structure formation indicate the possibility of a small-scale velocity damping of perturbations of the type in warm dark matter such a sterile neutrino dark matter.  This work presents a calculation of cosmological parameter constraints including the possibility of a warm dark matter candidate to investigate the largest scale of the suppression of small scale power allowed by linear measures of the cosmological structure from the inferred matter power spectrum.  I find the lower limits on a sterile neutrino dark matter particle mass placed by observations of fluctuations in the cosmic microwave background, clustering of galaxies and in the Ly$\\alpha$ forest in the Sloan Digital Sky Survey (SDSS), and clustering of the Ly$\\alpha$ forest in high-resolution spectroscopic observations by the Keck and VLT Observatories.  The lower mass bound in the most conservative case, using only the SDSS main-galaxy 3D power-spectrum shape is $m_s > 0.11\\rm\\ keV$; including the SDSS Ly$\\alpha$ forest, this limit improves to $m_s > 1.7\\rm\\ keV$, while also including high-resolution Ly$\\alpha$ forest data, this improves to $m_s > 3.0\\rm\\ keV$, with all limits at $95\\%$ CL. These results are consistent with the results of Viel et al.~\\cite{Viel:2005qj}, but more constrained than their results ($m_s > 2.0\\rm\\ keV$) due to the inclusion here of additional CMB and large scale structure data. The limits from the high-resolution Ly$\\alpha$ forest data are subject to large systematic uncertainties, and therefore I adopt the SDSS Ly$\\alpha$ data as the current WDM structure benchmark.  These constraints can also be inferred to place limits on gravitino warm dark matter candidates as well as other models with a primordial fluctuation small scale cutoff, $\\alpha$. The central values and confidence intervals for the other cosmological parameters do not change significantly with the inclusion of the possibility of WDM.  The sterile neutrino WDM candidate is well constrained by the linear clustering observations considered here, and by limits on its radiative decay by observations of the Virgo cluster by XMM-Newton, which provides an upper mass constraint~\\cite{Abazajian:2001vt,Boehringer2001}.  There also exist upper mass constraints from the diffuse X-ray background~\\cite{Mapelli:2005hq,Boyarsky:2005us}.  The combination of these constraints allows a narrow range for this particle dark matter candidate mass: \\begin{equation} 1.7{\\rm\\ keV} < m_s < 8.2\\rm\\ keV. \\end{equation} It remains for further work to detect or constrain X-ray lines from nearby clusters of galaxies, as well as use the presence or lack of a cutoff scale in the linear matter power spectrum to detect or exclude this dark matter candidate."
    },
    "0512/astro-ph0512407_arXiv.txt": {
        "abstract": "We present the results of time-resolved spectroscopy of 13 O-type stars in the Cas OB6 stellar association.  We conducted a survey for radial velocity variability in search of binary systems, which are expected to be plentiful in young OB associations. Here we report the discovery of two new single-lined binaries, and we present new orbital elements for three double-lined binaries (including one in the multiple star system HD~17505). One of the double-lined systems is the eclipsing binary system DN~Cas, and we present a preliminary light curve analysis that yields the system inclination, masses, and radii. We compare the spectra of the single stars and the individual components of the binary stars with model synthetic spectra to estimate the stellar effective temperatures, gravities, and projected rotational velocities. We also make fits of the spectral energy distributions to derive $E(B-V)$, $R=A_V/E(B-V)$, and angular diameter. A distance of 1.9~kpc yields radii that are consistent with evolutionary models.  We find that 7 of 14 systems with spectroscopic data are probable binaries, consistent with the high binary frequency found for other massive stars in clusters and associations. ",
        "introduction": "%  The frequency of spectroscopic binaries (SBs) in star forming regions has been a topic of interest for many years.  \\citet{mas98} review the fraction of binary systems found among O stars in clusters/associations, the field, and runaway systems.  The results from their large survey agree with previous studies \\citep*[e.g.][]{gar80,lev91} and indicate that a large fraction of O stars in clusters and associations are binaries while fewer field stars show multiplicity and very few runaway stars have companions.  The values from \\citet{mas98} are broken down into spectroscopic and resolved binaries.  For O stars in clusters and associations they give a total of 75\\% as binaries; 61\\% of the binaries are spectroscopic and 41\\% are resolved (with some obvious overlap).  The spectroscopic binary group includes, however, a significant number of cases labeled as SB1? or SB2? (systems suspected of radial velocity variability or showing double lines), so the fraction of confirmed spectroscopic binaries is then 34\\% among O stars in clusters and associations.  We have conducted a radial velocity survey of O-type stars in the Cas OB6 stellar association in search of spectroscopic binary systems. The Cas~OB6 region is one of the largest complexes in the Perseus arm of the Galaxy and is located at a distance of about 2.3~kpc \\citep*{gar92,mas95,fos03}. The warm dust and ionized gas in this region extends over 150~pc along the arm and is dominated by the W3/W4/W5 chain of \\ion{H}{2} regions and the HB3 supernova remnant \\citep*{car00,ter03}. The open cluster IC~1805 (OCL~352 = Melotte 15) is situated at the western end in the W4 region and contains 8 O-type stars \\citep{mas95}. The three brightest of these, HD~15558, 15570, and 15629, are probably the ionization source that powers the blowout of gas in the Galactic chimney \\citep*{nor96,ter03} and the emission arc above W4 \\citep*{rey01}. The open cluster IC~1848 (OCL~364) is found in the eastern part of the association in the W5 region. Its primary source of ionizing flux is the multiple O-star system HD~17505 \\citep{sti01}. In addition to these two clusters, \\citet{car00} find evidence of 19 other young clusters that are still embedded in their primordial clouds.  Our knowledge of the binary star population among the young, massive stars of Cas~OB6 is quite limited at present. There are a number of radial velocity investigations of the sample but the overall small number of measurements has made the detection of binary stars difficult \\citep*{hay32,und67,ish70,con77,liu89,liu91,rau04}.  Spectroscopic orbits are available for only three systems, HD~15558 \\citep{gar81}, HD~16429 \\citep{mcs03}, and BD$+60^\\circ497$ \\citep{rau04}, while a fourth star, BD$+60^\\circ470$ = DN~Cas, is a known eclipsing binary \\citep{fra74}. However, results for O-stars in other clusters and associations suggest that many more binaries may exist \\citep{mas98,gar01}. It is important to determine the extent of the binary star population in Cas~OB6 in order to estimate accurately the ionizing flux of the key O-type stars and to assess the role of gravitational encounters with binary stars that might lead to ejection of runaway stars \\citep*{gie86,hoo01}.  Here we present a reconnaissance of the radial velocity variations observed in 13 bright stars in Cas~OB6 that we obtained over a 5 night run with the KPNO 4~m Mayall telescope in 2003 October.  We also present the first radial velocity study of the eclipsing binary system DN~Cas, which we combine with published photometry to determine the component masses. We discuss our observations and reduction techniques in \\S2. The radial velocity measurements and the derivation of orbital elements are summarized in \\S3. We use a tomography algorithm to separate the spectra of the binary star components, and we describe a comparison of the stellar spectra with model spectra in \\S4 that we use to estimate the basic stellar parameters. We examine in \\S5 the spectral energy distributions to determine reddening, ratio of total-to-selective extinction, and angular size for each target, and we discuss the probable distance to Cas~OB6. Details about the individual objects follow in \\S6. Finally we summarize in \\S7 our derived fraction of multiple and single systems for Cas~OB6 in the context of past results. ",
        "conclusions": "%  We have monitored thirteen O stars in the Cas~OB6 association in search of multiple star systems.  Five of the thirteen were found to have radial velocity variations related to orbital motion.  Of these five, we have determined orbital solutions for three double-lined systems: the eclipsing binary DN~Cas, BD$+60^\\circ497$, and the spectroscopic triple HD~17505A.  The remaining two binaries, BD$+60^\\circ594$ and HD~17520, are single-lined binaries with periods longer than 5 days. In addition to these five, one of our targets, HD~15558, is probably a long period binary \\citep{gar81,sti01}. One other bright O-star in Cas~OB6 is HD~16429, which was found by \\citet{mcs03} to be a spectroscopic triple. If we include all of these objects in the known sample of observed O-stars in the Cas~OB6 association, then we find that 7 of the 14, or $50 \\%$, are spectroscopic binaries.  This value is well within the range found by \\citet{mas98} for spectroscopic binaries among stars in clusters and associations.  We note that two systems, HD~17505 and DN~Cas, are of special interest and would benefit from follow-up observations. We are planning to obtain new three-color photometry of the eclipsing binary DN~Cas in order to make a better fit of the light curve and to determine an accurate mass of the O8~V primary star. The multiple system HD~17505 is a fascinating system with at least four O stars (with a total mass close to $100~M_\\odot$) which are apparently gravitationally bound.  High resolution studies via interferometry to separate the Aa and Ab components would produce a more accurate kinematic model for the system and would help us better understand stellar interactions in these systems and their role in the formation and ejection of massive stars."
    },
    "0512/astro-ph0512473.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract {} % aims {RXS J113155.4-123155 (z=0.66) is a quadruply imaged lensed quasar with a resolved Einstein Ring. The goal of this paper is to provide a full characterization of this system, and more particularly accurate astrometry and photometry. These observational constraints constitute a mandatory ingredient for the precise determination of the lens mass profile, the derivation of the Hubble constant $H_0$ from time delay measurements and investigations on the presence of massive substructures in the lensing galaxy.} % Method {Visible and near-infrared imaging observations of RXS J113155.4-123155 were carried out at various epochs using several ground based telescopes and the HST. The frames have been deconvolved using the MCS algorithm. A Singular Isothermal Ellipsoid (SIE) + external shear has been used to model the lensing galaxy potential. } %Results { MCS deconvolution enables us to separate the flux of the QSO (point-like images) from that of its host galaxy and to accurately track the flux variations of the point-like images in various filters. The deconvolved frames unveil several multiply imaged structures in the Einstein ring and an unidentified object in the vicinity of the lensing galaxy.  We discuss the lightcurves and the chromatic flux ratio variations and deduce that both intrinsic variability and microlensing took place during a span longer than one year.  We demonstrate that microlensing may easily account for the so called anomalous flux ratios presented in the discovery paper. However, the observed flux ratios are still poorly reproduced when modeling the lens potential with a SIE+shear. We argue that this disagreement can hardly be explained by milli-lensing caused by substructures in the lensing galaxy. A solution proposed in Paper II consists in a more complex lens model including an octupole term to the lens gravitational potential.  } % conclusions heading (optional), leave it empty if necessary {}  \\titlerunning{Multi-wavelength and multi-epoch imaging of RXS J113155.4-123155.}  \\authorrunning{Sluse D. et al.} ",
        "introduction": "\\label{sec:intro}  RXS J113155.4-123155 (hereafter J1131) is one of the nearest confirmed multiply imaged AGN. The source at $z_s = 0.658$ is lensed by an elliptical galaxy at $z_l = 0.295$ (Sluse et al. \\cite{SLU03}; hereafter {\\it{Paper I}}). This system is a long axis quad with an image configuration very similar to B1422+231 (Patnaik et al. \\cite{PAT92}): three merging images (B-A-C; typical of a source lying close to a cusp caustic) face the faint saddle-point image D lying close to the lensing galaxy (G). For such a system (i.e. cusp configuration system), the magnification behaviour is well understood - the flux of the middle image should be equal to the total fluxes of the two outer images (Schneider \\& Weiss \\cite{SCH92a}). However, this generic prediction was strongly violated at the epoch of the discovery, suggesting the likely presence of small-scale structures in the lensing galaxy (Keeton, Gaudi \\& Petters ~\\cite{KEE03b}, ~\\cite{KEE05}; hereafter KGP03, KGP05). Thanks to the data set presented here, we re-examine the values of the flux-ratios with regard to their temporal and chromatic variations. Additionally, because this quad is quite bright and shows lensed images with a wide angular separation ($2 \\theta_E \\sim$ 3.6\\arcsec), it is potentially a target of interest for time delay measurements. Also, the bright Einstein ring detected from the optical to the near-infrared offers unique constraints on the lens modeling (Kochanek et al.~\\cite{KOC01}) that may break the degeneracy between the intrinsic lens galaxy shape and the external shear (Keeton et al. \\cite{KEE00}). However, the bright Einstein ring introduces an additional complexity to the photometry measurement because it is superimposed over the lensed point-like images. This motivates our extensive discussion of the sources of systematic photometric errors affecting the individual lensed QSO images and of the photometric accuracy that can be reached for J1131.  The observations presented here constitute the first part of a follow-up study of this system that also includes detailed lens models and source reconstruction from HST imaging, X-ray imaging as well as optical and NIR spectroscopy. Sects.~\\ref{sec:Obs} \\&~\\ref{sec:astrophot} describe the data reduction and analysis. Section~\\ref{sec:results} includes a description of the ring morphology, a simple modeling of the lensing galaxy and a discussion of the flux (time and chromatic) variability of the lensed images. Finally, Sect.~\\ref{sec:conclusions} summarizes the main results and presents our conclusions.  We adopt throughout this paper $H_0=$ 65~km s$^{-1}$ Mpc$^{-1}$, $\\Omega_0=$ 0.3 and $\\Lambda_0=$ 0.7; magnitudes are computed in the Vega system.  %__________________________________________________________________ ",
        "conclusions": "\\label{sec:conclusions}  We have presented direct optical and NIR imaging of the quadruply imaged quasar RXS J113155.4-123155 obtained with ground-based telescopes and with HST at various epochs. The HST NICMOS images unveil many details in the Einstein ring of this system as well as a putative companion object (X) lying at $\\sim$1\\arcsec~from the lensing galaxy. Accurate positions have been derived for the 4 lensed images, for X and for the lensing galaxy based on the HST NIC2 data and on the ground-based VLT observations. Both sets of measurements do agree within 6 mas. A SIE+$\\gamma$ model has been fitted to the observed image configuration. Nevertheless, the positions of image D and of the lens galaxy do not match satisfactorily the observed ones. Additionally, we have used the SIE+$\\gamma$ model to derive the expected point-like flux ratios when taking into account the extended nature of the source.  The images have been deconvolved using the MCS method known to preserve photometry. Hence, we derived for the first time relative and absolute photometry of J1131 from the optical to the NIR range. Since the photometric measurements of the point-like images in J1131 are complicated by the superimposed Einstein ring, we have thoroughly discussed the potential sources of photometric errors. We have shown that the MCS method allows to reach a photometric accuracy close to the photon noise limit in the $B$ and $V$ filters if a good quality PSF can be constructed. On the other hand, small systematic errors affect the flux ratios in the $R$ band due to the Einstein ring.  The small chromatic variations of the flux ratios between the $B$ and $R$ bands suggest a low level of differential extinction between the lensed images of J1131. Based on sparse photometric measurements obtained between November 2002 and April 2004, we have shown the likely presence of microlensing for image A together with intrinsic quasar flux variability up to 0.3 mag. Microlensing of A is also supported by our Chandra observations of J1131 in the X-rays (see Paper III, Claeskens et al., in preparation; Blackburne et al.~\\cite{BLA05}). Additionally, based on the chromatic variations of the flux ratios, we have shown that A should likely be de-amplified (scenario S1), in agreement with the fact that it is a saddle-point of the arrival-time surface{\\footnote{During the referee process, Rozo et al. (\\cite{ROZ05}) demonstrated that small scale structures, if they can be considered as linear perturbations, might not preferentially demagnify saddle point images. However, large perturbations from point-mass microlenses are likely caused by non linear perturbers and the results of Rozo et al. (\\cite{ROZ05}) may not be valid in such a case.}} (Schechter \\& Wambsganss~\\cite{SCH02},  Keeton~\\cite{KEE03a}). Alternatively, our data cannot rule out that two independent microlensing amplifications (scenario S2) with roughly the same amplitude occur for both images B and C.  We have used the $R_{\\rm{cusp}}$ relation (Mao \\& Schneider, \\cite{MAO98}; Eq.(~\\ref{equ:rcusp})) to investigate the evidence for substructures in the lensing galaxy. Once lower/upper bounds on the flux ratios are imposed from our knowledge of microlensing occuring in this system, we find a lower limit on $R_{\\rm{cusp}}$ close to zero. This shows that the small scale structure(s) identified in KGP03 from the $R_{\\rm{cusp}}$ relation, was (were) indeed due to (a) microlens(es).   Nevertheless, the observed flux ratios do not agree with the model predictions. We argue that milli-lensing effects cannot provide the best explanation under S1 while under S2, observations suggest that a massive substructure should {\\it magnify} the saddle point image A. This is however a low probability scenario (Schechter \\& Wambsganss~\\cite{SCH02}). An alternative to the milli-lensing explanation is to introduce in the macro-lens potential higher order multipole terms (e.g. Evans \\& Witt \\cite{EVA03}). This is described in Paper II, where we successfully added an octupole ($m=4$) to the fiducial model and reproduced both the observed flux ratio $I_{\\rm B}/I_{\\rm C}$ (under S1) and the observed relative lens position. Therefore, we found  that  the conjunction of microlensing and of higher order terms in the lens potential could explain the flux ratios observed in J1131. The preference of a more complex macro model over the presence of substructure is opposite to what Kochanek \\& Dalal (\\cite{KOC04}) have found for other lenses. Of course the solution of adding an octupole might not be unique and other models such as those including several sources of shear (e.g. Morgan et al.~\\cite{MOR05}), or with two mass components (e.g. Dye \\& Warren~\\cite{DYE05}) should be explored.    %                                     Two column figure (place early!) %______________________________________________ Gamma_1 (lg rho, lg e)"
    },
    "0512/astro-ph0512127_arXiv.txt": {
        "abstract": "In the six billion years between redshifts $z=1$ and $z=0.1$, galaxies change due to the aging of their stellar populations, the formation of new stars, and mergers with other galaxies. Here I explore the relative importance of these various effects, finding that while mergers are likely to be important for the red galaxy sequence they are unlikely to affect more than 10\\% of the blue galaxy sequence.  I compare the galaxy population at redshift $z \\sim 0.1$ from the Sloan Digital Sky Survey (SDSS) to the sample at $z \\sim 1$ from the Deep Extragalactic Evolutionary Probe 2 (DEEP2).  Galaxies are bluer at $z\\sim 1$: the blue sequence by about 0.3 mag and the red sequence by about 0.1 mag, in \\umg. I evaluate the change in color and in the luminosity functions of the two sequences using some simplistic stellar population synthesis models. These models indicate that the luminous end of the red sequence fades less than passive evolution allows by about 0.2 mag. Due to a lack of luminous blue progenitors, ``dry'' mergers betweeen red galaxies then must create the luminous red population at $z\\sim 0.1$ if stellar population models are correct. The blue sequence colors and luminosity function are consistent with a reduction in the star-formation rate since $z\\sim 1$ by a factor of about three, with no change in the number density to within 10\\%. These results restrict the number of blue galaxies that can fall onto the red sequence by any process, and in particular suggest that if mergers are catastrophic events they must be rare for blue galaxies. ",
        "introduction": "\\label{motivation}  How galaxies form is naturally an important and interesting topic for astronomers, living as we do in the disk of a relatively large and typical galaxy. A well-developed theory of galaxy formation exists, in which galaxies form at the centers of the potential wells of the dark matter halos that are a natural prediction of the current standard Cold Dark Matter (CDM) cosmology. Gasdynamic simulations (e.g. \\citealt{kerev05a,springel05b,nagamine05a}) and simpler ``semi-analytic'' models (e.g. \\citealt{benson03a,nagashima05a}) attempt to make predictions of this theory.  The class of predictions I address here involve the change of galaxies (and consequently of the galaxy population) over time. Of course, given the finite age of the Universe, galaxies must change over time, but the predictions are more specific than that and qualitatively point mostly in the same direction: that star formation was higher in galaxies in the past and that their luminosities were larger (\\citealt{springel03a, nagamine01a}).  High redshift galaxy surveys of the past ten or twenty years or so have revealed such changes --- indeed, the observations predate the theoretical predictions. The first such observations were in clusters, and indicated that the fraction of blue galaxies has declined markedly in the last 3--4 Gyrs (\\citealt{butcher84a}). Observations of galaxies in the field have indicated that this process occurs in the field as well (\\citealt{lilly95b, cowie96a, cohen02a, im02a, gabasch04a, bell04a, faber05a, willmer05a}). Taken as a whole, these observations have indicated that between redshifts $z=1$ and $z=0$ galaxies became dimmer and less star-forming.  Here, I present a detailed comparison of the change in the luminosity and color distribution of galaxies between redshifts around $z\\sim 1$ and $z\\sim 0.1$.  At redshift $z\\sim 0.1$ I use the Sloan Digital Sky Survey (SDSS; \\citealt{york00a}) to measure the color and absolute magnitude distribution of galaxies. As \\citet{blanton03d} demonstrated, the structure of this distribution consists of the well-known red sequence of galaxies (\\citealt{baum59, faber73a, visvanathan77, terlevich01a}) and a similar but much broader blue sequence of galaxies. In this paper, I compare this distribution to the corresponding distribution in the Deep Evolutionary Extragalactic Probe 2 (DEEP2; \\citealt{davis03a, faber03a}).  In Section \\ref{data}, I describe the two data sets. In Section \\ref{convert} I describe the procedure for the conversion of the SDSS sample into the DEEP2 sample. In Section \\ref{results} I compare the SDSS and DEEP2 samples. In Section \\ref{model} I present a simple model of the results. In Section \\ref{discussion}, I discuss the implications of the models and the data. In Section \\ref{conclusions}, I present conclusions.  Where necessary, I have assumed cosmological parameters $\\Omega_0 = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 100$ $h$ km s$^{-1}$ Mpc$^{-1}$ (with $h=1$). All magnitudes are (unless otherwise noted) AB-relative. Throughout, I will refer to a piece of software, {\\tt kcorrect}\\footnote{\\tt http://cosmo.nyu.edu/blanton/kcorrect}, for converting between various bandpass systems (Blanton et al. in preparation).  Except where noted, I will always be referring to the version of the software labeled {\\tt v4\\_1\\_4}. Finally, note that most of my comparisons will be in the \\band{0.1}{g} band, which is the SDSS $g$ band blueshifted by a factor 1.1, and has virtues extolled below. ",
        "conclusions": "\\label{conclusions}  I have presented a comparison of the low redshift ($z\\sim 0.1$) SDSS data with the high redshift ($z\\sim 1$) DEEP2 data. In general, the data suggest a quiet life for most galaxies on the blue sequence, with fewer than 10\\% being destroyed between those two epochs. It remains for better quantitative predictions of $\\Lambda$CDM than currently exist to determine whether that theory is consistent with this relatively peaceful history.  In particular, the data alone indicate that: \\begin{enumerate} \\item the red sequence of galaxies is redder by about 0.1 mag in the low redshift data; \\item the blue sequence of galaxies is redder by about 0.3 mag in the low redshift data; \\item the change in the blue sequence luminosity function is consistent with pure fading; and \\item the change in the red sequence luminosity function requires something more complex, perhaps an increase in number density. \\end{enumerate}  I have further interpreted these results in terms of simple models for the star-formation histories of galaxies. From these models I conclude: \\begin{enumerate} \\item that galaxies migrate quickly (1 Gyr) to the the red sequence when their star formation is sharply cut off; \\item the colors and luminosities of blue sequence galaxies are consistent with a decline in star-formation rate by about a factor of three in the last 8 Gyrs; \\item that consistency on the blue sequence suggests that to within $\\sim 10\\%$ mergers do not reduce the blue sequence of galaxies, contradicting predictions but in less obvious contradiction to direct observations of mergers; and \\item absent errors in the models or photometry at the 0.2 mag level, the data requires mergers to produce the luminous end of the red sequence. \\end{enumerate}  It is worth contrasting the change of the galaxy population with time to the change of the population with environment, since it illuminates a mystery regarding the latter. As \\citet{hogg04a} and others find, the red and blue sequences of galaxies do not change with environment, they are simply differently populated. In contrast, I find here that the red and blue sequences do change with time. Thus, the analogy often invoked that underdense regions are in some sense ``younger'' regions of the Universe is not a good analogy --- we can observe the Universe when it was younger and it looks quite different than do underdense regions today. This result supports the notion that the blue galaxies in dense regions had on average about the same formation epoch and subsequent history as blue galaxies in underdense regions."
    },
    "0512/gr-qc0512160_arXiv.txt": {
        "abstract": "Newly formed black holes are expected to emit characteristic radiation in the form of quasi-normal modes, called ringdown waves, with discrete frequencies. \\lisa~should be able to detect the ringdown waves emitted by oscillating supermassive black holes throughout the observable Universe.   We develop a multi-mode formalism, applicable to any interferometric detectors, for detecting ringdown signals, for estimating black hole parameters from those signals, and for testing the no-hair theorem of general relativity.  Focusing on \\lisa, we use current models of its sensitivity to compute the expected signal-to-noise ratio for ringdown events, the relative parameter estimation accuracy, and the resolvability of different modes. We also discuss the  extent to which uncertainties on physical parameters, such as the black hole spin and the energy emitted in each mode, will affect our ability to do black hole spectroscopy. ",
        "introduction": "\\label{intro}  The Laser Interferometer Space Antenna (\\lisa) is being designed to observe gravitational waves in the low-frequency regime, between $10^{-5}$ and $10^{-1}$ Hz.  A leading candidate source of detectable waves is the inspiral and merger of pairs of supermassive black holes (SMBHs). The signal should comprise three pieces: an inspiral waveform, a merger waveform and a ringdown waveform.  The inspiral waveform, originating from that part of the decaying orbit leading up to the innermost stable orbit, has been analyzed using post-Newtonian theory and black-hole perturbation theory, and extensive studies of the detectability of this phase of the signal have been carried out (see eg. \\cite{BBW,BC} and references therein). The nature of the merger waveform is largely unknown at present, and is the subject of work in numerical relativity.  The ringdown waveform originates from the distorted final black hole, and consists of a superposition of quasi-normal modes (QNMs).  Each mode has a complex frequency, whose real part is the oscillation frequency and whose imaginary part is the inverse of the damping time, that is uniquely determined by the mass $M$ and angular momentum $J$ of the black hole.  The amplitudes and phases of the various modes are determined by the specific process that formed the final hole.  The uniqueness of the modes' frequencies and damping times is directly related to the ``no hair'' theorem of general relativistic black holes, and thus a reliable detection and accurate identification of QNMs could provide the ``smoking gun'' for black holes and an important test of general relativity in the strong-field regime~\\cite{dreyer}.  In a pioneering analysis, Flanagan and Hughes (\\cite{fh}, henceforth FH) showed that, independently of uncertainties in the black hole spin and in the relative efficiency of radiation into various modes, the signal-to-noise ratio (SNR) for black hole ringdown could be comparable to that for binary inspiral. Consequently, both ringdown and inspiral radiation from SMBHs should be sufficiently strong relative to the proposed sensitivity of \\lisa~that they may both be detectable with high SNR throughout the observable universe.  With high SNR comes high accuracy, and hence the potential to {\\em measure} ringdown QNMs and to test general relativity.  The FH analysis provided some insight into the issue of {\\it detectability} of ringdown waves. However, to our knowledge, the problem of {\\it parameter estimation} from black hole ringdown with \\lisa~has not been discussed in depth in the literature to date.  Most existing studies have referred specifically to high-frequency ringdown sources and Earth-based interferometers~\\cite{echeverria,finn,kaa,creighton,nakano,tsunesada}.  The main purpose of this paper is to discuss the {\\em measurability} of ringdown QNM frequencies using \\lisa~by carefully developing a framework for analyzing QNM radiation, and then applying the standard ``Fisher matrix'' formalism for parameter estimation \\cite{finn}. We will treat both single-mode and multi-mode cases.  From a mathematical point of view the excitation amplitude of QNMs is an ill-defined concept, because QNMs are not complete~\\cite{Beyer,Beyer2,NollertPrice,Nollert,Szpak}. Following~\\cite{leaver2,A94,Nils97,GA} we will associate with each QNM an ``excitation coefficient'' that quantifies in a pragmatic (as opposed to rigorous) way the response of a black hole to perturbations with some given angular dependence. We will also define useful energy coefficients to characterize the energy deposited into various QNMs. Unfortunately, there is only sketchy information at present from numerical and perturbative simulations of distorted black holes, gravitational collapse or head-on collisions of two black holes as to what might be expected for the amplitudes, phases or energies of QNMs.  It is clear that the QNM content of the waveform will depend strongly on the initial conditions and on the details of the distortion.  In the absence of such information we will consider appropriate ranges of energy coefficients, and ranges of relative QNM amplitudes and phases as a way to assess the measurability of ringdown modes in some generality.  Although we will consider a range of SMBH masses from $10^5$ to $10^8 \\,M_\\odot$, we note that, for masses smaller than about $10^6~M_\\odot$, the damping time of the waves may be shorter than the light-travel time along the \\lisa~arms, and as a consequence the number of observable oscillations will be so small that a Fisher matrix approach may not be fully reliable. We will also consider the full range of SMBH angular momenta, from zero to near extremal.  \\begin{figure*}[t] \\begin{center} \\epsfig{file=l2m2eps.ps,width=6cm,angle=-90} \\caption{Value of $\\epsilon_{\\rm rd}$ required to detect the fundamental mode with $l=m=2$ (detection being defined by a SNR of $10$) at $D_L=3~$Gpc. For illustrative purposes here we pick the fundamental mode with $l=m=2$, but the dependence on $(n,l,m)$ is very weak.  The three curves correspond to $j=0$ (solid), $j=0.8$ (dashed) and $j=0.98$ (dot-dashed), where $j=J^2/M=a/M$ is the dimensionless angular momentum parameter of the hole. The ``pessimistic'' prediction from numerical simulations of head-on collisions is $\\epsilon_{\\rm rd}=10^{-3}$ (as marked by the dashed horizontal line), so we should be able to see all equal-mass mergers with a final black hole mass larger than about $\\sim 10^5~M_\\odot$ (the vertical line is just a guide to the eye).  The dip in the curves is a consequence of white-dwarf confusion noise in the \\lisa~noise curve. \\label{epsrd}} \\end{center} \\end{figure*}  One of our conclusions is that the prospects for detection of ringdown radiation by \\lisa~are quite encouraging.  Figure \\ref{epsrd} shows the value of the fraction of ``ringdown energy'' $\\epsilon_{\\rm rd}$ (defined as the fraction of the black hole mass radiated in ringdown gravitational waves) deposited in the fundamental ``bar mode'' with $l=m=2$ (assuming that mode dominates) that is required for the mode to be detectable by \\lisa~with a SNR of 10 from a distance of 3 Gpc. Three values of the dimensionless angular momentum parameter $j \\equiv J/M^2 = a/M$ are shown: zero, 0.8 (an astrophysically interesting value), and 0.98. Recall that $0 \\le j \\le 1$, spanning the range from Schwarzschild to extremal Kerr black holes. For SMBH masses between $10^6$ and $10^7 \\,M_\\odot$, deposition energies as small as $10^{-7}$ of the mass should be detectable.  We show that this conclusion is not strongly dependent on $(l,m)$, or on the overtone index $n$.  \\begin{figure*}[t] \\begin{center} \\epsfig{file=l2m2all.ps,width=6cm,angle=-90} \\caption{Errors (multiplied by the signal-to-noise ratio $\\rho$) in measurements of different parameters for the fundamental $l=m=2$ mode as functions of the angular momentum parameter $j$. Solid (black) lines give $\\rho \\sigma_j$, dashed (red) lines $\\rho \\sigma_M/M$, dot-dashed (green) lines $\\rho \\sigma_A/A$, dot-dot-dashed (blue) lines $\\rho \\sigma_{\\phi}$, where $\\sigma_k$ denotes the estimated rms error for variable $k$, $M$ denotes the mass of the black hole, and $A$ and $\\phi$ denote the amplitude and phase of the wave. \\label{errs-intro}} \\end{center} \\end{figure*}  We also find that accurate measurements of SMBH mass and angular momentum may be possible.  For detection of the fundamental $l=m=2$ bar mode, for example, Figure \\ref{errs-intro} shows the estimated error (multiplied by the SNR, $\\rho$) in measuring the SMBH mass $M$, angular momentum parameter $j$, QNM amplitude $A$, and phase $\\phi$ (see Sec. \\ref{onemode} for definitions). For example for an energy deposition of $10^{-4} M$ into the fundamental mode of a $10^6 \\,M_\\odot$ SMBH with $j=0.8$ at 3 Gpc ($\\rho \\sim 200$), $M$ and $j$ could be measured to levels of a percent; if the energy deposition is only $10^{-6}$, they could still be measured to 10 percent. Generalizing to multi-mode detection (specifically to detection of two modes with a range of relative amplitudes), we find similar results.  \\begin{figure*}[t] \\begin{center} \\epsfig{file=overweak.ps,width=6cm,angle=-90} \\caption{``Critical'' SNR $\\rho_{\\rm crit}$ required to resolve the fundamental mode ($n=0$) from the first overtone ($n'=1$) with the same angular dependence ($l=l'$, $m=m'$). We assume the amplitude of the overtone is one tenth that of the fundamental mode. Solid lines refer to $m=l,..,1$ (bottom to top), the dotted line to $m=0$, and dashed lines to $m=-1,..,-l$ (bottom to top). \\label{milanfar-fig2}} \\end{center} \\end{figure*}  \\begin{figure*}[t] \\begin{center} \\begin{tabular}{ccc} \\epsfig{file=overstrongl2.ps,width=4.5cm,angle=-90} & \\epsfig{file=overstrongl3.ps,width=4.5cm,angle=-90} & \\epsfig{file=overstrongl4.ps,width=4.5cm,angle=-90} \\\\ \\end{tabular} \\caption{``Critical'' SNR $\\rho_{\\rm both}$ required to resolve {\\it both} the frequency and the damping time of the fundamental mode ($n=0$) from the first overtone ($n'=1$) with the same angular dependence ($l=l'$, $m=m'$). We assume the amplitude of the overtone is one tenth that of the fundamental mode.  Solid lines refer to $m=l,..,1$ (top to bottom), the dotted line to $m=0$, and dashed lines to $m=-1,..,-l$ (top to bottom, unless indicated).  In the color versions, we used black for the modes with $l=|m|$, red for those with $0<|m|<l$ and blue for $m=0$. For $l=m$ the critical SNR grows monotonically as $j \\to 1$. \\label{milanfar-fig3}} \\end{center} \\end{figure*}  In order to test the no-hair theorem, it is necessary (though not sufficient) to resolve two QNMs~\\cite{dreyer}; roughly speaking, one mode is used to measure $M$ and $j$, the other to test consistency with the GR prediction.  Using a simple extension of the Rayleigh criterion for resolving sinusoids, we estimate the SNR required to resolve the frequencies and/or the damping times of various pairs of modes, as a function of the angular momentum parameter $j$.  For example, to resolve the fundamental $(n=0)$ mode for a given $(l,m)$ from the first overtone $(n=1)$ for the same $(l,m)$, the critical SNR required to resolve either frequency or damping time is shown in Fig. \\ref{milanfar-fig2}, while the SNR required to resolve both is shown in Fig. \\ref{milanfar-fig3}. These values assume that the amplitude of the overtone is $1/10$ that of the fundamental mode.  Comparing Figs. \\ref{milanfar-fig2} and \\ref{milanfar-fig3} with Fig. \\ref{insp-rd} in Sec. \\ref{singlemodenum}, we infer that tests of the no-hair theorem should be feasible, even under rather pessimistic assumptions about the ringdown efficiency $\\epsilon_{rd}$, as long as the first overtone radiates a fraction $\\sim 10^{-2}$ of the total ringdown energy.  However, resolving {\\em both} frequencies and damping times typically requires a SNR greater than about $10^3$. This is only possible under rather optimistic assumptions about the radiative efficiency, and it can be impossible if the dominant mode has $l=m=2$ and the black hole is rapidly spinning (solid black line in the left panel of Fig.~\\ref{milanfar-fig3}). Requiring SNRs at least as large as $10^2$ implies that resolving QNMs will be impossible for redshifts larger than about 10.  The remainder of this paper provides details. In Sec.~\\ref{preliminaries} we introduce our notation and formalism, and clarify some conceptual issues related to the QNM decomposition of gravitational waveforms. In Sec.~\\ref{onemode} we compute the single-mode SNR assuming that only one mode dominates the ringdown. As a first step we update the FH analysis of \\lisa's SNR for detection of single-mode waveforms.  With respect to FH we use a better semianalytic approximation of the \\lisa~noise curve, we consider different QNM frequencies, and we compare with the expected SNR for inspiral as computed in \\cite{BBW}. We explore the angular momentum dependence by considering black holes with $j\\equiv a/M=0,~0.8,~0.98$, and confirm the FH expectation that angular momentum does not have a big effect on the SNR. Uncertainties in the ringdown efficiency $\\epsilon_{\\rm rd}$ have a larger impact, since $\\rho\\sim \\sqrt{\\epsilon_{\\rm rd}}$. In Sec.~\\ref{onemode-pe} we assess the accuracy of parameter estimation in single-mode situations, revisiting the analyses in \\cite{echeverria,finn,kaa}, and show that it is in general very good. A more detailed analysis shows that, for counterrotating ($m<0$) and axisymmetric ($m=0$) modes, rotation doesn't necessarily help parameter estimation, and for some counterrotating modes the error can even blow up at ``critical'' values of the black hole angular momentum. Following preliminary remarks in Sec.~\\ref{excitation} describing mathematical issues in and model predictions for multi-mode ringdowns, in Sec.~\\ref{multimode} we generalize to a two-mode analysis, computing the SNR and errors on parameter estimation. We find that it is computationally convenient to treat separately the case of waveforms in which the modes have different angular dependence ($l\\neq l'$ or $m\\neq m'$) and the case where the overtones have the same angular dependence as the fundamental mode. In Sec.~\\ref{nohair} we determine the minimum signal-to-noise ratio required to discriminate between different quasi-normal mode pairs. Conclusions and perspectives for future work are presented in Sec.~\\ref{conclusions}.  In the Appendices we collect various numerical results and technical calculations. Appendix \\ref{deltafun} discusses the accuracy of a useful semianalytic approach to the calculation of the SNR (the $\\delta$-function approximation, valid for ringdown signals with large quality factor). In Appendix \\ref{app:FisherOne} we present an explicit calculation of the Fisher matrix using different conventions and different numbers of free parameters in the ringdown waveform. Appendix \\ref{app:noise} describes our semianalytical model of the \\lisa~noise.  Appendix \\ref{recover-rayleigh} discusses a particular aspect of our criterion for resolving normal modes.  Finally, Appendix \\ref{app:QNM} lists for reference the values of the complex frequencies and the angular separation constants of the spin-weighted spheroidal harmonics for selected normal modes, and also displays analytical fits of the QNM frequencies accurate within a few percent. ",
        "conclusions": "\\label{conclusions}  We have presented a general framework for analysing the detectability of quasi-normal ringdown gravitational waves from massive black holes, and for using them to estimate parameters of the hole and to test the general relativistic no-hair theorem.  In this initial work, we made a number of simplifying assumptions, including the reality of the spin-weighted spheroidal harmonics (thus simplifying the angle averaging), the restriction to one-mode or two-mode situations, and the assumption of large quality factors in deriving analytic expressions.  In future work, we hope to explore the consequences of relaxing some of these assumptions. Although we focused on detection of modes using \\lisa, our framework is equally applicable to ground-based interferometers.  For example, the non-angle-averaged case can be studied using Monte Carlo simulations in place of angle averaging, thereby permitting inclusion of the fully complex spheroidal harmonics. Multi-mode calculations involving more than two modes should be carried out. Our discussion of testing the no-hair theorem has focused only on resolvability of modes; we hope to use our tools to perform in detail the no-hair test suggested in \\cite{dreyer}. A crucial point is to determine the excitation of modes.  When they have different angular dependence, we must use numerical relativistic simulations of mergers, or minimally simulations of distorted black holes with initial data that mimics the merger.  For overtones with the same angular dependence, we can use perturbative and analytical techniques to estimate the relative amplitudes; this work is in progress.  It will be important to extend the analysis from matched-filtering to more sophisticated data-analysis techniques (eg. the tiling method used by the TAMA group~\\cite{nakano,tsunesada}) and combine them with time-delay interferometry (TDI), especially in the case of short damping times.  Another important question is the effect of combining parameter estimation from ringdown with parameter estimation from inspiral.  This may make it possible to improve the measurement of mass or angular momentum by taking advantage of ``prior'' information, and may permit one to deduce how much energy and angular momentum is radiated in the merger phase."
    },
    "0512/astro-ph0512311_arXiv.txt": {
        "abstract": "We investigate the stability of differentially rotating proto-neutron stars (PNSs) with a toroidal magnetic field. Stability criteria for nonaxisymmetric MHD instabilities are derived using a local linear analysis. PNSs are expected to have much stronger radial shear in the rotation velocity compared to normal stars. We find that nonaxisymmetric magnetorotational instability (NMRI) with a large azimuthal wavenumber $m$ is dominant over the kink mode ($m=1$) in differentially rotating PNSs. The growth rate of the NMRI is of the order of the angular velocity $\\Omega$ which is faster than that of the kink-type instability by several orders of magnitude. The stability criteria are analogous to those of the axisymmetric magnetorotational instability with a poloidal field, although the effects of leptonic gradients are considered in our analysis. The NMRI can grow even in convectively stable layers if the wavevectors of unstable modes are parallel to the restoring force by the Brunt-V\\\"ais\\\"al\\\"a oscillation. The nonlinear evolution of NMRI could amplify the magnetic fields and drive MHD turbulence in PNSs, which may lead to enhancement of the neutrino luminosity. ",
        "introduction": "Neutron stars are formed in the aftermath of supernova explosions associated to the gravitational core-collapse of massive stars (8 -- $30M_{\\sun}$) at the end of their evolutions. The binding energies released by core-collapse are stored in the interior of a newly born hot neutron star, and those energies are emitted as the neutrinos in the explosion stage. Neutrino emission from proto-neutron stars (PNSs) is one of the key processes of the delayed explosion scenario of core-collapse supernovae, in other words, the size of neutrino luminosities and the energy deposition efficiency determine whether the explosion will succeed or not (Bethe 1990).  Convection in PNSs is considered to play an important role in enhancing the neutrino luminosities (Epstein 1979; Lattimer \\& Mazurek 1981; Arnett 1987; Burrows \\& Lattimer 1986; Burrow \\& Fryxell 1993). Other hydrodynamic instabilities may also contribute to the enhancement of neutrino luminosities (Bruenn \\& Dineva 1996; Mezzaacappa et al. 1998a; Miralles et al. 2000, 2002, 2004). Mayle (1985) and Wilson \\& Mayle (1988, 1993) have argued that the material in the PNSs would be unstable to a double diffusive instability which is referred to as ``neutron finger''. Bruenn et al. (2004) have found a new double diffusive instability called ``lepto-entropy finger instability''. However, despite considering various hydrodynamic instabilities, recent numerical simulations have not succeeded in producing sufficient neutrino luminosities leading to the delayed supernova explosions (Burrows et al. 1993; Keil et al. 1996; Mezzacappa et al. 1998b; Buras et al. 2003).  One of the primary causes of the inefficient neutrino luminosity is the strong radial gradients of the entropy and leptonic fraction. Stable stratification of these quantities appears at the outer layer of PNSs and suppresses convective motions (Janka \\& M\\\"uller 1996). Convective regions are located deep inside of PNSs and surrounded by a convectively stable layer in which the surface fluxes of $\\nu_e$ and $\\tilde{\\nu_e}$ are mainly built up (Buras et al. 2003). Thus the convection in PNS has little influence on the emission of neutrinos and is irrelevant for the supernova dynamics.  On the other hand, MHD instabilities could be a seed of turbulence even in the stably stratified layer. Magnetic effects on the dynamics of supernovae have been examined actively since the importance of magnetorotational instability (MRI) in supernova cores is pointed out (Akiyama et al. 2003). Two-dimensional MHD simulations of the collapse of supernova cores indicate that the shapes of shock waves and the neutrino sphere can be modified by the effects of magnetic fields (Kotake et al. 2004; Yamada \\& Sawai 2004; Takiwaki et al. 2004). Thompson et al. (2005) suggest that the turbulent viscosities sustained by the MRI convert the rotational energies of supernova core into the thermal energies and may affect the supernova dynamics. However, effects of magnetic phenomena inside of the PNSs have not been understood yet and still have many uncertainties.  Therefore, in this paper we analyze the linear growth of MHD instabilities in the stably stratified layer of PNSs, and discuss the possibility of enhancement of neutrino luminosities. Especially, we focus on the instabilities caused by a strong differential rotation. We assume the toroidal component of the field is dominant in the supernova cores. Local stability of stellar objects with a rotation and toroidal magnetic fields have been investigated by many authors (Fricke 1969; Tayler 1973; Acheson 1978; Pitts \\& Tayler 1985; Schmitt \\& Rosner 1986; Spruit 1999, 2002). Spruit (1999) concluded that the kink-type instability (Tayler instability) grows at first among various MHD instabilities in stellar radiative zone. However, most of the previous work assume a weak differential rotation because their motivation is for the interiors of main sequence stars like the Sun. The stability analysis for the cases of a strong differential rotation, which is expected in PNSs, has not been carried out sufficiently. In such cases, nonaxisymmetric magnetorotational instability (NMRI) could be dominant over the kink-type instability (Balbus \\& Hawley 1992). The nonlinear evolutions of the NMRI can initiate and sustain MHD turbulence and amplify magnetic fields, which may lead to the enhancement of neutrino luminosity via magnetoconvection, viscous heating, and the other MHD phenomena.  The paper is organized as follows. In \\S2 we obtain the dispersion equation that determines the stability of PNSs. In \\S3 we discuss the general properties of the dispersion relation, and in \\S4 stability criteria and the maximum growth rate are derived by an analytical approach. In \\S5 we apply our results to the interiors of PNSs, and examine the linear growth of MHD instabilities. The conditions that the NMRI can be dominant over the kink-type instability and the nonlinear effects of MHD instabilities are discussed in \\S6. Finally, we summarize our main findings in \\S7. ",
        "conclusions": "\\label{sec:discussion} \\subsection{Effects of Diffusion Processes}  In the interior of PNSs, neutrino radiation plays an important role in the heat and lepton transports. The thermal and lepton diffusions caused by the neutrino radiation can reduce both the thermal and leptonic buoyant frequencies significantly. The coefficients of the thermal and chemical diffusivities in PNSs are given by \\begin{equation} \\kappa_{\\rm{T}} \\simeq 10^9 \\ Y^{-1}_{1/3} l^2_{30} \\frac{T_4}{\\mu^{6}_{20}} \\ {\\rm cm}^2 \\ \\sec^{-1} \\;, \\ \\ \\ \\nu_{\\rm{chem}} \\simeq 10^9 \\ l^2_{30} \\frac{T_4}{\\mu^{6}_{20}} \\ \\ {\\rm cm}^2 \\ \\sec^{-1} \\;, \\label{eq45} \\end{equation} where $Y_{1/3}$ is the electron fraction in units of $1/3$, $l_{30}$ is a length scale in units of 30 km, $T_4$ is the temperature in units of 4 MeV, and $\\mu_{20}$ is the electron chemical potential normalized by 20 MeV (Socrates et al. 2005).  These diffusion processes must be considered if the wavelength of fluctuations is shorter than the critical length scale that can be estimated by equating the diffusion time and the period of the buoyant oscillations. The critical length scale of the thermal diffusion is given by $\\lambda_{\\rm{T}} \\sim \\sqrt{\\kappa_{\\rm{T}}/ N_{\\rm{T}}} \\simeq 10^3$ cm, where we assume $N_{T} \\simeq 1000$ sec$^{-1}$. The critical length of the chemical diffusion may be a few times larger than that of the thermal one, because the leptonic buoyant frequency is smaller than $N_{\\rm T}$ (Thompson \\& Murray 2000). Thus, our analysis is valid for the wavelength of the perturbations longer than the diffusion length $\\lambda_{\\rm T}$.  However, the diffusive processes could affect the growth of the NMRI for the shorter wavelength $\\lambda < \\lambda_{\\rm{T}}$. The primary effect of the thermal and chemical diffusions is to reduce the restoring force caused by the stable stratification (Acheson 1978; Spruit 1999). Therefore, the maximum growth rate of NMRI increases and approaches to  the case of $N/\\Omega \\rightarrow 0$ when $\\lambda \\ll \\lambda_{\\rm T}$. Thus, the growth of these shorter wavelength modes may be dominant in the stably stratified layers of PNSs.     \\subsection{NMRI Modes vs. Kink Modes} The conditions that the NMRI modes dominate over the kink modes can be derived from our simplified dispersion equation~(\\ref{eq27}). The NMRI modes grow faster if the azimuthal wavenumber $m_{\\max}$ is more than 2. The shaded (white) area in Figure~\\ref{figure7}a represents the NMRI (kink) dominant region in the $q$-$\\theta$ diagram. The boundary of these two areas is given by where $m_{\\rm{max}} = 2$ is fulfilled in equation~(\\ref{eq36}). In this figure, we assume $N/\\Omega=10$ and $\\Omega/\\omega_A=30$ which are the typical values in PNSs. The NMRI modes are dominant at most of the regions if the shear parameter $q$ is relatively large ($q \\gtrsim 10^{-2}$). The kink modes become important only when the shear parameter is extremely smaller than unity. The boundary curve is little changed by the size of the Brunt-V\\\"ais\\\"al\\\"a frequency.  For comparison, a similar picture for the solar radiative zone is shown by Figure~\\ref{figure7}b. In the solar radiative zone, the Brunt-V\\\"ais\\\"al\\\"a frequency is much larger than the angular velocity $N/\\Omega \\simeq 10^3$ (Spruit 2002). Here we assume the Alfv\\'en frequency $\\Omega/\\omega_A = 5$ as the typical value in the solar radiative zone. The parameter region where the kink modes are dominant extends to the larger $q$ regime because of the smaller ratio of $\\Omega/\\omega_A$. Only when the shear rate is about unity, the NMRI modes can be important for this case. According to the helioseismology, the solar radiative zone exhibits a nearly rigid rotation ($q \\ll 1$). Thus, the kink modes are dominant and Tayler-Spruit dynamo would be amplifying the magnetic fields (Spruit 2002). \\subsection{Nonlinear Growth of the NMRI} The NMRI can grow in the stably stratified layer of PNSs regardless of the size of the Brunt-V\\\"ais\\\"al\\\"a frequency. The nonlinear growth of the NMRI could initiate and sustain MHD turbulence and amplify magnetic fields. As is well known in the contexts of accretion disks, MHD turbulence leads to the angular momentum transport. Thus the rotational configurations within PNSs evolve toward a rigid rotation. If the PNSs will rotate rigidly, the neutrino sphere become oblate and the neutrino luminosies could be enhanced in the polar region (Kotake et al. 2003). Thus, the polar region would be heated by neutrinos preferentially.  Turbulent viscosity generated by the NMRI can convert the rotational energies into the thermal energies, so that the neutrino luminosities from the PNSs may be increased (Thompson et al. 2005). Ramirez-Ruiz \\& Socrates (2005) suggest that the heating of materials through magnetic reconnection may also lead to the enhancement of neutrino luminosity. We would like to stress that the magnetoconvection and magnetic Rayleigh-Taylor instability of magnetic flux tube would assist the mixing of materials in the layers where is heretofore thought to be convectively stable (Parker 1974, 1975, 1979).  Although these effects are absolutely a matter of speculation, it is sure that these effects are not considered in the present scenarios of core-collapse supernovae. Therefore, we anticipate that these multiple magnetic effects can enhance the neutrino luminosity sufficiently to yield explosions in models that would otherwise fail. Because these magnetic phenomena are essentially three-dimensional, we must perform three-dimensional MHD simulations to reveal those effects quantitatively. This is our future work."
    },
    "0512/astro-ph0512582_arXiv.txt": {
        "abstract": "We perform numerical simulations of cosmic string evolution with intercommuting probability $P$ in the range $5\\times 10^{-3}\\le P\\le 1$, both in the matter and radiation eras, using a modified version of the Allen-Shellard code. We find that the dependence of the scaling density on $P$ is significantly different than the suggested $\\rho\\propto P^{-1}$ form. In particular, for probabilities greater than $P\\simeq 0.1$, $\\rho(1/P)$ is approximately flat, but for $P$ less than this value it is well-fitted by a power-law with exponent $0.6^{+0.15}_{-0.12}$. This shows that the enhancement of string densities due to a small intercommuting probability is much less prominent than initially anticipated. We interpret the flat part of $\\rho(1/P)$ in terms of multiple opportunities for string reconnections during one crossing time, due to small-scale wiggles. We also propose a two-scale model incorporating the key physical mechanisms, which satisfactorily fits our results over the whole range of $P$ covered by the simulations. ",
        "introduction": "Introduction}  Cosmic strings have recently been reincarnated in the form of `cosmic superstrings' or, more specifically, fundamental strings and one-dimensional $D$-branes produced at the end of brane inflation \\cite{DvalTye,Quevedo,SarTye,DvalVil} (for reviews see for example \\cite{PolchIntro,Kibblerev,DavKib,Majumdar}). These objects have properties which can be significantly different than usual field theory strings, and this opens up the possibility that cosmic string observations could yield information about string theory, while providing new ways to constrain various brane inflation models \\cite{Babichev,cycloops}. Cosmic superstrings typically have smaller tensions, in the range $10^{-12} < G\\mu < 10^{-7}$ \\cite{SarTye,JoStoTye2,PolchIntro}, and reconnect with probabilities which can be significantly less than unity \\cite{JoStoTye2,PolchProb}, unlike ordinary cosmic strings \\cite{Shell_Recon}. This reduced reconnection probability results in an enhancement of the string number density today \\cite{JoStoTye2,DvalVil,EDVOS}. Cosmic superstring networks can also consist of more than one type of string, which can zip together to produce trilinear vertices with links stretching between them \\cite{DvalVil,PolchStab}.  The evolution of cosmic superstring networks has recently been studied in Refs.~\\cite{EDVOS,Sak,Mart_nonint,MTVOS,CopSaf}. The key issue is whether these networks reach a \\emph{scaling regime}, that is one in which their characteristic lengthscale stays constant relative to the horizon. This self-similar evolution is a well-known property of usual strings, which reconnect with probability of order unity, at least in the abelian case where no trilinear vertices can be formed (though some evidence for non-abelian scaling has also been found \\cite{VachVil,McGraw,McGraw1,MTVOS}). In all the above studies of various models of cosmic supersting networks, evidence for scaling behaviour was found, though the issue is not yet completely resolved.  Assuming that an attractor scaling solution exists, an interesting question to ask is how the density of the scaling network depends on the intercommuting probability $P$. Using a simple one-scale model \\cite{Kibble}, Jones, Stoica and Tye have argued in Ref.~\\cite{JoStoTye2} that the expected behaviour is $\\rho \\propto P^{-2}$. This is a dramatic effect, as a probability of $10^{-2}$, for example, would lead to an enhancement of the string density by a factor of $10^{4}$ compared to ordinary strings: we could live in a universe packed with $F/D$-string relics from a brane inflation era, which could in principle be observed to extract information about physics at the string scale! However, string intercommuting is a small-scale process and can be expected to depend crucially on small-scale wiggles on strings, which are not captured by the simple one-scale model. One could still use a one-scale (or VOS \\cite{vos0,vosk}) model but introduce an \\emph{effective} intercommuting probability $P_{\\rm eff}=f(P)$ \\cite{EDVOS}, whose form should be determined by numerical simulations or an analytic model for small-scale structure. One should expect that, since small-scale wiggles would mean that strings may have more than one opportunity to reconnect in each encounter, the effect of $P$ on string density could be somewhat counterbalanced by the strings' small-scale-structure. Thus a power law weaker than $\\rho \\propto P^{-2}$ may be anticipated.  Indeed, Sakellariadou recently performed flat space simulations of cosmic strings with reconnection probabilities in the range $10^{-3}\\le P \\le 1$ and found a power law $\\rho \\propto P^{-1}$ in the range $10^{-3}\\le P \\le 0.3$ \\cite{Sak}. The purpose of this letter is  to investigate the dependence of the string energy density $\\rho$ on the intercommuting probability $P$ for strings evolving in more realistic cosmological backgrounds. We perform numerical simulations in both the matter and radiation era, and find a power law significantly weaker than that of Ref.~\\cite{Sak}. We explain our results in terms of string small-scale structure and the importance of a second scale in the problem. We thus propose a two-scale analytic model which provides a good fit to our numerical results. ",
        "conclusions": "Discussion}  We can interpret the flat part of Fig.~\\ref{rho_of_P_fig} in terms of string small-scale-structure: if two colliding strings are wiggly, they may have more than one opportunity to reconnect during their encounter.  On large scales, correlation length segments move together slowly and coherently ($v_{\\rm c}\\approx 0.15$ for $P=1$ \\cite{MartShell_sss}), but their small-scale structure is relativistic and rapidly oscillates ($v_{\\rm rms} \\approx 1/\\sqrt2$). One therefore expects the strings to cross several times as they pass, thus the overall process should be described by an \\emph{effective} probability of reconnection $P_{\\rm eff}>P$. Reducing $P$ by a factor of a few, for example, should not have a big impact on the string density \\cite{EDVOS}. Indeed, we can understand the flat part of Fig.~\\ref{rho_of_P_fig}, and in particular explain the position where the slope starts to change, by using the following simple model for string interactions.  Consider two colliding straight strings, described on small-scales by a monochromatic mode of amplitude $A$ and period $T$. If $v_{\\rm c}$ is their relative coherent speed, the duration of the collision is approximately $\\Delta t\\approx 2A/v_{\\rm c}$. Now on small scales the strings are oscillating with period $T$, so the number of reconnection opportunities during their crossing time is $N\\approx \\Delta t / T = 2A / v_{\\rm c} T$. But since on these scales string modes are propagating with velocity $c\\sim 1$, $T$ is equal to the wavelength of the mode, and so $A/T$ is a measure of string wiggliness. In particular, for large amplitudes we may write $1+A/T \\sim \\tilde\\mu$ from which we can infer that \\be\\label{N_int} N\\sim \\frac{2(\\tilde\\mu - 1)}{v_{\\rm c}}. \\ee Interestingly, this simple estimate gives about the right value for the position where the curve changes slope in Fig.~\\ref{rho_of_P_fig}. Indeed, from Fig.~\\ref{mu_tilde_fig} we see that for $P>0.1$, $\\tilde\\mu$ is in the range $1.6\\le \\tilde\\mu\\le 1.7$, while we know that string velocities at the scale of the correlation length are approximately $v\\approx 0.15$. This gives $N\\sim10$ intercommuting opportunities. One therefore expects that to see a significant effect on the scaling string density, $P$ should be reduced by a factor greater than about $10$.  Furthermore, we can estimate the effective probability by considering the  $N$ reconnection opportunities, assuming each has an independent probability $P$, that is \\be\\label{P_eff} P_{\\rm eff}=1-(1-P)^N. \\ee For $N\\simeq 10$, a probability as low as $0.2$ yields $P_{\\rm eff}\\simeq 1$, and significant deviations from unity only occur for $P\\simeq 0.1$ for which $P_{\\rm eff}\\simeq 0.65$, as observed in Fig.~\\ref{rho_of_P_fig}. In the limit $P\\ll 1$, Eq.~(\\ref{P_eff}) gives $P_{\\rm eff}\\simeq NP\\approx [2(\\tilde\\mu-1)/v_{\\rm c}]P$ and so $P$ can only be enhanced by a factor of order $10$ (see Fig.~\\ref{mu_tilde_fig}), which results in an effective reconnection probability that is still much less than unity.  We now turn to explaining the constant slope part of Fig.~\\ref{rho_of_P_fig} for small probabilities. First recall that a simple one-scale model with $P<1$ predicts $\\rho\\propto P^{-2}$ in direct contradiction with our numerical results. As discussed above, the introduction of an effective intercommuting probability can dramatically change this for relatively large $P$ (say $P>0.1$), thus explaining the flat part of Fig.~\\ref{rho_of_P_fig}. However, to accommodate the constant slope part for $P\\ll 1$ one would need $P_{\\rm eff}\\propto P^{\\kappa}$, with $\\kappa$ a constant less than $1/2$. At present there seems no motivation for such a $P_{\\rm eff}$; instead, Eq.~(\\ref{P_eff}) suggests that $P_{\\rm eff}$ is linear in $P$ for $P\\ll 1$. Thus, it appears that no one-scale model can fit our numerical results for $P\\ll 1$: as the string intercommuting probability decreases, the one-scale approximation becomes increasingly poor.  This is in agreement with the results of Ref.~\\cite{ShellAll_GHV} in which  long string intercommutings were switched off in numerical simulations of evolving strings, while small loop production was allowed. It was found that this prevented the string density from scaling, that is, the inter-string distance $L$ was no longer proportional to $t$ (the solution being $\\rho\\propto L^{-2} \\propto t^{-7/8}$).  Nevertheless, the actual correlation length $\\xi$ along the string did scale at approximately the size of the horizon ($\\xi\\propto t$).  This is very similar to the situation we observe in our simulations. Reducing the intercommuting probability leads to an increase of the string density and therefore a decrease of the characteristic length associated with it (the interstring distance $L$). The correlation length $\\xi$, however, is comparatively unaffected by this, and stays at a scale of order the horizon (Fig.~\\ref{mu_tilde_fig}). This can be understood by considering the two distinct mechanisms for producing loops: (i) self-intersections of the same string and (ii) collisions between long strings. The former tends to chop off small loops, straightening the strings out (thus affecting mainly the correlation length $\\xi$) while also controlling the amount of small-scale structure $\\tilde\\mu$. Small loop production alone is a significant energy loss mechanism, but it is not sufficient to ensure scaling \\cite{ShellAll_GHV}. On the other hand, long string reconnections have much more dramatic effects introducing large-scale `bends' in the strings which catalyse the collapse of large regions of string and the formation of many more loops.  In contrast, these energy losses are sufficient to govern the interstring distance $L$ and cause scaling.  Small loop production, the first mechanism, is not greatly affected by reducing the intercommuting probability: once the string is sufficiently wiggly, left and right moving modes which fail to interact due to a small $P$ will keep propagating and meet more incoming modes, with which they will eventually interact and form loops. It might take longer, but such interactions are inevitable even for $P\\ll 1$. The relevant question is whether enough of these interactions can take place in each Hubble time in order to straighten the strings out at the horizon scale, but this seems reasonable given the much shorter timescale on which small loop production operates.  It appears to be confirmed by Fig.~\\ref{mu_tilde_fig}, where there is only a very weak build-up of small scale structure (a factor of 2-3 in $\\tilde \\mu$) while $P$ changes by over two orders of magnitude. On the other hand, long string intercommuting, the second mechanism, depends more crucially on $P$: two colliding strings have a given interval of time in which to interact, so intercommuting is no longer inevitable for relatively small $P<0.1$. A reduced $P$ necessarily means less string interactions and less energy dumped from the long string network in the form of loops. This leads to an increase in the long string density and thus a significant decrease of the characteristic length relative to the correlation/horizon scale.  To obtain an analytic model for such networks, it is therefore necessary to introduce two scales: a characteristic length $L$ (roughly the interstring distance) quantifying the energy density in strings, and a correlation length $\\xi$, defined to be the distance beyond which string directions are not correlated. Indeed, similar models have appeared in the literature (e.g. the three-scale model of Ref.~\\cite{AusCopKib}), but for normal $P=1$ cosmic strings it has been established that these two scales are comparable so the one-scale (velocity-dependent) approximation \\cite{vos0,vos,vosk}  can be surprisingly accurate \\cite{vostests}. Here, we develop a two-scale velocity-dependent model which fits our numerical results, again surprisingly well given its simplicity.  Consider a string network, characterised by tension $\\mu$, correlation length $\\xi$ and energy density $\\rho$. We define the characteristic lengthscale $L$ of the network by $\\rho\\equiv \\mu/L^2$ and note that the number of strings per correlation volume $V=\\xi^3$ is $N_{\\xi}=\\rho V/\\mu\\xi=\\xi^2/L^2$. If $v$ is the typical string velocity, each string intersects $N_{\\xi}-1$ other strings in time $\\Delta t=\\xi/v$, so we have $N_{\\xi}^2-N_{\\xi}$ intercommutings per correlation volume per time $\\Delta t$. Assuming that each intercommuting produces a loop of length $\\tilde c \\xi$ (this can be done formally by integrating an appropriate loop production function over all relevant loop sizes, which introduces the loop production parameter $\\tilde c$ \\cite{book}) the energy loss due to the formation of loops can be written as \\be\\label{delta_rho_loops_2s} \\left(\\frac{\\delta\\rho}{\\delta t}\\right)_{\\rm 2-scale} = \\frac{(N_{\\xi}^2-N_{\\xi})v}{\\xi^4}\\mu\\tilde c \\xi = \\tilde c \\rho \\left(\\frac{\\xi}{L^2}-\\frac{1}{\\xi}\\right). \\ee This is to be contrasted with the corresponding result for the one-scale model \\be\\label{delta_rho_loops_1s} \\left(\\frac{\\delta\\rho}{\\delta t}\\right)_{\\rm 1-scale} = \\tilde c \\frac{\\rho}{L}. \\ee As an interesting aside, we note that the last term in Eq.~(\\ref{delta_rho_loops_2s}) $\\propto 1/\\xi$ has the same form (though opposite sign) as the term which would need to be introduced to account for direct small loop production (or string radiation); it cannot itself cause $L$ to scale as $t$.  Our velocity dependent two-scale model can therefore be constructed by using the usual VOS model equations \\cite{vosk}, derived by performing a statistical averaging procedure on the Nambu-Goto equations of motion and energy momentum tensor, but using the phenomenological loop production term (\\ref{delta_rho_loops_2s}) instead of (\\ref{delta_rho_loops_1s}). The result is a system of two coupled ODEs, governing the time evolution of the characteristic length $L$ and the average velocity $v$ of string segments: \\begin{eqnarray}\\label{Ldt_2s} 2\\frac{dL}{dt} & =& 2\\frac{\\dot a}{a}L(1+v^2) + \\tilde c v \\left(\\frac{\\xi}{L}-\\frac{L}{\\xi}\\right) \\\\ \\label{vdt_2s} \\frac{dv}{dt} &=& (1-v^2)\\left(\\frac{k}{\\xi}-2\\frac{\\dot a}{a}v\\right) \\end{eqnarray} In Eq.~(\\ref{vdt_2s}), $k$ is the so-called momentum parameter, which is a measure of the angle between the curvature vector and the velocity of string segments and thus is related to the smoothness of strings \\cite{vos0}. Note that we have made no attempt to derive an evolution equation for the correlation length $\\xi$ as the simulations show that this depends weakly on $P$ and moreover, unlike $L$, it remains comparable to the horizon (Fig.~\\ref{mu_tilde_fig}). We will neglect this weak dependence of $\\xi$ on $P$ and take $\\xi=t$ in Eqs.~(\\ref{Ldt_2s}-\\ref{vdt_2s}).  To apply the model (\\ref{Ldt_2s}-\\ref{vdt_2s}) to our simulated networks, we have also introduced an effective intercommuting probability given by (\\ref{P_eff}), that is, we have replaced $\\tilde c$ in Eq.~(\\ref{Ldt_2s}) by $P_{\\rm eff}\\tilde c$. As a first approximation, we have neglected the dependence of $N$ on $P$ (this can be understood in terms of the dependence of $\\tilde\\mu$ on $P$, see Eq.~(\\ref{N_int}) and Fig.~\\ref{mu_tilde_fig}) and we have simply taken $N=10$ (a rough average value over the range of $P$). For the loop production parameter, we have used the value $\\tilde c=0.23$ of Refs.~\\cite{vostests,vosk} which fits both radiation and matter era runs in the $P=1$ case.  In Fig.~\\ref{2scale_fig} we plot the matter era scaling energy densities, obtained both from the simulations and our two-scale model (solid line), for string networks with intercommuting probabilities in the range $5\\times 10^{-3}\\le P\\le 1$. We see that the model provides a surprisingly good fit  (given its simplicity and the approximations made) of the numerical data and it reproduces the observed change of slope around $P\\approx 0.1$. We stress that we have made no special parameter choices to obtain this fit, and by modifying parameters further we could do much better.  Nevertheless, it is clear from the dashed line in Fig.~\\ref{2scale_fig} how the model could be improved by taking into account the dependence of the collision number $N$ on $\\tilde\\mu(P)$.  In this case, we have estimated the effective string tension $\\tilde\\mu$ as a function of $P$ from Fig.~\\ref{mu_tilde_fig}, which yields $N$ from Eq.~(\\ref{N_int}) (that is, 6--18 over the full range of $P$) and then $P_{\\rm eff}(P)$ from Eq.~(\\ref{P_eff}) which is finally input `by hand' in our two-scale model. The striking agreement that results provides strong motivation for understanding the key physical mechanisms governing the small-scale structure parameter $\\tilde\\mu$. Indeed, modifying $P$ seems to provide a useful testbed for observing the dynamics of $\\tilde\\mu$, but we leave a more sophisticated analysis for a future publication \\cite{AMS_inprep}. We note the need also to take into account the dependence of $\\xi$ on the reconnection probability, and a direct small-loop production term (as discussed above). However, here, we just highlight the fact that a simple version of our two-scale model with the addition of the effective probability of Eq.~(\\ref{P_eff}) seems to provide a satisfactory fit to the numerical data.   \\begin{figure}[!h] \\includegraphics[height=2.7in,width=3.0in]{Two_Scale.eps} \\caption{\\label{2scale_fig} Scaling string density obtained from simulations (data points with errors) and from the analytic two-scale model (\\ref{Ldt_2s}-\\ref{vdt_2s}) (solid line). The fit is improved further (dashed line) by phenomenologically incorporating the dependence of the effective string tension $\\tilde \\mu$ on the reconnection probability $P$.} \\end{figure}"
    },
    "0512/astro-ph0512257_arXiv.txt": {
        "abstract": "In an investigative 16 hour L band observation using the MERLIN radio interferometric array, we have resolved both the pulsar PSR B1951+32 and structure within the flat spectral radio continuum region, believed to be the synchrotron nebula associated with the interaction of the pulsar and its `host' supernova remnant CTB 80. The extended structure we see, significant at $\\sim$ 4.5 $\\sigma$, is of dimensions 2.5\" $\\times$ 0.75\", and suggests a sharp bow shaped arc of shocked emission, which is correlated with similar structure observed in lower resolution radio maps and X-ray images. Using this MERLIN data as a new astrometric reference for other multiwavelength data we can place the pulsar at one edge of the HST reported optical synchrotron knot, ruling out previous suggested optical counterparts, and allowing an elementary analysis of the optical synchrotron emission which appears to trail the pulsar. The latter is possibly a consequence of pulsar wind replenishment, and we suggest that the knot is a result of magnetohydrodynamic (MHD) instabilities. These being so, it suggests a dynamical nature to the optical knot, which will require high resolution optical observations to confirm. \\\\ ",
        "introduction": "Recent subarcsecond Hubble Space Telescope (HST) \\& Chandra observations of the inner Crab Nebula have yielded evidence of dynamic activity in close proximity to the pulsar with various shocks and `wisps' evolving in terms of  position, morphology and luminosity over timescales of days to weeks, implying local velocities of order 0.7c (Hester et al. 1996; Weisskopf et al. 2000). Observations in the radio using the Very Large Array (VLA) show evidence for similar time-varying morphological changes within the nebula but at scales of order 1\" (Bietenholz et al. 2001).  Such multiwavelength observations challenge our understanding of these \\\\ pulsar/plerion/supernova remnant (SNR) associations. Considering the Crab Nebula, the estimated current rate of particle injection derived via nebular X-ray luminosity differs markedly from the historical average rate determined from the radio emitting particles (Arons, 1998; Atoyan, 1999). The nebula's emission may be modelled as a function of the particle injection spectrum related to the spin down energy of the pulsar, yet the exact processes which link the thermalising pulsar wind to the observed synchrotron emission are not clear (Reynolds \\& Chevalier, 1984), neither is our understanding of the particle acceleration processes within the pulsar wind in the first place (Gallant \\& Arons, 1994; Lou, 1993).  Another `young' pulsar/SNR/X-ray plerion association that provides a working laboratory with which to test our understanding of such interactions following on from the Crab is that of the young radio, X-ray and $\\gamma$-ray pulsar PSR B1951+32 associated with the CTB 80 SNR. Early radio and X-ray observations noted a central plerion/spectrally flat region to the SW of the SNR,  within which PSR B1951+32 was detected, located within a concentration of nebular emission towards one edge of the flat spectral region (e.g. Strom et al. 1984). The association is valid based on this clear interaction between pulsar and remnant, with similar pulsar canonical age (107 kyr)  and dynamically derived SNR age (9.6 $\\times$ 10$^{4}$d$_{2.5kpc}$).  Strom (1987) performed a detailed survey of this central flat region with the VLA at various frequencies and baselines, all yielding maps with resolutions $\\sim$ 1\", and recent surveys using the WSRT at 6 cm (Strom \\& Stappers, 2000) have indicated considerable complexity within the spectrally flat core. Strom  was the first to comment on the `hot spot', located to the SW of the pulsar, suggesting that it was likely a consequence of the pulsar wind interacting with the remnant and creating a `wisp' like structure similar to that observed in close proximity to the Crab pulsar.  In contrast to the young Crab, it appears that the older PSR B1951+32 has caught up with its expanding remnant, resulting in the observed complex multiwavelength emission. As such, it represents an extremely interesting stage in the life cycle of a pulsar, when the neutron star penetrates and interacts with the remnant/swept up ISM region. Precisely what radiation gets emitted, and where, will define the constraints to any subsequent modelling efforts.   Moon et al. (2004) have recently reported an optical/X-ray analysis of the system using archival HST and Chandra data, in addition to new ground based optical and IR observations. The Chandra data clearly shows a cometary pulsar wind nebula which appears to be confined by a bow shock produced by the high velocity motion of the pulsar, which corresponds to Strom's previously defined radio `hot spot'. Optical/IR photometry have also indicated the presence of a synchrotron 'knot', as originally reported by Hester (2000) embedded within the cometary head observed by Chandra. Previous attempts to isolate putative optical counterparts to PSR B1951+32 have been concentrated around this structure, which is believed to be closely related to the pulsar - thus far two candidates have been proposed which are within the VLA error ellipse (Butler et al. 2002). Moon et al. (2004), based on their analysis, argue that only one is likely, and even then it is generally agreed that the `point source' involved may well be a background star or non-uniformity in the knot itself.   The key to resolving many of the outstanding aspects of the the system's multiwavelength geometry is a rigorous astrometric reference frame. In this short letter we report a 16 hour L band observation with the MERLIN radio interferometer. Both objects were resolved at a resolution of 150 mas, and despite the relatively low signal to noise of the extended emission associated with the shock front, we have been able to re-examine existing multiwavelength data using the MERLIN data as the astrometric reference. ",
        "conclusions": "We report the first high resolution radio observation of the inner PSR B1951+32 plerion using MERLIN at L band.We have resolved both the pulsar and apparent fine structure within the `hot spot' identified at lower resolution and believed to be a consequence of the pulsar wind interacting with swept up ISM/SNR material.  We have used our MERLIN data to register the astrometrically corrected archival HST observations of the field. Combined, these data indicate that the previously identified optical `knot' of synchrotron emission extends behind the pulsar, along a line that bisects the shock front emission. The dimensions of the optical knot and the VLA determined proper motion argue for a synchrotron cooling time that is consistent with particle replenishment from the pulsar wind.  The formation of the knot can also be attributed to the mechanisms outlined in Lou (1998) with the interaction of MHD wind streams, whilst the knot's luminosity can be maintained by particle injection from the pulsar wind. Variations in the knot's luminosity and morphology are anticipated as successive quasi-periodic disturbances emanate from the pulsar. This being so it argues for a fundamentally dynamical nature to the observed synchrotron knot which may only be really discernible using future HST or ground-based adaptive optics observations.  Finally, the MERLIN data definitively rules out the putative optical counterparts to PSR B1951+32 suggested by Butler et al. (2002) and Moon et al. (2004), and provides an unambiguous error box with which to assist future high time resolution searches."
    },
    "0512/astro-ph0512531_arXiv.txt": {
        "abstract": "{We present the first time resolved medium resolution optical spectroscopy of the recently identified peculiar Intermediate Polar (IP) \\object 1RXS\\,J154814.5-452845, which allows  us to precisely determine the binary orbital period ($\\rm P_{\\Omega}=9.87\\pm$0.03\\,hr) and the white dwarf spin period ($\\rm P_{\\omega}=693.01\\pm$0.06\\,s). This system is then the third  just outside the  purported $\\sim$6-10\\,hr IP orbital period gap and the fifth of the small group of long period IPs, which has a relatively high degree of asynchronism. From the presence of weak red absorption features, we identify the secondary star with a spectral type  K2$\\pm$2\\,V, which appears to have evolved on the nuclear timescale. From the  orbital radial velocities of emission and the red absorption lines a mass ratio $q=0.65\\pm0.12$ is found. The masses of the components are estimated to be $\\rm M_{WD} \\geq 0.5 \\,M_{\\odot}$ and $\\rm M_{sec}=0.4 -0.79\\,\\,M_{\\odot}$ and the binary inclination  $25^o < i \\leq 58^o$. A distance between 540-840\\,pc is estimated. At this distance, the  presence of peculiar absorption features surrounding Balmer emissions cannot be due to the contribution of the white dwarf photosphere and their spin modulation suggests an origin in the magnetically confined accretion flow. The white dwarf  is also not accreting at a particularly high rate ($\\rm \\dot M < 5\\times 10^{16}\\,g\\,s^{-1}$), for its orbital period. The  spin-to-orbit period ratio $\\rm P_{\\omega}/P_{\\Omega}$=0.02 and the low mass accretion rate suggest that  this system is far from spin equilibrium. The magnetic moment of the accreting white dwarf is found  to be  $\\mu < 4.1\\times10^{32}$\\,G\\,cm$^3$, indicating a low magnetic field system. ",
        "introduction": "Intermediate Polars (IP) are magnetic Cataclysmic Variables (mCVs) characterized by strong and hard ($\\rm kT \\sim$ 15\\,keV)   X-ray pulses usually at the rotational period of the accreting white dwarf (WD) ($\\rm P_{\\omega} << P_{\\Omega}$). The X-ray pulsations indicate that the accretion flow is magnetically channelled towards the WD polar regions. However, the lack of detectable optical-IR circular polarization in  most of these systems prevents the measure of the WD magnetic field.   It indicates  that the magnetic field of the accreting WD is lower (B$<$5-10\\,MG) than  that detected in the strongly magnetized (B$\\sim$10-230\\,MG) synchronous ($\\rm P_{\\omega} = P_{\\Omega}$) polar systems (see Warner \\cite{Warner95} for a comprehensive review of mCVs). The orbital periods of IPs are generally long ($\\rm P_{orb}>$4\\,hr), with only three systems confirmed to be IPs by X-rays below the  so-called 2-3\\,hr orbital period gap, whilst the polars are typically found at shorter ($<$4\\,hr) orbital periods, with most systems below the period gap. The different location of the two subclasses in the orbital period distribution of mCVs suggests that IPs will synchronize in their evolution path towards  short orbital periods and hence they might be progenitors of  polar systems (Norton et al. \\cite{norton04}). However, the difference in the magnetic field strengths was the prime reason for rejecting the hypothesis of IPs evolving into polars, although  recent works on magnetic field evolution in accreting WD might explain the discrepancy  (Cumming \\cite{Cumming02}). The  large variety of  observational properties found in the 40 or so systems known to date, still  need to be understood in terms of accretion and evolutionary state. In particular a wide  range of asynchronism seems to characterize this class (Woudt \\& Warner \\cite{woudtwarner}). Such a wide range was recently discussed by Norton et al. (\\cite{norton04}) in terms of spin equilibrium for magnetic accretion, which can take place on a wide variety of ways ranging from magnetized  accretion streams to extended accretion discs. Also, in a few systems the presence of a \"soft X-ray\" emission component, similar  to that observed in the polars (Buckley \\cite{Buckley}) raised  the evolutionary question of whether these soft X-ray IPs are the  true progenitors of the polars. Furthermore the apparent lack of IPs in the range of orbital periods $\\sim$6-10\\,hr (the so-called \"IP period gap\") was only recently noticed (Schenker et al.  \\cite{schenker}), with only one recently discovered system at 7.2\\,hr (Bonnet-Bidaud et al. \\cite{bonnet-bidaud05}) and with a handful of long period IPs (Buckley \\cite{Buckley}; G\\\"ansicke et al. \\cite{gaensicke05}), being the most peculiar ever known. Hence, our understanding of the evolutionary  relationships among mCVs and in particular of IPs is still very poor and the addition of new systems and the study of their properties has a great  potential to alleviate this problem.\\\\   The X-ray source \\object 1RXS\\,J154814.5-452845 (henceforth RX\\,J1548) was identified as an  Intermediate Polar (IP) by Haberl et al. (\\cite{Haberl02}) using optical and X-ray observations. The 693\\,s  WD rotational period  was identified in both optical and X-ray light curves, whilst sparse optical spectroscopy and photometric data did not allow an unambiguous identification of the orbital period. The two possible values of 9.37$\\pm$0.69\\,hr or 6.72$\\pm$0.32\\,hr, would add RX\\,J1548 to the small group of IPs with long orbital periods, both periods locating it in the IP period gap. Furthermore, RX\\,J1548 possesses a highly absorbed hot black-body soft X-ray component similar to that observed in the IP V2400\\,Oph, but different from that found in the soft IPs (PQ\\,Gem, V405\\,Aur) (de Martino et al. \\cite{demartino04}) which instead do not suffer from strong absorption. RXJ\\,1548 is also peculiar in  that its optical spectrum, showing broad  absorption features underneath Balmer emissions, is similar to those observed  in IP V\\,709 Cas (Bonnet-Bidaud et al \\cite{Bonnet-Bidaud01}), which might suggest the contribution of the underlying accreting WD atmosphere (Haberl et al. \\cite{Haberl02}).  \\noindent In this work we present extensive high temporal resolution optical spectroscopy  aiming to determine the true orbital period, to elucidate the nature of the peculiar absorption features and that of the stellar components of this binary.   \\begin{table}[t!] \\caption{Summary of the observations of RX\\,J1548.} \\label{obslog} \\centering \\begin{tabular}{c c c c c } \\hline \\hline \\noalign{\\smallskip} Date &  Start &  Spectral  Range     & $\\rm T_{expo}$ & N. Spectra\\\\ UT    &  UT   &  $\\AA$         & sec & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} May 27, 2003 & 01:38:32 & 4038 - 7020 & 90 & 92  \\\\ May 27, 2003 & 02:38:26 & 5900 - 8600 & 200 & 4  \\\\ May 28, 2003 & 00:10:56 & 4038 - 7020 & 90 & 160 \\\\ May 28, 2003 & 01:26:21 & 5900 - 8600 & 200 & 7 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} ",
        "conclusions": "Our optical spectroscopy of RX\\,J1548 has allowed to refine the orbital  and WD spin periods detected by Haberl et al. (\\cite{Haberl02})  and to derive information on this binary  as follows:  \\begin{itemize}  \\item  The orbital period of 9.87\\,hr locates RX\\,J1548 just outside the IP period gap, thus joining other four long period systems (AE\\,Aqr, V1062 Tau, RX\\,J173021.5-0559 and GK\\,Per). Because of lack of long term monitoring,  we are unable to assess whether RX\\,J1548 also shows a peculiar optical behaviour as observed in the well monitored AE\\,Aqr, V1062 Tau and GK\\,Per. We only detect a  strong  variability in the radial velocities of emission lines and a long term luminosity variation on a five years timescale.\\\\  \\item The 693\\,s spin period is detected in the radial velocities of emission lines and of the absorption of Balmer lines.  Though not truly spectrophotometric, our data do not reveal a strong spin or  beat periodicity in the continuum contrary to  previous photometric observations. \\\\  \\item From weak absorption features in the red portion of the spectrum we identify the spectral type of the secondary star as a K2$\\pm$2\\,V. The orbital period and the spectral type suggest that the donor star in RX\\,J1548 has undergone nuclear evolution as also found in other long period CVs.\\\\  \\item The amplitude of the orbital radial velocites of emission and weak absorption lines allows us to determine the mass ratio $q=0.65\\pm 0.12$. From the secondary star spectral type, allowing for uncertainties in the mass-spectral type relation for CV donors, we find  $\\rm M_{WD} \\geq 0.5\\,M_{\\odot}$ and $\\rm M_{sec}=0.4-0.79\\,M_{\\odot}$ and a binary inclination  $25^o < i \\leq 58^o$.\\\\   \\item We estimate the distance in the range 540-840\\,pc which limits to only 3-4\\,$\\%$ the contribution of the WD to the observed flux at H$_{\\gamma}$ continuum and to 20-25\\,$\\%$ the contribution of the secondary star at 6100\\,$\\AA$. This implies that the observed absorption features underneath Balmer emissions are not due to the WD atmosphere. RX\\,J1548 is not accreting at a high rate which does  not favour an optically  thick disc as instead observed in high transfer rate systems such as  dwarf novae during outbursts or in SW Sex stars. The observed orbital and spin radial velocities of H$_{\\gamma}$ absorption instead suggest an origin  in the accretion flow close to the WD surface which needs further investigation via spectropolarimetric observations.\\\\  \\item The  spin-to-orbit period ratio $\\rm P_{\\omega}/P_{\\Omega}$=0.02 suggests  that this system is probably far from equilibrium. Also, the WD appears to be weakly magnetized with $\\rm B < $ 2\\,MG, much lower than typical values found in the other known soft X-ray/highly polarized IPs. Polarimetric observations are then needed to clarify whether this  system really contains a low field WD.  \\end{itemize}"
    },
    "0512/astro-ph0512194_arXiv.txt": {
        "abstract": "Observations of the centers of galaxies continue to evolve, and it is useful to take a fresh look at the constraints that exist on alternatives to supermassive black holes at their centers.  We discuss constraints complementary to those of Maoz (1998) and demonstrate that an extremely wide range of other possibilities can be excluded.  In particular, we present the new argument that for the velocity dispersions inferred for many galactic nuclei, even binaries made of point masses cannot stave off core collapse because hard binaries are so tight that they merge via emission of gravitational radiation before they can engage in three-body or four-body interactions. We also show that under these conditions core collapse  leads inevitably to runaway growth of a central black hole with a significant fraction of the initial mass, regardless of the masses of the individual stars.  For clusters of noninteracting low-mass objects (from low-mass stars to elementary particles), relaxation of stars and compact objects that pass inside the dark region will be accelerated by interactions with the dark mass.  If the dark region is instead a self-supported object such as a fermion ball, then if stellar-mass black holes exist they will collide with the object, settle, and consume it. The net result is that the keyhole through which alternatives to supermassive black holes must pass is substantially smaller and more contrived than it was even a few years ago. ",
        "introduction": "High-resolution observations of the nuclei of many galaxies have revealed large dark masses in small regions.  These are most naturally interpreted as supermassive black holes, but as emphasized by Maoz (1998) it is important to take stock of how rigorously we can rule out other possibilities.  Here we present arguments showing that under extremely general conditions almost all other options are ruled out, further emphasizing that supermassive black holes are by far the least exotic and most reasonable explanations for the data in many specific sources.  In \\S~2 we lay out our assumptions, making them as conservative as possible so that our conclusions are robust.  In \\S~3 we show that for many observed galactic nuclei, binaries are unable to heat the stellar distribution effectively because if they are hard then they merge quickly via gravitational radiation.  This important constraint, which depends only on dynamics and not the detailed properties of the specific objects, was not presented by Maoz (1998) or elsewhere as far as we are aware.  In \\S~4 we explore the consequences of core collapse and demonstrate that a very significant mass will inevitably coalesce even for point masses.  In \\S~5 we investigate for the first time the consequences if stellar-mass black holes exist outside the nucleus.  We show that enough of them will find their way to the center that they will have serious effects on the nuclear region, likely consuming a significant amount of mass and leading to a supermassive black hole.  We discuss the consequences of this analysis in \\S~6. ",
        "conclusions": "In the past decade, thanks to many observational developments, the case for supermassive black holes in the centers of many galaxies has gone from strong to essentially inescapable.   We have shown that for many specific galactic nuclei, the observational constraints are strong enough to rule out binary heating, hence the relevant evolution time is the time to core collapse.  This is a factor of $\\sim 20$ less than the time to evaporation, which has previously been used as the conservative standard for stellar cluster persistence. For many individual galactic nuclei, therefore, the combination of time to core collapse and lack of binary heating rules out dense stellar clusters as an alternate explanation for the inferred dark mass. Specifically, the Galaxy, NGC~4208, and M31 have core collapse times $<$2~Gyr for $0.5\\,M_\\odot$ objects and cannot be stabilised by binaries less than $100\\,M_\\odot$.  M32 also has a core collapse time $<$2~Gyr, but could in principle be stabilised by stellar-mass binaries. All other sources currently have core collapse times $>$200~Gyr for $0.5\\,M_\\odot$ objects.  The only remaining possibilities are concentrated regions of noninteracting low-mass particles or self-supported exotic objects such as a fermion balls \\citep{TV98}.  Even in this case, we have shown that dynamical evolution of the stars and black holes near the centers of galaxies will cause multiple stellar-mass black holes to fall to the center of the potential, if black holes exist at all.  Such black holes would consume any high-mass exotic pressure-supported objects, and would also accrete a noninteracting cluster of particles if allowed to move around freely. Therefore, the existence of stellar-mass black holes would lead to the production of supermassive black holes in many specific sources even if the supermassive holes did not form in other ways.  When combined with the high redshifts inferred from Fe~K$\\alpha$ lines in some Seyfert galaxies \\citep{RN03}, dramatic deviations from standard physics are required to explain observations in ways not involving black holes."
    },
    "0512/astro-ph0512641_arXiv.txt": {
        "abstract": "We find the solution for the scale factor in a flat Universe driven by dust plus a component characterized by a constant parameter of state which dominates in the asymptotic future. We also present an analytic formula (in terms of hypergeometric functions) for the past light cone in such a universe. As applications of this result, we give analytic expressions for the Luminosity Distance and for the Acoustic Scale, where the latter determines the peaks positions in the pattern of CMB anisotropies. ",
        "introduction": "\\setcounter{equation}{0} \\label{intro}  The past light cone, defined as the comoving distance $r$ travelled by a light signal from time $t$ to the present time $t_{0}\\,(>t)$, in cosmology is given by \\cite{KT} \\be r (t, t_0) = \\frac{1}{H_0 \\sqrt{|\\Omega_K|}} S_K \\left( H_0 \\sqrt{|\\Omega_K|} \\int_t^{t_{0}} \\frac{d t'}{a (t')} \\right) \\label{general} \\, , \\ee where $a(t)$ is the scale factor, $H_0$ is the Hubble parameter (defined as $H=d \\ln a/dt $) computed at $t_0$ and $\\Omega_K=-K/(H_0^2 a_0^2)$ is the present curvature energy density ($K=1, 0, -1$ for closed, flat, open spatial sections) with $a_0=a(t_0)$. Depending on $K=1\\,,0\\,,-1$ respectively, the function $S_K$: \\begin{displaymath} S_k (x)= \\left\\{ \\begin{array}{l} \\sin(x)  \\\\ x \\\\ \\sinh(x) \\, \\end{array} \\right. \\end{displaymath} contains information on the geometry of space.  For flat spatial sections Eq. (\\ref{general}) simplifies to \\be r (t, t_0) = \\int_t^{t_{0}} \\, d t'/ a(t') \\, . \\label{cd} \\ee Changing variable of integration from $t$ to the redshift $z$ (defined as $z=1/a - 1$, with $a(t_{0})=1$), one obtains \\be r (z) = \\int_0^{z} \\, d z' / H(z') \\, . \\label{cdz} \\ee Since $H(z)$ is given by the Friedmann equation (see Sec.\\ref{fried}), Eq.~(\\ref{cdz}) allows the computation of $r$ without solving for $a(t)$. Of course, once $r(z)$ is known, one still has to solve the Friedmann equation in order to obtain the dependence on time.  The integrations (\\ref{cd}),(\\ref{cdz}) are the starting points for the computation of some distances of astrophysical interest as the Proper Distance (or Horizon), the Angular Diameter Distance and the Luminosity Distance \\cite{peebles}. The comoving distance (\\ref{cdz}) is also needed for the computation of the Acoustic Scale that is the characteristic angular scale of the peaks in the Power Spectrum of the Cosmic Microwave Background (henceforth CMB) anisotropies.  An analytic form for $r$ is known for few simple cases \\cite{KT} because is difficult to perform the integration either in form of Eq.~(\\ref{cd}) or Eq.~(\\ref{cdz}). In the case of dust plus curvature, the time evolution for $a$ and $r$ are known \\cite{KT}.  Here we present the analytic solution of the scale factor {\\em and} the exact computation of the past light cone in a flat Universe driven by Dust (whose energy density $\\rho_C(t)\\propto a^{-3}$) and Dark Energy ($\\rho_X(t)\\propto a^{\\alpha}$) with a parameter of state (i.e. pressure over energy) constant in time. The important case of $\\Lambda$CDM is obtained setting $\\alpha =0$. We apply this formula to the following cases: first, we provide the exact expression of the Luminosity Distance (and we give an expansion of the past light cone at large $z$) and then we write down the analytic expression for the Acoustic Scale, relevant for the peaks analysis of the angular power spectrum of CMB anisotropies.  The paper is organized as follows: in section \\ref{fried} we solve the Friedmann equation introducing a new time coordinate. In section \\ref{plc} we find the analytic formula for the past light cone. We apply our results to the luminosity distance in section \\ref{ldistance} and to the acoustic scale in section \\ref{sectionas}. We conclude in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} We have found the analytic solution for the scale factor {\\em and} past light cone for a flat universe driven by dust and another $X$ component with a parameter of state constant in time which dominates in the future. If $w_X < -1/3$, $X$ is the simplest possible parametrization for Dark Energy. While Eq. (\\ref{solution}) generalizes the previously known solution for $\\Lambda$CDM \\cite{staro} to $\\alpha \\ne 0$, the result (\\ref{mainformula}) extends \\cite{DSW} for finite $z$, which is important for accurate quantitative predictions.   Eq.~(\\ref{mainformula}) has been applied to the following cases: the computation of the luminosity distance, relevant for supernovae, and the acoustic scale, which determines the position of the peaks in the pattern of CMB anisotropies. Both of these results are original. For the second case, our analytic prediction is in agreement with the value reported by the WMAP team. Note that this latter application requires a correct evaluation of the past light cone at high redshift, where a perturbative expansion does not work and only numerical methods were used so far.  A parameter of state constant in time cover a wide range of cosmological models which include dust. From a fundamental physics point of view, $X$ corresponds to a cosmic string network or a curvature term when $\\alpha=-2$ or to a domain walls background when $\\alpha=-1$. A quintessence model with an exponential potential which does not have tracking behaviour leads to a component with $ -2 < \\alpha < 0$ (for $\\alpha > 0$ one has to reverse sign of the kinetic term of a canonical scalar field, i.e. consider a phantom with an exponential potential)."
    },
    "0512/astro-ph0512655_arXiv.txt": {
        "abstract": "We consider the early cooling evolution of strongly magnetized strange stars in a CFL phase with high gap $\\Delta \\gsim 100$~MeV. We demonstrate how this model  may explain main features of the gamma-ray burst phenomena and also yield a strong star kick. The mechanism is  based on beaming of neutrino emission along the magnetic vortex lines. We show that for sufficiently high initial temperatures $T_0\\sim 30$ to $60~$MeV and surface magnetic fields $B_s \\sim 10^{15}$ to $10^{17}~$G, the energy release within the narrow beam is up to  $10^{52}$~erg with a magnetic field dependent time scale between $10^{-2}$~s (for a smaller magnetic field) to $10$~s. The above mechanism together with the parity violation of the neutrino-producing weak interaction  processes in a magnetic field allow for the strange star  kick. The higher the magnetic field the larger is the stars kick velocity. These velocities may cover the same range as observed pulsar kick velocities. ",
        "introduction": "Compact stellar objects are subdivided into three groups: usual neutron stars (having only hadronic matter interiors and an extended crust), hybrid stars (with quark core, hadronic shell and a crust similar to that in neutron stars) and strange stars (hypothetical self-bound objects composed of $u$, $d$ and $s$ quark matter with only a thin crust if any). The possibility of strange stars has been extensively discussed in the literature, see \\cite{Bombaci:2001uk} and Refs. therein.  Gamma ray bursts (GRB) are among  the most intriguing phenomena in the Universe, see \\cite{Piran} and Refs. therein. If the energy is emitted isotropically, an energy release  of the order of $10^{53}$ to $10^{54}$~erg is needed to power the GRB. However, there is now compelling evidence that the gamma ray emission is not isotropic, but displays a jet-like geometry. In the case of beaming a  smaller energy, of the order of $10^{51}$~erg, would be sufficient to power the GRB \\cite{Frayl}. Although typically one may speak about a short GRB with a time scale $t\\sim 10^{-3}$ to $ 10^{-1}~$s and a long GRB with a time scale $t\\sim ~1$ to $ 100$~ s, energies, durations, time scales involved in a burst, pulse shape structures, sub-burst numbers, etc, vary so much that it is hard to specify a typical GRB.  Many scenarios for the GRB have been proposed: the ``collapsar'', or ``hypernova'' model linking the GRB with ultra-bright type Ibc supernovae (hypernovae) and the subsequent black hole formation \\cite{BB}; neutron star mergers \\cite{ros}, or accretion of matter onto a black hole; strange star collisions \\cite{Haensel:1991um}; the model  \\cite{Usov} assuming a large surface magnetic field up to $ 10^{16}$~G of a millisecond pulsar produced in the collapse of a  $\\sim 10^9~$G white dwarf, with a powerful pulsar wind, as the source of the GRB; the model \\cite{SW1} suggesting a $ee^+$ plasma wind between heated neutron stars in close binary systems as consequence of the $\\nu\\bar{\\nu}$ annihilation; the model \\cite{RK} of a steadily accreting $\\sim 10^6~$G white dwarf collapse to a millisecond pulsar with a $\\sim 10^{17}$~G interior toroidal field, causing the GRB; an (isotropic) first order phase transition \\cite{Bombaci} of a pure hadronic compact star to a strange star (see also \\cite{Aguilera}); asymmetric core combustion \\cite{Lug} in neutron stars due to the influence of the magnetic field generating an acceleration of the flame in the polar direction, etc.  Pulsars are rotating neutron stars with rather high magnetic fields causing observable radio dipole signals. Most of the known pulsars are born in the neighborhood of the galactic plane and move away from it with natal kick velocities which are typically higher than those of their progenitors \\cite{Lyne94}. This implies that the birth process of pulsars also produces their high velocities and thus cannot be entirely isotropic \\cite{Spruit:1998sg}. Up to now, the mechanisms driving this asymmetry are far from clear. There are several hypotheses, ranging from asymmetric supernova explosions \\cite{lai01}, neutron star instabilities \\cite{colpi03,imshennik04} and magneto-rotational effects \\cite{moiseenko03,ardeljan04} to the model of an electromagnetic or neutrino rocket \\cite{lai01}. It is also not clarified whether the distribution of pulsar kick velocities is bimodal with a low-velocity component of $v \\leq 100 ~{\\rm km/s}$ (20 \\% of the known objects) and a high-velocity component of $v \\geq 500 ~{\\rm km/s}$ (80\\%) as suggested by \\cite{Arzoumanian:2002} or whether it can be explained by a one-component distribution \\cite{Hobbs:2005yx}. Most of the models are capable of explaining kick velocities of $v \\sim 100 ~{\\rm km/s}$, but it is a nontrivial problem to explain the highest measured pulsar velocities around $1600 ~{\\rm km/s}$.  In this work we assume  that a strongly magnetized strange star has been formed in the color superconducting color-flavor-locked (CFL) state with a large gap $\\Delta (T)$, $\\Delta (T=0)\\gsim 100$~MeV, cf. \\cite{Rajagopal:2000wf}, \\footnote{Our results are also relevant for the mCFL phase \\cite{IMTH}. We only need all quarks to be gapped with large gaps} as the result of a phase transition. The latter could be caused by the accretion of the matter from a companion star, the neutron star angular momentum decrease owing to the  gravitational and electromagnetic radiation after the supernova event had occurred, or during the collapse of a magnetized white dwarf to the neutron star state, the proto-neutron star collapse to the new stable state, or by some other reason. We also can deal with a hybrid star instead of the strange star. However, cf. \\cite{RK}, there is experimental evidence that the GRB  carries only a tiny baryon load of mass $\\lsim 10^{-4}M_{\\odot}$. Therefore the  hadronic shell and the crust should be rather thin or there should be a special reason for a low baryon loading. We suggest to explain the beaming by the presence of a strong magnetic field. {\\em Thus the beaming and the low baryon loading stimulate us to conjecture about a magnetized strange star,} which has a tiny  hadron shell, if any, and only a thin crust. Varying  the value of the surface magnetic field $B_s$, we search for an optimal configuration to explain the GRB characteristics  and to estimate a range of velocities of the star kicks. We use units $\\hbar =c=1$. ",
        "conclusions": "\\label{conclusion}  We have shown that in the presence of a strong magnetic field $B_{in,s} \\sim 10^{15} - 10^{17} $~G, the early cooling evolution of strange  stars with a color superconducting quark matter core in the CFL phase right after a short transport era is characterized by anisotropic neutrino emission, collimated within a beaming angle $\\theta_{\\nu}$ around the magnetic axis. Initial temperatures of the order of $T_0 \\sim 30 - 60~$MeV allow for the energy release $\\sim 10^{53}$~erg. An energy  $\\sim 10^{52}$~erg is released within a narrow beam on time scales from $10^{-2}$~s to $10$~s, appropriate for the phenomenology of both short and long GRB. Exploring the beaming mechanism one may obtain a wide range of strange star kick velocities up to $10^{3} ~{\\rm km/s}$, in dependence on the magnetic field, pairing gap, radius, mass and burst duration of the star. The range of strange star kick velocities is thereby in a qualitative agreement with recent observational data on pulsar kick velocity  distribution. All our estimates are very rough and essentially vary with the parameters of the model. Thus we just have shown a principal possibility of the application of the model to the GRB and kick description. A more systematic investigation, together with a consistent modeling of the compact star structure will be developed along the lines described in this contribution \\cite{Berdermann}.     \\subsection*{Acknowledgements} J.B. and D.B. acknowledge support by DAAD grant No. DE/04/27956 and by the DFG Graduate School 567 ``Strongly Correlated Many-Particle Systems'' at University of Rostock. The work of D.N.V. has been supported by DFG grant 436 RUS 113/558/0-3 and by an RF President grant NS-5898.2003.02. H.G has been supported by DFG under grant 436 ARM 17/4/05. D.B. is grateful for the invitation to participate at the School for Nuclear Physics in Erice and for the support by DFG."
    },
    "0512/astro-ph0512463_arXiv.txt": {
        "abstract": "We present measurements of the redshift-space luminosity-weighted or `marked' correlation function in the SDSS.  These are compared with a model in which the luminosity function and luminosity dependence of clustering are the same as that observed, and in which the form of the luminosity-weighted correlation function is entirely a consequence of the fact that massive halos populate dense regions. We do this by using mock catalogs which are constrained to reproduce the observed luminosity function and the luminosity dependence of clustering, as well as by using the language of the redshift-space halo-model. These analyses show that marked correlations may show a signal on large scales even if there are no large-scale physical effects---the statistical correlation between halos and their environment will produce a measureable signal. Our model is in good agreement with the measurements, indicating that the halo mass function in dense regions is top-heavy; the correlation between halo mass and large scale environment is the primary driver for correlations between galaxy properties and environment; and the luminosity of the central galaxy in a halo is different from (in general, brighter than) that of the other objects in the halo. Thus our measurement provides strong evidence for the accuracy of these three standard assumptions of galaxy formation models.  These assumptions also form the basis of current halo-model based interpretations of galaxy clustering.  When the same galaxies are weighted by their $u-$, $g-$, or $r-$band luminosities, then the marked correlation function is stronger in the redder bands.  When the weight is galaxy color rather than luminosity, then the data suggest that close pairs of galaxies tend to have redder colors.  This wavelength dependence of marked correlations is in qualitative agreement with galaxy formation models, and reflects the fact that the mean luminosity of galaxies in a halo depends more strongly on halo mass in the $r-$band than in $u$.  The $u-$band luminosity is a tracer of star formation, so our measurement suggests that the correlation between star formation rate and halo mass is not monotonic.  In particular, the luminosity and color dependence we find are consistent with models in which the galaxy population in clusters is more massive and has a lower star formation rate than does the population in the field.  The virtue of this measurement of environmental trends is that it does not require classification of galaxies into field, group and cluster environments. ",
        "introduction": "In hierarchical models of structure formation, there is a correlation between halo formation and abundances and the surrounding large scale structure---the mass function in dense regions is top-heavy (Mo \\& White 1996; Sheth \\& Tormen 2002). Galaxy formation models assume that the properties of a galaxy are determined entirely by the mass and formation history of the dark matter halo within which it formed. Thus, the correlation between halo properties and environment induces a correlation between galaxy properties and environment. The main goal of the present work is to test if this statistical correlation accounts for most of the observed trends between luminosity and environment (luminous galaxies are more strongly clustered), or if other physical effects also matter.  We do so by using the statistics of marked correlation functions (Stoyan \\& Stoyan 1994; Beisbart \\& Kerscher 2002) which have been shown to provide sensitive probes of environmental effects (Sheth \\& Tormen 2004; Sheth, Connolly \\& Skibba 2005). The halo model (see Cooray \\& Sheth 2002 for a review) is the language currently used to interpret measurements of galaxy clustering.  Sheth (2005) develops the formalism for including marked correlations in the halo model of clustering, and Skibba \\& Sheth (2005) extend this to describe measurements made in redshift space.  This halo model provides an analytic description of marked statistics when correlations with environment arise entirely because of the statistical effect.  Section~\\ref{mocks} describes how to construct a mock galaxy catalog in which the luminosity function and the luminosity dependence of clustering are the same as those observed in the Sloan Digital Sky Survey.  In these mock catalogs, any correlation with environment is {\\em entirely} due to the statistical effect. Section~\\ref{model} shows that the halo model description of marked statistics provides a good description of this effect, both in real and in redshift space. Section~\\ref{sdss} compares measurements of marked statistics in the SDSS with the halo model prediction.  The comparison provides a test of the assumption that correlations with environment arise entirely because of the statistical effect.  A final section summarizes our results, and shows that marked statistics provide interesting information about the correlation between galaxies and their environments without having to separate the population into the two traditional extremes of `cluster' and `field'. ",
        "conclusions": "We showed how to generate a mock galaxy catalog which has the same luminosity function (Figure~\\ref{lf}) and luminosity dependent two-point correlation function as the SDSS data. We used the mock catalog to calculate the luminosity-weighted correlation function in a model where all environmental effects are a consequence of the correlation between halo mass and environment (Figures~\\ref{xirVLS} and~\\ref{xisVLS} show results in real and redshift space). We then showed how to describe this luminosity-weighted correlation function in the language of the redshift-space halo model (equation~\\ref{Wk1h2h}). The analysis showed that estimates of the luminosity dependence of clustering constrain how the mass-to-light ratio of halos depends on halo mass (equation~\\ref{m2l} and Figure~\\ref{compareML}). The central galaxy in a halo is predicted to be substantially brighter than the other objects in the halo, and although the luminosity of the central object increases rapidly with halo mass, the mean luminosity of the other objects in the halo is approximately independent of the mass of the host halo.  Our analysis also showed that measurements of clustering as a function of luminosity completely determine the simplest halo model description of marked statistics.  In addition, measurements of the marked correlation function allow one to discriminate between models which treat the central object in a halo as special, from those which do not (Figures~\\ref{xirVLS} and~\\ref{xisVLS}). Also, in hierarchical galaxy formation models, the marked correlation function is expected to show a signal on large scales if the average mark of the galaxies in a halo correlates with halo mass. The signal arises because massive halos populate the densest regions; it is present even if there are no physical effects which operate to correlate the marks over large scales.  We compared this halo model of marked statistics with measurements in the SDSS (Figures~\\ref{xiSDSS} and~\\ref{projMp}). The agreement between the model and the measurements on scales smaller than a few Mpc provides strong evidence that central galaxies in halos are a special population---in general, the central galaxy in a halo is substantially brighter than the others. (Berlind et al. 2004 come to qualitatively similar conclusions, but from a very different approach.)  Substantially better agreement is found for a model in which the minimum halo mass required to host a luminous central galaxy does not change abruptly with luminosity. This is in qualitative agreement with some semi-analytic galaxy formation models, which generally predict some scatter in central luminosity at fixed halo mass (e.g. Sheth \\& Diaferio 2001; Zheng et al. 2005).  The agreement between the halo model calculation and the data on scales larger than a few Mpc indicates that the standard assumption in galaxy formation models, that halo mass is the primary driver of correlations between galaxy luminosity and environment, is accurate. In particular, these measurements are consistent with a model in which the halo mass function in large dense regions is top-heavy, and, on these large scales, there are no additional physical or statistical effects which affect the luminosities of galaxies. In this respect, our conclusions are similar to those of Mo et al. (2004), Kauffmann et al. (2004), Blanton et al. (2005), and Abbas \\& Sheth (2006), although our methods are very different.  We note in passing that there is a weak statistical effect for which the halo-model above does not account:  at fixed mass, haloes in dense regions form earlier (Sheth \\& Tormen 2004). Gao, Springel \\& White (2005) show that this effect is more pronounced for low mass haloes (related marked correlation function analyses by Harker et al 2006 and Wechsler et al. 2006 come to similar conclusions). The agreement between our halo-model calculation and the measurements in the SDSS suggests that this correlation between halo formation and environment is not important for the relatively bright galaxy population we have studied here.  This is presumably because these SDSS galaxies populate more massive haloes.  Comparisons with larger upcoming SDSS datasets, with a fainter luminosity threshold (such as $M_r<-18$), may bear out the correlation between low-mass halo formation and environment.  As a final indication of the information contained in measurements of marked statistics, Figure~\\ref{ugrMs} compares $M(s)$ when the $u-$, $g-$ and $r-$band luminosities are used as the mark.  The underlying population is the same as that for Figures~\\ref{rppiSDSS}--\\ref{projMp}:  the sample is volume limited to $M_r<-20.5$. Thus, $\\xi(s)$ is fixed, and only $W(s)$ changes with wave-band. Notice that there is a clear trend with wavelength:  there is a slight anti-correlation % in the $u-$band, whereas $M(s)$ rises slightly with decreasing scale when $g-$band luminosity is the mark, and the signal is even stronger when $L_r$ is the mark. This trend with wavelength is qualitatively consistent with the predictions of semi-analytic galaxy formation models (Sheth, Connolly \\& Skibba 2005) and indicates that the mean $u$-band luminosity of the galaxies in a halo depends less strongly on halo mass than does the mean $r-$band luminosity in a halo (Sheth 2005). In the models, the $u-$band luminosity is an indicator of the star formation rate; our measurements suggest that the correlation between star formation rate and halo mass is weak---if it is an increasing function of halo mass at low masses, then it decreases at larger masses.  \\begin{figure} \\centering \\includegraphics[width=\\hsize]{fig8.ps} % \\caption{Redshift-space luminosity-weighted correlation functions measured in volume limited catalogs with $M_r<-20.5$ in the SDSS. Circles, triangles and squares show $M(s)$ when the weight is $r-$, $g-$ and $u-$band luminosity respectively. For clarity, jack-knife error bars are only shown for the $r-$band measurement, since the uncertainties are similar in the other bands.} \\label{ugrMs} \\end{figure}  Figure~\\ref{umrgmr} shows the result of weighting these same galaxies by their colors.  The top panel shows results where the weight is the difference in the absolute magnitudes, $M_u-M_r$ and $M_g-M_r$, whereas the weights in the bottom bottom panel were the ratios of the luminosities in two bands.  Comparison of the two panels shows the effect on $M(s)$ of rescaling the weights while preserving their rank-ordering---while there are quantitative differences, the results in both panels are qualitatively similar.  The $M(s)$ measurements shown in the bottom panel are more widely separated because the luminosity ratio involves $10^{\\rm color}$, which has the effect of weighting the redder galaxies more heavily. In particular, this analysis indicates clearly that close pairs of galaxies tend to be redder than average. Sheth, Connolly \\& Skibba (2005) show that this is also the case in semi-analytic galaxy formation models.  The measurements shown in Figures~\\ref{ugrMs} and~\\ref{umrgmr} are consistent with models in which galaxies in clusters are more massive and have smaller star formation rates than galaxies in the field. In effect, these figures demonstrate the environmental dependence of galaxy luminosities and colors, without having to divide the galaxy sample up into discrete bins of `field', `group', and `cluster'. Thus, marked statistics allow one to study correlations with environment over a continuous range in density, rather than in somewhat arbitrary discrete bins in environment.  In this respect, our use of marked statistics to quantify and interpret environmental trends is very different from recent approaches which address the same problem. Since marked statistics are simple to measure and interpret, we hope that they will become standard tools for quantifying the correlation between the properties of galaxies and their environments.  \\begin{figure} \\centering \\includegraphics[width=\\hsize]{fig9.ps} % \\caption{Redshift-space color-weighted correlation functions measured in volume limited catalogs with $M_r<-20.5$ in the SDSS.  The top panel shows results when the color weight is the difference in absolute magnitudes.  In the bottom panel, galaxies were weighted by the ratio of the luminosities in the two bands, so they span a greater range around the mean value. Both panels show that close pairs of galaxies tend to have redder colors, although the difference is clearer when the weights span a greater range around the mean value.} \\label{umrgmr} \\end{figure}"
    },
    "0512/astro-ph0512180_arXiv.txt": {
        "abstract": "An observational approach is presented to constrain the global structure and evolution of the intracluster medium based on the {\\it ROSAT} and {\\it ASCA} distant cluster sample. From statistical analysis of the gas density profile and the connection to the $L_{\\rm X}-T$ relation under the $\\beta$-model, the scaled gas profile is found to be nearly universal for the outer region and the luminosity evaluated outside $0.2r_{500}$ is tightly related to the temperature through $\\propto T^{\\sim 3}$ rather than $T^{2}$. On the other hand, a large density scatter exists in the core region and there is clearly a deviation from the self-similar scaling for clusters with a small core size. A direct link between the core size and the radiative cooling timescale, $t_{\\rm cool}$ and the analysis of X-ray fundamental plane suggest that $t_{\\rm cool}$ is a parameter to control the gas structure and the appearance of small cores in regular clusters may be much connected with the thermal evolution. We derive the luminosity-`ambient temperature' relation ($L_{\\rm X}-T'$), assuming the universal temperature profile for the clusters with short cooling time and adopting a correction $T'=1.3T$, and find the dispersion around the relation significantly decreases in comparison to the case of the $L_{\\rm X}-T$ and the slope becomes less steep from $3.01^{+0.49}_{-0.44}$ to $2.80^{+0.28}_{-0.24}$. $L_{\\rm 1keV}$, which is defined as a normalization factor for each cluster, can be regarded as constant for a wide range of the cooling time.  We further examined the $L_{\\rm X}-T\\beta$ and $L_{\\rm X}-T'\\beta$ relations and showed a trend that merging clusters segregate from the regular clusters on the planes. A good correlation between the cooling time and the X-ray morphology on the $L_{\\rm 1keV}-(t_{\\rm cool}/t_{\\rm age})$ plane leads us to define three phases according to the different level of cooling, and draw a phenomenological picture: after a cluster collapses and $t_{\\rm cool}$ falls below the age of the universe, the core cools radiatively with quasi-hydrostatic balancing in the gravitational potential, and the central density gradually becomes higher to evolve from an outer-core-dominant cluster, which marginally follows the self-similarity, to inner-core-dominant cluster. ",
        "introduction": "The X-ray luminosity-temperature ($L_{\\rm X}-T$) relation of galaxy clusters is one of the most fundamental parameter correlations, established from previous X-ray observations \\citep[e.g.][]{Edge_etal_1990}. Since the X-ray luminosity reflects temperature and density profiles of hot intracluster medium (ICM), the $L_{\\rm X}-T$ relation should contain information on physical status and evolution of the ICM.  Observationally, the correlation is well approximated with a power-law function: $L_{\\rm X} \\propto T^{\\alpha}$, with $\\alpha\\sim3$ \\citep[e.g.][]{Edge_etal_1990,David_etal_1993, Markevitch_1998, Arnaud_Evrard_1999}.  In addition, it shows a significant scatter around the mean power-law relation and little redshift evolution. On the other hand, the self-similar model \\citep[e.g.][]{Kaiser_1986} predicts $\\alpha=2$. Thus the inconsistency between the observations and the simple theoretical model has been debated for many years and various possibilities including non-gravitational heating \\citep[e.g.][]{Evrard_Henry_1991, Cavaliere_etal_1997} and dependence of gas mass or gas-mass fraction on the temperature have been proposed \\citep[e.g.][]{David_etal_1993,Neumann_Arnaud_2001}.  Recently, hydrodynamical simulations have indicated that the effect of radiative cooling plays an important role to reproduce the observed $L_{\\rm X}-T$ relation \\citep[e.g.][]{Muanwong_etal_2002}.  They also suggested a further requirement of significant non-gravitational heating mechanism so as to account for the observed gas-mass fraction, however, the physical origin of the additional heating is yet to be understood.  Under the virial theorem and the isothermal $\\beta$-model \\citep{Cavaliere_Fusco_1976}, an ICM temperature $T_{\\rm gas}$ is proportional to a virial temperature, $T_{\\rm vir}$, and differs by a factor of $\\beta$, $T_{\\rm vir}\\sim T_{\\rm gas}\\beta$.  However, since $T_{\\rm gas}$ is usually an emission-weighted temperature measured from X-ray spectroscopy and strongly reflects the temperature of the cluster core region, a temperature decrease due to radiative cooling may have much influence on the $T_{\\rm gas}$ measurement.  The {\\it ASCA} spectroscopy of central region of clusters particularly with cD galaxy provided important results that cool emission from the core is systematically less luminous than that estimated in the past mainly on the basis of imaging observations \\citep[e.g.][]{Makishima_etal_2001}.  After the advent of the {\\it XMM-Newton} and {\\it Chandra} X-ray satellites, more detailed temperature profiles have been measured for nearby `cooling-flow' clusters, which typically show central temperature decrement by a factor of $\\sim 2$ \\citep[e.g.][]{Kaastra_etal_2004}.  These facts gave a piece of evidence against the standard cooling-flow model \\citep{Fabian_1994}, which triggered explorations of a variety of new scenarios for gas heating: heat conduction from outer hot layers \\citep[e.g.][]{Narayan_Medvedev_2001}, AGN heating \\citep[e.g.][]{Bohringer_etal_2002}, magnetic reconnections \\citep{Makishima_1997}, Tsunami \\citep{Fujita_etal_2004} etc. Recently \\cite{Masai_Kitayama_2004} proposed a quasi-hydrostatic model for cluster gas under radiative cooling, which predicts a characteristic temperature profile with an asymptotic temperature for the central region being $\\sim1/3$ of the non-cooling outer region. Their calculation agrees well with that observed in the cooling-flow clusters.  If the `universal' temperature profile emerges in all or a fraction of clusters with short cooling timescales, the effect of temperature drop should be properly taken into account in the discussion of the $L_{\\rm X}-T$ relation.  Another point is due to a connection of the $L_{\\rm X}-T$ relation to the cluster core sizes.  \\cite{Ota_Mitsuda_2002} found based on the X-ray data analysis of 79 distant clusters with {\\it ROSAT} and {\\it ASCA} that the distribution of the cluster core radius shows two distinct peaks at $\\sim 50$~kpc and $\\sim 200$~kpc for $H_0=70$. When dividing the sample into two subgroups corresponding to the two peaks, they show a significant difference in the normalization factor of the $L_{\\rm X}-T$ relation. Some possibilities have been discussed to understand the origin of the double-peaked distribution \\citep{Ota_Mitsuda_2002, Ota_Mitsuda_2004}, however, it is yet to be clarified. As long as one rely on the hydrostatic equilibrium and the isothermal $\\beta$-model, their result may indicate that underlying dark matter distribution is likely to have two preferable scales of 50 kpc and 200 kpc.  On the other hand, it is also found that the core radius is tightly related to the radiative cooling timescale \\citep[e.g.][]{Ota_Mitsuda_2004, Akahori_Masai_2005}. Thus the core structure might largely be a result of the thermal evolution of the ICM.  In this paper, using a large number of {\\it ROSAT} and {\\it ASCA} clusters compiled by \\cite{Ota_2001,Ota_Mitsuda_2004}, we investigate correlations between the fundamental cluster parameters.  Since {\\it ASCA} has a high sensitivity to measure the X-ray spectrum in the wide energy band while {\\it ROSAT} is good at imaging in the soft X-ray band, the two observatories are an excellent combination to study properties of the intracluster medium.  At present the {\\it XMM-Newton} and the {\\it Chandra} satellites are in orbit and generate much cluster data with higher sensitivities. However, the data set used in the present paper will be one of the best existing to construct the largest sample of distant clusters and study global X-ray structures.  While \\cite{Ota_Mitsuda_2002} mainly focused on the discovery of the typical two core sizes based on the analysis of 79 distant clusters with {\\it ROSAT} and {\\it ASCA}, \\cite{Ota_Mitsuda_2004} described the uniform analysis, and showed the X-ray parameters for the individual clusters and the results on scaling relations. The present paper addresses the $L_{\\rm X}-T$ relation in more detail, which is organized as below. In \\S~\\ref{sec:sample}, the sample and parameter estimation are described. In \\S~\\ref{sec:lx-t_density}, we show gas density profiles of the sample and how they are related to a scatter about the $L_{\\rm X}-T$ relation. Since redshift evolution is not clearly seen in the present distant sample and our analysis is based on a conventional $\\beta$-model, a study of redshift evolution in the $L_{\\rm X}-T$ relation is beyond the scope of the present paper.  In \\S~\\ref{sec:lx-t-rc} and \\S~\\ref{sec:lx-t-tcool}, we mainly focus on the $L_{\\rm X}-T$ relation and show its connection with the cluster core radius and the cooling time.  In \\S~\\ref{sec:lx-tbeta} we also derive a $L_{\\rm X}-T\\beta$ relation and show a trend that possible merging clusters segregate from regular clusters on the plane. Finally in \\S~\\ref{sec:discussion}, major results are summarized and the evolution of the ICM are discussed particularly in the light of radiative cooling.  We use $\\Omega_{\\rm M} = 0.3$, $\\Omega_{\\Lambda}=0.7$ and $h_{70}\\equiv H_0/(70~{\\rm km\\,s^{-1}Mpc^{-1}})=1$. The quoted parameter errors are the 90\\% confidence range throughout the paper unless otherwise noted.  The $1\\sigma$ error bars are plotted in all figures. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Summary of the Results} From the statistical analysis of 69 distant clusters in $0.1<z<0.56$ with the {\\it ROSAT} and {\\it ASCA} X-ray catalog \\citep{Ota_Mitsuda_2004}, we have obtained the following observational results.  \\begin{enumerate}  \\item Regarding the cluster outer region, the density profiles are nearly universal and the scatter around the $L_{\\rm X}(>0.2r_{500})-T$ relation is surprisingly small.  However, the slope of the relation is significantly larger than the self-similar expectation of 2.  \\item For the central region, the gas density exhibits a significant scatter and the self-similar condition is not satisfied particularly in the small core clusters.  \\item We studied the redshift evolution of the $L_{\\rm X}-T$ relation. If restricted the luminosity integration area to the central $0.2r_{500}$ region, there might be a weak, negative evolution, however, the normalization factor of the relation does not significantly change with redshift and is consistent with the nearby sample within the data scatter.  \\item We investigated the parameter correlations focusing on the $L_{\\rm X}-T$ relation and the connections to $r_c$ and $t_{\\rm cool}$, and suggested based on the X-ray fundamental plane analysis that $t_{\\rm cool}$ is likely to be a control parameter for the ICM structure of the cluster core region.  \\item For all small core clusters, $t_{\\rm cool}<t_{\\rm age}$ and the central electron density is systematically higher as $n_{e0}\\propto r_c^{-1.9}$, which is suggestive of a central concentration of the gas as cooling progresses. On the other hand, we confirmed that the temperature decrease is mild in the small core clusters; it is estimated as $T=1.3T'$ using the universal temperature profile. Then we derived the luminosity-`ambient temperature' relation and showed that the dispersion of $L_{\\rm 1keV}$ significantly decreases and can be regarded as constant against $t_{\\rm cool}$ within the data scatter.  \\item From the $L_{\\rm X}-T\\beta$ and $L_{\\rm X}-T'\\beta$ relations, we showed that $L_{\\rm 1keV}$ is related to the X-ray morphology, and that there is a marginal trend for the merging clusters segregating from the regular clusters on their planes. We indicated the possible three phases of the gas property along the $t_{\\rm cool}$ axis in Fig.~\\ref{fig14}d, according to the different level of cooling and ICM density structure. \\end{enumerate}  \\subsection{Impact of Radiative Cooling on the $L_{\\rm X}-T$ Relation and Comparison to the Quasi-hydrostatic Model} \\label{subsec:discuss_quasi_static}  In the context of isobaric cooling flow model, the local density of the gas increases against decreasing the temperature so that the thermal pressure $P(r)$ at given $r$ is kept constant within the cooling radius, i.e. $P(r) = n_e(r) kT(r) = \\mbox{const.}$, and then $n_e\\propto T^{-1}$ would follow.  If this is the case, since $L_{\\rm X} \\propto n_e^2 T^{1/2}$ for the gas hotter than $kT\\sim2$~keV, the $L_X-T$ relation would exhibit $L_{\\rm X}\\propto T^{-3/2}$.  For a typical cooling radius of $r_{\\rm cool}\\sim 0.1r_{500}$ ($\\sim 100$~kpc), $L_{\\rm X}(<r_{\\rm cool})$ contributes about 50\\% of the total luminosity. When we derived the $L_{\\rm X}(<0.1r_{\\rm 500})-T$ relation, however, we did not find any tendency for such an inverse correlation (see also the $L_{\\rm X}(<0.2r_{500})-T$ relation in Fig.~\\ref{fig4}a).  The observational results indicate that the gas in the core region of small core clusters is undergoing radiative cooling but the temperature decline towards the center seems mild.  The gas may be considered close to a steady state.  We will discuss such a possible state of the gas, quasi-hydrostatic cooling, proposed by \\cite{Masai_Kitayama_2004}.  The model predicts a moderate and smooth gas inflow with hydrostatic balancing.  In the context, unlike isobaric cooling flows that increase the local density so $P(r)$ is kept against local cooling, quasi-hydrostatic cooling allows the gas to modify its profile or core size so $\\nabla P(r)$ matches with the force by gravitational potential.  In \\S~\\ref{subsec:lx-t-tcool} we derived the $L_{\\rm X}-T'$ relation by correcting the emission-weighted temperature for small-core clusters with short $t_{\\rm cool}$ based on the idea of the universal temperature profile, and showed that the bolometric luminosity normalized by the temperature dependence of $(kT')^{2.8}$, $L_{\\rm 1keV}$ is constant within the data scatter over a wide range of $t_{\\rm cool}/t_{\\rm age}$.  This strongly suggests that the rate of thermal energy loss is kept nearly constant after the onset of cooling in the cluster core.  According to the quasi-hydrostatic model, the temperature starts to decrease at the cooling radius and approaches a constant of $\\sim 1/3$ the ambient temperature towards the center.  The temperature profile does not depend on its absolute value.  This picture may explain the observations of nearby cooling-flow clusters.  The mass inflow rate from the outer region is controlled by cooling so as to maintain the quasi-hydrostatic balance and is expected to vary little through the flow.  The condition is well satisfied in the case that the temperature is $kT \\ga 2$~keV where bremsstrahlung dominates cooling more than line emission.  The present sample meets this condition since it has a temperature range of $2<kT~{\\rm [keV]}<12$.  Therefore our results can be consistently understood within a framework of the quasi-hydrostatic cooling model.  At the same time, we have found that the gas density profile for the outer part of the cluster does not vary very much from cluster to cluster, and is almost independent of the gas cooling at the center. This would also imply that the energy loss due to radiative cooling at the center is only a portion of the total energy of the thermal gas, and the outer region is likely to offer a great {\\it reservoir} of the heat, which is carried a fraction into the inner cooling region advectively by the gas inflow under quasi-hydrostatic balance.  As described above, cooling with quasi-hydrostatic balancing modifies the core structure of the gas.  This may give a clue to the origin of the small core of $\\sim 50$~kpc or the inner core component of the double-$\\beta$ cluster.  In addition the study would also have an impact on the present understanding of the origin of the dispersions around the $L_{\\rm X}-T$ relation because it provided the observational evidence that they are directly linked to the thermal evolution of the ICM and showed how the relation would be modified when the effect of the cooling is considered.   \\subsection{Evolution of Density Structure of X-ray Clusters} \\label{subsec:discuss_evolution}  In this subsection we attempt to present an overall picture of the ICM evolution on the basis of the observational results.  First we consider about the $r_c$ distribution and the origin of the double-$\\beta$ nature of the clusters from the phenomenological point of views.  We have reported that the cluster core radius has a large cluster-to-cluster dispersion of more than an order of magnitude, even if restricted the sample to the regular clusters \\citep{Ota_2001}. This is partly due to the presence of two peaks separated by a factor of $\\sim4$. For 9 (7) out of the 35 regular clusters, the surface brightness is significantly (marginally) better fitted with the double-$\\beta$ model, and the average values of two core radii for the double-$\\beta$ clusters coincide well with the two $r_c$ peaks of the single-$\\beta$ clusters.  Then a question may arise: are there three subclasses of clusters, small-core single-$\\beta$, large core single-$\\beta$, and double-$\\beta$ clusters, or are we looking at some different aspects of a single evolutionary process?  From the imaging analysis, we noticed that the double-$\\beta$ clusters are further subdivided into inner-core dominant and outer-core dominant ones; if we are to perform a single $\\beta$ model fitting for these clusters, we would pick up either of the core radii as the best-fit result.  Thus the similarity in the $r_c$ distributions for the single-$\\beta$ and double-$\\beta$ clusters can be interpreted in the following way: the (regular) clusters have more or less double-$\\beta$ nature in their structure, and when X-ray emission from one of the two cores dominates the whole cluster emission, we recognize it as the single-$\\beta$ case and pick up either the inner or outer core as the core radius. This will account for an apparent large dispersion of core radius. Accordingly the single-$\\beta$ clusters with a small/large core radius may respectively correspond to extreme cases of inner/outer core dominant double-$\\beta$ clusters.  Then why are there inner-core dominant and outer-core dominant clusters?  Such double structures have often been attributed to the cooling-flow phenomenon at the cluster center. However, the present analysis showed that the standard cooling-flow model with dynamic inflow is not preferable to account for the results, in particular the $T-n_{e0}$ and $L_{\\rm 1keV}-t_{\\rm cool}/t_{\\rm age}$ relations.  On the other hand, in \\S~\\ref{sec:lx-tbeta} we suggested the possible three phases of the ICM property on the $L_{\\rm 1keV}-t_{\\rm cool}/t_{\\rm age}$ plane in the light of the radiative cooling and the relevance to a variety of density structure.  It is likely that the double-$\\beta$ clusters are those in the transient phase of their cooling cores, i.e., core size transition from outer-core dominated to inner-core dominated, mentioned in the quasi-hydrostatic cooling model.  As the interpretation of the double core nature, we consider in the following way: when some time passed after the collapse and the cooling becomes important, the core is radiatively cooled and the central density gradually becomes higher to evolve from an outer-core-dominant cluster toward an inner-core-dominant cluster.  The above discussion on the evolution of gas structure focused on the effect of radiative cooling and the possibility of quasi-hydrostatic state since the X-ray analysis showed that a majority of small core clusters are distributed in a narrow range of $L_{\\rm 1keV}$ once the temperature bias is corrected (Fig.~\\ref{fig12}) and thus supports the existence of a typical value. Moreover, a long tail at the smaller $L_{\\rm 1keV}$ is seen in the histogram for $r_c < 100$~kpc (Fig.~\\ref{fig12}c). This indicates existence of some peculiar clusters or extra physical mechanism that produces such a variation. Recently, \\cite{O'hara_etal_2005} examined a relation between substructure and cool cores and the scatter about cluster scaling relations using 45 nearby clusters. They found that cool core clusters, which usually have less morphological substructure, exhibit higher intrinsic scatter about scaling relations even after correcting for cooling core effect, suggesting a possibility that a more global process is at work. Then the variation in the $L_{\\rm 1keV}$ histogram for the small $r_c$ clusters may be related with the scatter found by \\cite{O'hara_etal_2005}. To further clarify both the structural evolution and the scatter of scaling relations, a systematic study with higher sensitivities will be important.  \\subsection{Related Issues and Future Observations}  We will remark on the effect of the cold front detected in nearby clusters with the {\\it Chandra} observations in regard to the present interpretation, and the assumption of the universal temperature profile with consideration on the recent progress of {\\it Chandra} and {\\it XMM-Newton} observations.  The cold front structure is characterized by a compact core with the density and temperature jumps in contrast to the surroundings, and interpreted as a remnant of subcluster merger \\citep{Markevitch_etal_2000}.  For higher redshifts, it is more difficult to directly identify such small structures due to the limitation of the sensitivity to measure the temperature maps, and some of them might possibly be fitted with the double-$\\beta$ model. However, as the cold front is a signature of recent merger, it is expected to be accompanied with the irregular X-ray morphology. Thus for relatively bright clusters, a small core of the cold front in an irregular cluster may be distinguishable from that of a regular cluster by referring to the morphological information.  For example, \\object{A115} and \\object{1E0657-56}, for which a possible remnant of cooling core of a merging subcluster was pointed out \\citep{Gutierrez_Krawczynski_2005, Markevitch_etal_2002}, were classified into irregular clusters from our image analysis, though they have different $r_c$ and are located in the different phases in Fig.~\\ref{fig14}d ( i) and iii), respectively).  For more accurate identification of the cold front in distant clusters, it is necessary to measure the 2 dimensional density and temperature maps with higher sensitivity observations.  We have utilized the universal temperature profile to estimate the effect of temperature decrease and relate the emission-weighted temperature to the ambient temperature for clusters with short cooling time and apparent luminosity enhancement in Fig.~\\ref{fig10}b. Due to the limited spatial resolution of {\\it ASCA}, the temperature profiles were not directly constrained for the distant sample. However for some of them, the results have been already published with the {\\it Chandra} and {\\it XMM-Newton} data. As for the double-$\\beta$ clusters or the small core cluster with very short cooling time of $-1.3 \\lesssim \\log{(t_{\\rm cool}/t_{\\rm age})} \\lesssim -0.7 $, \\object{A2204}, \\object{PKS0745-191}, \\object{A2390}, \\object{A1835}, the gradual temperature decline within the central region of roughly $r<100 - 200$~kpc were reported \\citep{Sanders_etal_2005, Hicks_etal_2002, Allen_etal_2001b, Majerowicz_etal_2002}. For the other double-$\\beta$ cluster, A1689, with slightly longer $\\log{t_{\\rm cool}/t_{\\rm age}}=-0.56$, the nearly isothermal profile was found with {\\it Chandra} \\citep{Xue_Wu_2002}, and a complex temperature structure was further reported with {\\it XMM-Newton} \\citep{Andersson_Madejski_2004}.  Therefore, from the overall standpoint, the assumption of the universal temperature profile for $\\log{t_{\\rm cool}/t_{\\rm age}} \\lesssim -0.5$ seems reasonable for many of the small core and double-$\\beta$ clusters.  To directly confirm the temperature structure in relation to the gas cooling, we need to analyze the temperature profile of individual clusters in various phases with using higher resolutions.  In \\S~\\ref{subsection:lx-t_outside}, we derived the $L_{\\rm X}(>0.2r_{500})-T$ relation without correcting for the temperature profile.  Adopting the corrected temperature $T' =1.3 T$ (Eq. \\ref{eq:temp_correction}), we obtain a similar slope, $L_{\\rm X}(>0.2r_{500})=2.34^{+2.35}_{-1.13}\\times10^{42} (kT')^{3.05^{+0.33}_{-0.35}}$ for 69 clusters ($\\chi^2/{\\rm d.o.f.}=610/67$) (Table~\\ref{tab1}) and a smaller standard deviation for $L_{\\rm 1keV}$, $\\sigma/\\mu=(4.5\\pm0.5)\\times10^{-3}$.  Once the ambient temperature $T'$ is directly measured for a large sample of clusters by future observations, we would suggest to use it instead of $T$ in the analysis of cluster scaling relations."
    },
    "0512/astro-ph0512149_arXiv.txt": {
        "abstract": "We report Swift Burst Alert Telescope (BAT) observations of the X-ray Flash (XRF) XRF 050416A.  The fluence ratio between the 15-25 keV and 25-50 keV energy bands of this event is 1.5, thus making it the softest gamma-ray burst (GRB) observed by BAT so far. The spectrum is well fitted by the Band function with $\\epo$ of 15.0$_{-2.7}^{+2.3}$ keV.  Assuming the redshift of the host galaxy (z $=$ 0.6535), the isotropic- equivalent radiated energy $\\eiso$ and the peak energy at the GRB rest frame ($\\eps$) of XRF 050416A are not only consistent with the correlation found by Amati et al. and extended to XRFs by Sakamoto et al., but also fill-in the gap of this relation around the 30 -- 80 keV range of $\\eps$.  This result tightens the validity of the $\\eps$ -- $\\eiso$ relation from XRFs to GRBs.  We also find that the jet break time estimated using the empirical relation between $\\eps$ and the collimation corrected energy $\\egamma$ is inconsistent with the afterglow observation by Swift X-ray Telescope.  This could be due to the extra external shock emission overlaid around the jet break time or to the non existence of a jet break feature for XRF, which might be a further challenging for GRB jet emission, models and XRF/GRB unification scenarios. ",
        "introduction": "The observations of X-ray flashes (XRF) are providing important information for understanding the nature of Gamma-Ray Bursts (GRB).  The detailed studies of XRFs started few years ago based on BeppoSAX observations \\citep{heise2000,kippen2002}, but X-ray rich events had already been detected by the $Ginga$ satellite.  \\citet{yoshida1989} reported that soft X-ray emission below 10 keV co-exists with $\\gamma$-ray emission of GRBs.  About 36\\% of the bright bursts observed by $Ginga$ have $\\epo$ energy, which is the photon energy at which the $\\nu$F$_{\\nu}$ spectrum peaks, around a few keV and also show large X-ray to $\\gamma$-ray fluence ratios \\citep{strohmayer1998}.  The Wide Field Cameras (WFC) on-board the $Beppo$SAX satellite observed 17 XRFs in five years \\citep{heise2000}. \\citet{kippen2002} searched for GRBs and XRFs which were observed in both WFC and BATSE. The WFC and BATSE joint spectral analysis of XRFs shows that their $\\epo$ energies are significantly lower than those of the BATSE $\\epo$ distribution \\citep{preece2000}. The systematic study of the spectral properties of XRFs observed by $HETE$-2 also supports this result \\citep{sakamoto2005}.  The afterglow detection and the redshift measurement from the host galaxy of XRF 020903, which is one of the softest XRF observed by $HETE$-2, shows the dramatic progress in understanding the nature of XRFs.  The prompt emission of XRF 020903 has $\\epo$ $<$ 5.0 keV which is two orders of magnitude smaller than that of typical GRBs.  The optical transient and the host galaxy of XRF 020903 were detected.  Further spectroscopic observation of the host galaxy suggests that the redshift is $0.25 \\pm 0.01$ \\citep{soderberg2004}.  \\citet{sakamoto2004} calculated the isotropic-equivalent energy $\\eiso$ and the peak energy at the source frame $\\eps$ using the redshift of the host galaxy, and found that XRF 020903 follows an extension of the empirical relationship between $\\eiso$ and $\\eps$ found by \\citet{amati2002} for GRBs (a.k.a. Amati relation). This result provides the observational evidence that XRFs and GRBs form a continuum and are a single phenomenon.  In this paper, we report the prompt emission properties of XRF 050416A observed by Burst Alert Telescope (BAT) on-board the $Swift$ satellite. The X-ray flash, XRF 050416A, was detected and localized by the $Swift$ \\citep{gehrels2004} Burst Alert Telescope (BAT; \\citet{barthelmy2005}) at 11:04:44.5 UTC on 2005 April 16 \\citep{sakamoto2005b,sakamoto2005c}. $Swift$ autonomously slewed to the BAT on-board position, and both $Swift$ X-Ray Telescope (XRT; \\citet{burrows2005}) and UV-Optical Telescope (UVOT; \\citet{roming2005}) detected the afterglow (Cusumano et al. (2005) in preparation, Holland et al. (2005) in preparation). The afterglow emission of XRF 050416A was also observed by ground observatories at various wavelengths \\citep{cenko2005a,anderson2005, li2005,kahharov2005,price2005,cenko2005b,soderberg2005}.   \\citet{cenko2005c} reported that the host galaxy is faint and blue with large amount of the star formation and its redshift is z = 0.6535 $\\pm$ 0.0002. Throughout this paper, the quoted errors are the 90\\% confidence level and the sky coordinates in J2000 unless we state otherwise in the text. ",
        "conclusions": "One of the most important discoveries related to XRF 050416A is the confirmation of the $\\eps$ -- $\\eiso$ relation \\citep{amati2002}. We calculate the $E_{\\rm peak}$ energy at the GRB rest frame, $\\eps$, and the isotropic-equivalent energy (1 -- 10$^{4}$ keV at the rest frame), $\\eiso$, using the redshift of the host galaxy (z=0.6535).  Assuming $\\alpha = -1$, $\\eps$ and $\\eiso$ of XRF 050416A are 25.1$_{-3.7}^{+4.4}$ keV and $(1.2 \\pm 0.2) \\times 10^{51}$ erg, respectively. Figure \\ref{fig:epeak_eiso} shows the data point of XRF 050416A with the known redshift GRBs of $Beppo$SAX and $HETE-2$ sample \\citep{amati2003,lamb2004, sakamoto2004}.  XRF 050416A not only follows the $\\eps$ $\\propto$ $\\eiso^{0.5}$ relation, but also fills in the gap of the relation around $\\eps$ of 30 -- 80 keV. This result tightens the validity of this relation at five orders of magnitude in $\\eiso$ and at three orders of magnitude in $\\eps$.  XRF 050416A bridges the gap between XRFs which have $\\eps$ of less than 10 keV and GRBs in the $\\eps$ -- $\\eiso$ relation.  The confirmation of $\\eps$ -- $\\eiso$ relation from XRFs to GRBs gives us a clear indication that XRFs and GRBs form a continuum and are a single phenomenon.  There are several jet models to explain a unified picture of XRFs and GRBs.   The off-axis jet model \\citep{yamazaki2004,toma2005}, the structured jet model \\citep{rossi2002,zhang2002,zhang2004}, and the variable jet opening angle model \\citep{lamb2005} are the most popular models in this aspect.  On the other hand, there are theoretical models to explain XRFs in the frame work of the internal shock model \\citep{mes2002,mochkovitch2003} and of the external shock model \\citep{dermer1999,huang2002,dermer2003}.  The cited jet models and internal/external shock models not only explain the existences of XRFs, under certain assumptions, but also, in some of their realizations or for some values of their parameters, they can predict the $\\eps$ -- $\\eiso$ correlation.  According to the XRT afterglow observation of XRF 050416A, the decay slope of the afterglow emission is $\\sim$ $-0.9$ from 0.015 days to $\\sim$ 34.7 days after the GRB trigger without any signature of a jet break (Cusumano et al. (2005) in preparation; \\citet{neusek2005}).  Using $\\eps$ and $\\eiso$ of XRF 050416A measured by BAT, we can estimate the jet break time using the relation between $\\eps$ and the jet collimation-corrected energy $E_{\\gamma}$ found by \\citet{ghirlanda2004} (Ghirlanda relation). However, there is a debate about the assumption of the jet model used by \\citet{ghirlanda2004} to derive the relationship between $\\eps$ and $E_{\\gamma}$ \\citep{xu2005,liang2005}. Based on this argument, we use the empirical relation between $\\eiso$, $\\eps$, and the jet break time at the rest frame, $t_{\\rm jet}^{\\rm src}$, derived by \\citet{liang2005}.  Note that there is no assumption of a jet model in the formula found by \\citet{liang2005}, and thus their relation is purely based on observational properties. When we use the equation (5) in \\citet{liang2005}, ($\\eiso/10^{52} \\,{\\rm erg}) = 0.85 \\times (\\eps/ 100 \\,{\\rm keV})^{1.94} \\times (t_{\\rm jet}^{\\rm src}/1 \\, {\\rm day})^{-1.24}$, the jet break time in the observer's frame is estimated to be $\\sim$ 1.5 days after the GRB on-set time. Note that this estimated jet break time is consistent with the estimation using the Ghirlanda relation assuming the circum-burst density of 3 cm$^{-3}$. Thus, the estimated jet break time using the empirical $\\eps$-$\\eiso$- $t_{\\rm jet}^{\\rm src}$ relation is inconsistent with the null detection of a jet break until more than 34.3 days after the trigger by XRT.  In the off-axis jet model \\citep{yamazaki2004,toma2005}, the null detection of the jet break in the XRT data of XRF 050416A could be difficult to explain. When we assume a bulk Lorentz factor of 100, $\\eps$ of 300 keV for an on-axis observer, and a jet opening angle of 2 degrees, the viewing angle from the jet on-axis is estimated to be $\\sim$ 4 degrees from the observed $\\eps$ of 25 keV.  According to \\citet{granot2002}, when observing the jet from an angle two times larger than the jet opening angle, we would expect to see a rise in the flux around one day after the burst.  It is possible to increase the bulk Lorentz factor and to reduce the off-axis viewing angle to achieve the same Doppler factor.  However, in this case, the afterglow light curve should be close to the on-axis case, thus, we would expect to see the jet break around the time we estimated.  On the other hand, the variable jet opening angle model \\citep{lamb2005} might work for XRF 050416A if $\\egamma$ is a constant value.  If we assume the values typical for GRBs ($\\eps$ = 300 keV and the jet opening angle of 5 degrees), the jet opening angle of XRF 050416A is calculated to be 52 degrees because of the inverse relation between $\\eps$ and the jet opening angle in the case in which $\\egamma$ is a constant. When we used the formulation of \\citet{sari1999} applying the estimated jet opening angle, the jet break time will be 64 days in the case of the circum-burst density of 10 cm$^{-3}$.  Both properties of the low $\\eps$ and the null detection of the jet break could be explained in the variable jet opening angle model if $\\egamma$ is constant.  However, as \\citet{ghirlanda2004} showed, $\\egamma$ is not a constant parameter, but has a good correlation with $\\eps$.  When we applied the Ghirlanda relation, $\\eps \\propto \\egamma^{0.7}$, in the variable opening angle model, and re-calculated the jet break time, the break time will be 0.7 days assuming the circum-burst density of 10 cm$^{-3}$.  In the variable jet opening angle model, there is no way to explain both the Ghirlanda relation and the null detection of the jet break by XRT simultaneously.  One natural way to explain the non-detection of the jet break feature is that extra components are overlaid around a jet break time period.  According to the afterglow calculations in the X-ray band by \\citet{zhang2005}, there are several possibilities to hide a jet break feature due to some kinds of emission by the external shock.  These are the external shock emission from 1) the dense clouds surrounding a GRB progenitor (e.g. \\citet{lazzati2002}), 2) a moderately relativistic cocoon component of a two-component jet (e.g. \\citet{granot2005}) , and 3) a jet with large fluctuations in angular direction (patchy jets; \\citet{{kumar2000}}).  On the other hand, it might be the case that XRFs indeed do not show the signature of a jet break in the afterglow. Indeed although the numbers in the sample are limited, there is no clear observational indication of a jet break in any XRF afterglow light curve so far. If the later case is true, we need to change our view of XRFs completely. Thus, the multi-wavelength observations of the XRF afterglows will be crucial to investigate whether a jet break feature exists in XRFs or not."
    },
    "0512/astro-ph0512239_arXiv.txt": {
        "abstract": "We report the discoveries of two, two-image gravitationally lensed quasars selected from the Sloan Digital Sky Survey: SDSS~J0806+2006 at $z_s=1.540$ and SDSS~J1353+1138 at $z_s=1.629$ with image separations of $\\Delta{\\theta}=1\\farcs40$ and $\\Delta{\\theta}=1\\farcs41$ respectively.  Spectroscopic and optical/near-infrared imaging follow-up observations show that the quasar images have identical redshifts and possess extended objects between the images that are likely to be lens galaxies at $z_l \\simeq 0.6$ in  SDSS~J0806+2006 and $z_l \\simeq 0.3$ in SDSS~J1353+1138.  The field of SDSS~J0806+2006 contains several nearby galaxies that may significantly perturb the system, and  SDSS~J1353+1138 has an extra component near its Einstein ring that is probably a foreground star. Simple mass models with reasonable parameters reproduce the quasar positions and fluxes of both systems. ",
        "introduction": "\\label{sec:intro}  Since the discovery of Q0957+561 \\citep*{walsh79}, about 80 gravitationally lensed quasars have been discovered \\citep{kochanek04}. Lensed quasars are not only intriguing phenomena but also have become indispensable astronomical tools, including probes of the cosmological parameters and the structure of galaxies \\citep[e.g.,][]{refsdal64,kochanek91}. In particular, the abundance of gravitational lenses in a well-defined source sample can be used to constrain dark energy \\citep[e.g.,][]{turner90,fukugita90,chae02}. Unfortunately, the largest existing survey, the Cosmic Lens All-Sky Survey \\citep[CLASS;][]{myers03,browne03}, contains only 22 lensed radio sources (with a well-defined statistical sample of 13 lenses) discovered from  ${\\sim}10,000$ radio sources, which is still insufficient to place tight constraints on dark energy models. The Sloan Digital Sky Survey \\citep[SDSS;][]{york00} should lead to a significantly larger lens sample for attacking the dark energy problem.  SDSS is expected to identify $10^5$ quasars spectroscopically \\citep[e.g.,][]{schneider05} and $\\approx10^6$ quasars photometrically \\citep[e.g.,][]{richards04}, which should lead to a sample of over $10^2$ lensed quasars given a standard lensing probability of $10^{-3}$ \\citep{turner84}.  Indeed, 12 new lensed quasars have been discovered from the SDSS quasars so far \\citep*{inada03a,inada03b,inada03c,inada05,johnston03,morgan03,pindor04, pindor05,oguri04,oguri05,burles05}, in addition to recovering 5 previously known lensed quasars \\citep{walsh79,weymann80,surdej87,bade97,oscoz97}. We can presently construct a well-defined statistical sample of 16 lensed SDSS quasars, but there remain many promising SDSS lensed quasar candidates for which the necessary follow-up observations are incomplete.  In this paper, we report on the discovery of two more gravitationally lensed quasars, SDSS~J080623.70+200631.9 (SDSS~J0806+2006) and SDSS~J135306.35+113804.7 (SDSS~J1353+1138). We present imaging and spectroscopic follow-up observations with the University of Hawaii 2.2-meter (UH88) telescope, the W. M. Keck Observatory's Keck I and II telescopes, and the Magellan Consortium's Landon Clay 6.5-m (LC6.5m) telescope \\footnote{The second telescope of the Magellan Project; a collaboration between the Observatories of the Carnegie Institution of Washington (OCIW), University of Arizona, Harvard University, University of Michigan, and Massachusetts Institute of Technology (MIT) to construct two 6.5 Meter optical telescopes in the southern hemisphere.}. We model the systems to check that their geometries are consistent with the lensing hypothesis.  The structure of this paper is as follows. We describe our lens candidate selections from the SDSS data in \\S \\ref{sec:sdss}. The follow-up observations and mass modeling of SDSS~J0806+2006 and  SDSS~J1353+1138 are presented in \\S \\ref{sec:0806} and \\S \\ref{sec:1353}, respectively. We finally present a summary and give a conclusion in \\S \\ref{sec:conc}.  Throughout the paper we assume a cosmological model with the matter density $\\Omega_M=0.27$, cosmological constant $\\Omega_\\Lambda=0.73$, and Hubble constant $h=H_0/100{\\rm km\\,sec^{-1}Mpc^{-1}}=0.7$ \\citep{spergel03}. ",
        "conclusions": "\\label{sec:conc}  We report the discovery of two doubly-imaged quasar lenses, SDSS~J0806+2006 and SDSS~1353+1138.  Both were selected from the SDSS spectroscopic quasar sample as lensed quasar candidates and confirmed in subsequent imaging and spectroscopic observations. SDSS~J0806+2006 consists of two $z_s=1.540$ quasar images separated by $\\Delta{\\theta}=1\\farcs40$ lensed by a galaxy at $z_l\\simeq0.6$. The lens galaxy is closer to the fainter image as expected, and its redshift, as suggested by its magnitude, colors, and the presence of a \\ion{Mg}{2} absorption feature, is $z_l=0.573$. Several nearby galaxies may perturb this system and indicate that the lens galaxy is part of a small group. SDSS~1353+1138 consists of two $z_s=1.629$ quasar images separated  by $\\Delta{\\theta}=1\\farcs41$ with a lens galaxy at $z_l{\\sim}0.3$. The redshift of the lens galaxy is estimated based on its magnitude, colors, and the spectral  flux ratio between the two quasar images. There is an additional component, which we have labeled C, superposed on this system whose nature is presently unexplained. Observations using the {\\em Hubble Space Telescope} are probably required to clarify its role in the lensed quasar system."
    },
    "0512/gr-qc0512081_arXiv.txt": {
        "abstract": "The emission spectrum from a simple accretion disk model around a compact object is compared for the cases of a black hole (BH) and a boson star (BS) playing the role of the central object. It was found in the past that such a spectrum presents a hardening at high frequencies; however, here it is shown that the self-interaction and compactness of the BS have the effect of softening the spectrum, the less compact the star is, the softer the emission spectrum at high frequencies. Because the mass of the boson fixes the mass of the star and the self-interaction the compactness of the star, we find that, for certain values of the BS parameters, it is possible to produce similar spectra to those generated when the central object is a BH. This result presents two important implications: (i) using this simple accretion model, a BS can supplant a BH in the role of compact object accreting matter, and (ii) within the assumptions of the present accretion disk model we do not find a prediction that could help distinguish a BH from a BS with appropriate parameters of mass and self-interaction. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512525_arXiv.txt": {
        "abstract": "In this paper, we present a first comparison of different Adaptive Optics (AO) concepts to reach a given scientific specification for 3D spectroscopy on Extremely Large Telescope (ELT). We consider that a range of $30\\%-50\\%$ of Ensquarred Energy (EE) in H band ($1.65$$\\mu$m) and in an aperture size from $25$ to $100$mas is representative of the scientific requirements. From these preliminary choices, different kinds of AO concepts are investigated : Ground Layer Adaptive Optics (GLAO), Multi-Object AO (MOAO) and Laser Guide Stars AO (LGS). Using Fourier based simulations we study the performance of these AO systems depending on the telescope diameter. ",
        "introduction": "The new era of astronomical telescopes with diameters reaching 30 to 100m will provide a dramatic advance in our understanding of the universe. The concept of ELTs is unique to complement the detection which will be made from Space. By accommodating high spectral resolution, 3D spectroscopy devices and relatively large fields of view (10 arc minutes), ELTs will be a privileged tool for the study of formation and evolution of galaxies.\\\\ Adaptive Optics enables large telescopes to provide diffraction limited images by real time correction of turbulence. Because the contamination by the interstellar medium light is one of the main issue in extragalactic studies it is required to observe in a direction far from our galactic plane. In that case, the density of stars becomes dramatically small, and because of anisoplanatism effects, classical AO working on Natural Guide Star (NGS) can not be used. To overcome this problem, new AO techniques (GLAO, MOAO) have emerged in the last few years to increase the corrected field using NGS. These new methods are based on a full measurement of the $3$D turbulent volume using several NGS. In other hand and to improve sky coverage, Laser Guide Stars (LGS) have been proposed \\cite{Foy85}; \\cite{Lelouarn98}. Should some critical issues be solved (cone effect, spot elongation, tilt indetermination) this last solution would allow the ultimate scientific requirements to be met.\\\\ In section 2 we present the specifications required to achieve scientific goals imposed by astrophysical studies of high redshifted galaxies. Section 3 is then dedicated to a first performance evaluation for AO systems, and some preliminary results from Fourier based simulations will be shown. ",
        "conclusions": "We have presented here some recent results concerning AO concepts for futurs ELTs and particularly for 3D spectroscopy. First numerical simulations show that even for quasi-ideal case, GLAO could not achieve the scientific goal, MOAO could be used with a limited sky coverage and only a LGS based system could offer a 100\\% sky coverage. This article also shows that the implementation of rather simple LGS systems without Tip-tilt correction could be of interest specifically for very large telescope diameters.\\\\ Additional simulations are mandatory to precise the science specifications and evaluate if coupling or aperture size could be relaxed. Further analysis should then lead to a detailled AO system definition."
    },
    "0512/hep-th0512262_arXiv.txt": {
        "abstract": "{ We consider a physically viable cosmological model that has a field dependent Gauss-Bonnet coupling in its effective action, in addition to a standard scalar field potential. The presence of such terms in the four dimensional effective action gives rise to several novel effects, such as a four dimensional flat Friedmann-Robertson-Walker universe undergoing a cosmic inflation at the early epoch, as well as a cosmic acceleration at late times. The model predicts, during inflation, spectra of both density perturbations and gravitational waves that may fall well within the experimental bounds. Furthermore, this model provides a mechanism for reheating of the early universe, which is similar to a model with some friction terms added to the equation of motion of the scalar field, which can imitate energy transfer from the scalar field to matter.\\\\   {KEYWORDS}: Inflation, Dark Energy Cosmologies, Classical Theories of Gravity }  \\begin{document}   \\noindent ",
        "introduction": "Although Einstein's theory has been proven to be remarkably simple and successful as a classical theory of gravitational interactions, there are several observational facts which it has failed to elucidate. These cosmological conundrums include both cosmic inflation, or a period of accelerated expansion in the early universe, and a recent acceleration in the expansion of the universe.  Inflation in the early universe is a very attractive proposal for explaining the present large scale homogeneity and high degree of isotropy of the universe (one part in $100,000$), in addition to the observed spectrum of density perturbations, which is usually attributed to a scalar field rolling down a shallow potential. Similarly, the current acceleration of the universe, as indicated by recent cosmological results~\\cite{Bennett03a}, is usually attributed to some form of cosmic fluid having a large and smoothly distributed negative pressure, usually called {\\it dark energy} or {\\it dark pressure}.  Cosmologists have long wondered why/how the universe has been recently accelerating: is it due to a pure cosmological constant term, or due to some sort of negative pressure generated by one or more dynamical scalar fields, or something else? In recent years, different explanations have been provided for both inflation and the current epoch of acceleration: some examples of recent interest include brane-world modification of Einstein's general relativity (GR), including a 5d DGP (Dvali-Gabadadze-Porrati) model~\\cite{Dvali:2000a}. The names of the dark energy candidates run the gamut from $f(R)$ gravity~\\cite{Carroll:2003a} (modifying in a very radical manner the Einstein's GR itself) to ghost condensates (the idea which ncludes a more or less disguised non-locality)~\\cite{Arkani:2003a}. Many of these new proposals are pathological and do not appear more appealing than the two long envisioned alternative models of dark energy: a cosmological constant~\\cite{Weinberg:1988cp} and a slowly varying $\\Lambda$-term~\\cite{Peebles:1988}.  The cosmological constant is a pure dark energy or vacuum energy, while the variable $\\Lambda$-term is some kind of exotic matter or a slowly varying potential of a scalar field, usually referred to as {\\it quintessence}~\\cite{Steinhardt:1998}. This last category can comprise a Casimir energy or vacuum polarization effect from additional compact or curved non-compact spatial dimensions, which only weakly couple to ordinary matter {\\it in contrast} to most tentative quintessence models, including k-essence~\\cite{Picon:2000} or curvature quintessence. In a brane world scenario~\\cite{Ish2001kd}, for instance, the vacuum energy (or the dark energy) may be viewed as a smooth {\\it brane-tension} if our universe is a 3-brane embedded in higher dimensional spacetimes.  Indeed, the past decade has witnessed significant progress in the building of inflationary models as extensions of standard cosmology, to accommodate the effects of dark energy. However, most of the inflationary type potentials studied in the literature are picked up in a very {\\it ad hoc} fashion, rather than constructing such a potential as a valid solution of the field equations that follows, for instance, from low energy string effective actions. In this paper we initiate work in this direction. ",
        "conclusions": "In this paper we presented an analysis of accelerating/inflationary cosmologies by introducing in the effective action a field dependent Gauss-Bonnet coupling, other than a standard field potential for the field $\\sigma$. We find that the dark energy hypothesis fits into a low energy gravitational action where a scalar field is coupled to the curvature squared terms in Gauss-Bonnet combination. It is established that a GB scalar-coupling can play an important and interesting role in explaining both the early and late-time evolutions of the universe as well as providing a mechanism for reheating.  That we are able to explain accelerating universes using exact cosmological solutions in a modified Gauss-Bonnet theory, leading to a small deviation from the $w=-1$ prediction of non-evolving dark energy (or a cosmological constant) is likely to have a serious impact in search of a viable dark energy model. Our work also provides extension of quintessence (or time-varying $\\Lambda$) model in which part of the dark energy comes from a field dependent Gauss-Bonnet interaction term.  One of the key results is this: in the absence of a Gauss-Bonnet coupling, the tensor/scalar ratio is usually non-zero. However, with a non-trivial scalar Gauss-Bonnet coupling, i.e., $f(\\sigma)\\neq $, or effectively, $u(\\sigma) \\equiv f(\\sigma) H^2 \\sim \\e^{\\alpha N} \\neq 0$, such a ratio can be negligibly small if the expansion parameter $\\alpha$ takes a small positive value, $\\alpha \\gtrsim 0.1$, and hence $n_s\\lesssim 1$ and $n_T\\simeq 0$, leading to Harrison-Zel'dovich spectrum. Unlike a naive expectation, the inclusion of a (scalar) field dependent Gauss-Bonnet coupling $f(\\sigma)$, in addition to a field potential $V(\\sigma)$, into the effective action, could make the observability of tensor/scalar ratio and related inflationary parameters more achievable.  We emphasize that, in contrast to previous analysis, our calculations have all been implemented by the functional forms of the scalar potential and GB scalar coupling, as suggested by the symmetry of the field equations, rather than choosing particular model dependant forms for them. We have given in the Appendix the exact solutions for some special cases, about which a general comparison can be made in terms of the homogeneous solutions we presented in the bulk part of the paper.  Regardless of whether the model studied here appears natural or otherwise, it should be observation that determines whether or not it is correct. The current and future observations might make stronger demand on theoretical precision of inflationary parameters, including, the scalar and tensor spectral indices, and are certain to constrain a number of parameters of our model tightly, including the Gauss-Bonnet coupling constants."
    },
    "0512/astro-ph0512243_arXiv.txt": {
        "abstract": "We study the physical properties derived from interstellar cloud complexes having a fractal structure. We first generate fractal clouds with a given fractal dimension and associate each clump with a maximum in the resulting density field. Then, we discuss the effect that different criteria for clump selection has on the derived global properties. We calculate the masses, sizes and average densities of the clumps as a function of the fractal dimension ($D_f$) and the fraction of the total mass in the form of clumps ($\\epsilon$). In general, clump mass does not fulfill a simple power law with size of the type $M_{cl} \\varpropto R_{cl}^{\\gamma}$, instead the power changes, from $\\gamma \\simeq 3$ at small sizes to $\\gamma < 3$ at larger sizes. The number of clumps per logarithmic mass interval can be fitted to a power law $N_{cl} \\varpropto M_{cl}^{-\\alpha_M}$ in the range of relatively large masses, and the corresponding size distribution is $N_{cl} \\varpropto R_{cl}^{-\\alpha_R}$ at large sizes. When all the mass is forming clumps ($\\epsilon = 1$) we obtain that as $D_f$ increases from $2$ to $3$ $\\alpha_M$ increases from $\\sim 0.3$ to $\\sim 0.6$ and $\\alpha_R$ increases from $\\sim 1.0$ to $\\sim 2.1$. Comparison with observations suggests that $D_f \\simeq 2.6$ is roughly consistent with the average properties of the ISM. On the other hand, as the fraction of mass in clumps decreases ($\\epsilon < 1$) $\\alpha_M$ increases and $\\alpha_R$ decreases. When only $\\sim 10\\%$ of the complex mass is in the form of dense clumps we obtain $\\alpha_M \\simeq 1.2$ for $D_f=2.6$ (not very different from the Salpeter value $1.35$), suggesting this a likely link between the stellar initial mass function and the internal structure of molecular cloud complexes. ",
        "introduction": "The fact that the interstellar medium (ISM) has a hierarchical and self-similar structure when observed with sufficiently high dynamic range, is interpreted as evidence of an underlying fractal structure \\citep{sca90}. The boundaries of the projected images of interstellar clouds are irregular curves whose fractal dimension is around $\\sim 1.3$ \\citep[e.g.][]{fal91,lee04}. Apparently this is a universal result which does not depend on whether tracers of atomic, molecular or dust components are used, whether clouds are selfgravitating or not, etc \\citep{wil00}. The fractal dimension of the projected boundaries is usually associated with a three-dimensional fractal dimension $D_f \\simeq 2.3$ for the ISM \\citep[e.g.][]{bee92}. It has been argued that a fractal ISM with $D_f \\simeq 2.3$ could account for the observed mass and size distributions of the interstellar clouds \\citep{elm96}, as well as for the intercloud medium properties \\citep{elm97a}, and even for the stellar initial mass function \\citep{elm97b,elm99}.  In a previous paper \\citep{san05} we studied the effect that the projection of clouds has on the estimation of the fractal dimension, and we concluded that a value around $\\sim 1.3$ for the projected boundaries is more consistent with three-dimensional clouds having $D_f \\sim 2.6 \\pm 0.1$. The application of $\\Delta$-variance techniques to Polaris Flare cloud by \\citet{stu98} yielded a fractal dimension for the cloud surfaces $\\simeq 2.6$ \\citep[see also][]{ben01}. \\citet{elm02} simulated fractal brownian motion clouds with average fractal dimension $\\sim 2.75$ and obtained properties in gross agreement with observations. A fractal medium with a relatively high $D_f$ value can reproduce observations of HII regions surrounding stars \\citep{woo05}. \\citet{hen86,hen91} used a gravitationally driven turbulence model to study the properties of giant molecular clouds suggesting that a dimension $\\simeq 2.7$ could be necessary to explain the observed properties. Also \\citet{fle96} analyzed the properties of the turbulent, non-self-gravitating, neutral component of the ISM by using a model of compressible turbulence, concluding that the compression parameter that better reproduces observations is such that $D_f \\simeq 2.5$.  It is important to quantify the degree of complexity (through, for example, the fractal dimension) as a first step towards understanding the physical mechanisms responsible for structuring the ISM. In models of fractally homogeneous turbulence the fractal dimension of fully developed turbulence is $2.5 < D_f < 2.75$ \\citep{hen82}. On the other hand, turbulent diffusion in a incompressible medium generates structures with $D_f \\sim 2.3$ for a Kolmogorov spectrum \\citep{men90}. It is not obvious, however, that the energy can cascade down without any dissipation or injection, and deviation from a Kolmogorov spectrum is, in general, expected in the ISM \\citep{bru02a,bru02b}. The ISM is a highly compressible and turbulent medium and its fractal dimension depends, among other factors, on the degree of compressibility \\citep{fle96} or on the Mach number \\citep{pad04}. Also, self-gravity could by itself explain many observed properties \\citep{dev96}, although this fact does not imply that turbulence is not an important factor in the ISM.  Here we are interested in understanding the relationship between the physical properties of the interstellar clouds and their fractal structure, and in verifying whether observed properties are in agreement with relatively high fractal dimension values \\citep[as suggested in][]{san05}. In other words, we investigate what physical properties are intimately connected to the cloud ``geometry\" and whether this geometry, no matter how it was originated, defines its next evolutionary stage. Our approach in this work is to calculate and analyze the properties resulting from a hierarchical structure with a known and perfectly defined fractal dimension, no mattering the physical processes behind this structure. In order to do this, we simulate interstellar clouds by using a simple algorithm explained in section~\\ref{sec_fractals}, which generates a distribution of points with a very well defined fractal dimension. In section~\\ref{sec_pdf} we address the density fields resulting from the generated fractals and in section~\\ref{sec_clump} we discuss the different ways to construct ``clumps\" from these density fields. Section~\\ref{sec_results} is devoted to study the properties (masses, densities, sizes, and mass and size distributions) of the derived clumps and to compare these results with observations. Finally, the main conclusions are summarized in section~\\ref{sec_conclusion}. ",
        "conclusions": ""
    },
    "0512/gr-qc0512032_arXiv.txt": {
        "abstract": "We present a numerical model of a collapsing radiating sphere, whose boundary surface undergoes  bouncing due to a  decreasing of its inertial mass density (and, as expected from the equivalence principle, also of the ``gravitational'' force term)   produced by the ``inertial'' term of the transport equation. This model exhibits for the first time  the consequences of  such an effect, and shows that under physically reasonable conditions this decreasing of the gravitational term in the dynamic equation may be large enough as to revert the collapse and produce a  bouncing of the boundary surface of the sphere. ",
        "introduction": "In the study of gravitational collapse of massive stars, the inclusion of dissipative processes (in particular neutrino emission) is enforced by the fact that they provide the only plausible mechanism to carry away the bulk of binding energy, leading to a neutron star or black hole \\cite{1}. On the other hand, in cores of densities about $10^{12}g\\,\\,cm^{-3}$ the mean free path of neutrinos becomes small enough as to justify the use of diffusion approximation \\cite{3,4}. This seems to be confirmed by the observational data collected from supernova 1987A, which indicates that the radiation transport regime prevailing during the emission process, is closer to the diffusion approximation than to the streaming out limit \\cite{5}.  Motivated by the comments above, in a recent paper \\cite{Herrera1}, the Misner and Sharp approach to the study of adiabatic gravitational collapse \\cite{MisnerSharp} was extended as to include dissipation in, both, the streaming out and diffusion approximation (for the case of pure free streaming approximation see \\cite{Misner}). Then from the coupling of the dynamical equation to a causal transport equation in the context of M\\\"uller--Israel--Stewart theory  \\cite{Muller67,IsSt76} it was obtained that the effective inertial mass density of a fluid element and the gravitational force term in the dynamical equation, reduce by a factor which depends on dissipative variables. This reduction, in its turn, might lead to the bouncing of the collapsing sphere, as discussed in \\cite{Herrera1}.  As can be seen from inspection of the transport equation, such an effect is directly related to the presence of the inertial term $Ta_{\\beta}$ in the transport equation.  This explains why we refer to  such a bouncing as ``thermo--inertial''.  It is our purpose in this work to present a numerical model of a radiating collapsing sphere, where the above mentioned effect produces the bouncing of the boundary surface of the sphere, for physically acceptable values of all variables.  Since we are mainly concerned with time scales of the order of magnitude of (or even smaller than) the hydrostatic time scale, as in the quick collapse phase  preceding neutron star formation, we cannot rely on the quasistatic approximation, and therefore the full dynamic description has to be used \\cite{8A,9A}. This implies that we have to  appeal to a hyperbolic theory of dissipation. The use of a hyperbolic theory of dissipation is further justified  by the necessity of overcoming the difficulties inherent to parabolic theories (see references \\cite{6}--\\cite{22} and references therein).  The plan of the paper is as follows. In the next section we define the conventions and present the dynamical equation  coupled to the transport equation. The model to be considered as well as the strategy for the numerical integration is presented in Section III. Finally, a discussion of results is presented in Section IV. ",
        "conclusions": "The influence of  pre--relaxation effects on gravitational collapse has been brought out in many works in last decade \\cite{pre}, however the specific effect of bouncing, associated with the decreasing of the effective inertial mass density, produced by the increasing of $\\alpha$, had not been illustrated until now. It is worth stressing that $\\alpha$--terms in Eq. (\\ref{V4}) come from the inertial factor $Ta_\\beta$ in Eq. (\\ref{21}).  In this work we provide a numerical model of such bouncing, by assuming an increasing of $\\alpha$ at the boundary surface. We have concentrated the increase of $\\alpha$ on the boundary surface to illustrate the effect, the remaining of the sphere is assumed to be dissipating at much lower values of $\\alpha$. Of course, the increasing of $\\alpha$ may in principle occur at any region of the sphere and even in more that one, simultaneously. The results of our integration is deployed in the figures {\\ref{fig:r}}--{\\ref{fig:t}}, which exhibit the evolution of different variables with respect to the dimensionless time $t/r_{\\Sigma}$.  Figure  {\\ref{fig:r}}   shows the evolution of $R_{\\Sigma}$ for different values of $\\alpha_m$ from $0$ to $1$, the bouncing is clearly exhibited as well as its dependence on $\\alpha$. Figure {\\ref{fig:Ralpha}} emphasizes further the link between the increasing of $\\alpha$ and the bouncing.  Figures {\\ref{fig:d}}--{\\ref{fig:t}}, shows the  behaviour of (dimensionless) energy density, pressure, heat flow and temperature, evaluated at the boundary surface. Their values are always regular and satisfy the physical conditions $\\rho>P>0$.  The dimensionless quantity $\\kappa T_{\\Sigma}$ plotted in Figure {\\ref{fig:t}}  is, in conventional units, \\begin{equation} \\kappa T_{\\Sigma}=2\\,\\, 10^6 \\frac{G}{c^5}[\\kappa][T_{\\Sigma}] \\label{temp1} \\end{equation} with $G$ and $c$ denoting the gravitational constant and the speed of light, and where $[\\kappa]$ and $[T]$ denote the numerical values of conductivity and temperature in $g \\,\\, cm^{-3}\\,\\, K^{-1}$ and $K$ respectively. Therefore the maximum values of $\\kappa T_{\\Sigma}$ reached just after the bouncing, correspond to \\begin{equation} [\\kappa][T_{\\Sigma}] \\approx 10^{46} \\label{temp2} \\end{equation} which may be obtained with $[T_{\\Sigma}]\\approx 10^{12}$ and $[\\kappa] \\approx 10^{34}$. These values are well within the acceptable range for those variables in a pre--supernovae event \\cite{Ma}.  Thus we have seen that a relatively simple model, whose physical variables exhibit good behaviour and have acceptable numerical values, may serve to illustrate the bouncing of a dissipating self--gravitating sphere, produced by the decreasing of its effective inertial mass density associated to an increasing of $\\alpha$.  Nevertheless, in spite of the appeal of the presented model, we are well aware that invoking such an effect to describe a specific observed phenomena, would require a much more detailed astrophysical setting. This, however,  is out of the scope of this paper."
    },
    "0512/astro-ph0512075_arXiv.txt": {
        "abstract": "{\\bf Recent observations\\cite{Burnes2002,Veillet2002,Margot2002a} have revealed an unexpectedly high binary fraction among the Trans-Neptunian Objects (TNOs) that populate the Kuiper Belt.  The TNO binaries are strikingly different from asteroid binaries in four respects\\cite{Veillet2002}: their frequency is an order of magnitude larger, the mass ratio of their components is closer to unity, and their orbits are wider and highly eccentric. Two explanations have been proposed for their formation, one assuming large numbers of massive bodies\\cite{Weidenschilling2002}, and one assuming large numbers of light bodies\\cite{Goldreich2002}.  We argue that both assumptions are unwarranted, and we show how TNO binaries can be produced from a modest number of intermediate-mass bodies of the type predicted by the gravitational instability theory for the formation of planetesimals\\cite{GoldreichWard1973}. We start with a TNO binary population similar to the asteroid binary population, but subsequently modified by three-body exchange reactions, a process that is far more efficient in the Kuiper belt, because of the much smaller tidal perturbations by the Sun.  Our mechanism can naturally account for all four characteristics that distinguish TNO binaries from main-belt asteroid binaries. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512305_arXiv.txt": {
        "abstract": "We have obtained 2D spectral data for a sample of 58 nearby S0 galaxies with the Multi-Pupil Spectrograph of the 6m telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences. The Lick indices H$\\beta$, Mgb, and $\\langle \\mbox{Fe} \\rangle$ are calculated separately for the nuclei and for the bulges taken as the rings between $R=4\\arcsec$ and $7\\arcsec$; and the luminosity-weighted ages, metallicities, and Mg/Fe ratios of the stellar populations are estimated by confronting the data to SSP models. Four types of galaxy environments are considered: clusters, centers of groups, other places in groups, and field. The nuclei are found to be on average slightly younger than the bulges in any types of environments, and the bulges of S0s in sparse environments are younger than those in dense environments. The effect can be partly attributed to the well-known age correlation with the stellar velocity dispersion in early-type galaxies (in our sample the galaxies in sparse environements are in average less massive than those in dense environments), but for the most massive S0s, with $\\sigma _*=170-220$ km/s, the age dependence on the environment is still significant at the confidence level of 1.5$\\sigma$. ",
        "introduction": "In classical morphological sequence by \\citet{hubble} lenticular galaxies occupy intermediate position between ellipticals and spirals: they have a smooth and red appearance as the ellipticals, but also have stellar disks, almost as large as those of the spirals. The most popular hypothesis of S0 origin is that of their transformation from the spirals by stopping global star formation and removing or consuming remaining gas \\citep{ltcs0}. In distant, $z\\sim 0.5$, clusters this transformation is now observed directly: the number of lenticulars in the clusters diminishes strongly with the redshift \\citep{fasano}, instead one can see `passive spirals' -- red spiral galaxies lacking star formation -- at the periphery (`infalling regions') of the intermediate-redshift clusters \\citep{goto, yamauchi}. Many theoretical works have been done to explain in detail what physical mechanisms may be involved into the process of spiral transformation into the lenticulars: tidally induced collisions of disk gas clouds \\citep{byrdval90}, harassment \\citep{moore96}, ram pressure by intercluster medium \\citep{quilis}, etc. For the S0s in the field, the scheme of their transformation from the spirals is not so clear, but common view is that some external action like minor merger may produce the necessary effect.  By reviewing the various mechanisms of secular evolution which may transform a spiral galaxy into a lenticular one we have noticed that most of them result in gas concentration in the very center of the galaxy, so that a nuclear star formation burst seems inavoidable circumstance of the S0 galaxy birth. If to refer to S0 statistics in the clusters located between $z=0$ and $z\\approx 1$, the main epoch of S0 formation is $z\\approx 0.4-0.5$, so the nuclear star formation bursts in the nearby S0s must not be older than 5 Gyr. Indeed, in my spectral study of the central parts of nearby galaxies in different types of environments \\citep{me93} I have found that $\\sim 50\\%$ of nearby lenticulars have strong absorption lines H$\\gamma$ and H$\\delta$ in their nuclear spectra so they are of `E+A' type, as it is presently called and are dominated by intermediate-age stellar population. In this respect the S0s have resembled rather early-type spirals than ellipticals. Here I aim to continue this study, with a larger sample and with panoramic spectral data in order to separate the nuclei and their outskirts (bulges) which is a substantial advantage with respect to aperture spectroscopy.  Another crucial point of the present study, and also of a global paradigm of galaxy formation, is environmental influence. The current hierarchical assembly paradigm predicts a younger age of galaxies in lower density environments -- for the most recent simulations see e.g. \\citet{lanzoni05} or \\citet{delucia05}. Observational evidences concerning early-type galaxies are controversial: some authors find differences of stellar population ages between the clusters and the field \\citep{terlfor,kunt,thogal}, some authors do not find any dependence of the stellar population age on environment density \\citep{gallens}. In order to check whether the mean ages of the stellar populations depend on environment density monotonously, as the hierarchical paradigm predicts, in this work I consider four types of environments separately: the cluster galaxies, the brightest (central) galaxies of groups, the second-ranked group members, and the field galaxies. ",
        "conclusions": "By considering the stellar population properties in the nuclei and the bulges of the nearby lenticular galaxies in the various types of environments, I have found certain differences between the nuclei and the bulges as well as between the galaxies in dense and sparse environments. The nuclei are on average younger than the bulges in any types of environments, and both the nuclei and the bulges of S0s in sparse environments are younger than those in dense environments. The results of the consideration of the Mg/Fe ratios suggest that the main star formation epoch may be more brief in the centers of the galaxies in dense environments."
    },
    "0512/astro-ph0512133_arXiv.txt": {
        "abstract": "Close binary systems of compact stars, due to the emission of gravitational radiation, may evolve into a phase in which the less massive star transfers mass to its companion. We describe mass transfer by using the model of Roche lobe overflow, in which mass is transferred through the first, or innermost, Lagrange point. Under conditions in which gravity is strong, the shapes of the equipotential surfaces and the Roche lobes are modified compared to the Newtonian case. We present calculations of the Roche lobe utilizing the second order post-Newtonian (2PN) approximation in the Arnowitt-Deser-Misner gauge. Heretofore, calculations of the Roche lobe geometry beyond the Newtonian case have not been available. Beginning from the general N-body Lagrangian derived by Damour and Sch\\\"affer, we develop the Lagrangian for a test particle in the vicinity of two massive compact objects. As an exact result for the transverse-traceless part of the Lagrangian is not available, we devise an approximation that is valid for regions close to the less massive star. We calculate the Roche lobe volumes, and provide a simple fitting formula for the effective Roche lobe radius analogous to that for the Newtonian case furnished by Eggleton. In contrast to the Newtonian case, in which the effective Roche radius depends only upon the mass ratio $q=m_1/m_2$, in the 2PN case the effective Roche lobe radius also depends on the ratio $z=2 (m_1+m_2)/a$ of the total mass and the orbital separation. ",
        "introduction": "During the evolution of a close binary system involving compact stars, the stellar separation shrinks due to the emission of gravitational waves. In the event that the stars are not of equal mass, and the less massive star has a larger radius than its companion, mass transfer may ultimately occur. Gravity wave emission generally causes the mutual orbit to circularize \\citep{PETERS64}. For circular orbits, conservative mass transfer can be modelled as Roche lobe overflow under the assumption that the star is not significantly disrupted due to tidal interactions. The Roche lobe is the innermost gravitational plus centrifugal equipotential surface encompassing both stars.  In the model Roche lobe overflow, the radius of the less massive star is compared to the effective radius of its Roche lobe. Once the two radii become equal, because the Roche lobe radius decreases due to orbital decay, the star fills its Roche lobe and mass transfer occurs through the first, or innermost, Lagrange point $L_1$. Lying on the Roche lobe, $L_1$ is located between the two stars on the axis connecting their centers and is also a saddle point of the gravitational plus centrifugal potential between the two stars. Due to its saddle point nature, the first Lagrange point acts as a gravitational funnel through which mass transfer occurs.  Values of the Roche lobe radii as a function of orbital separation and mass ratio $q=m_1/m_2$, where $m_1$ refers to the lighter star, have been tabulated by \\citet{KOPAL1} for the Newtonian case. \\citet{PACZYNSKI1} and \\citet{EGGLETON1} have given analytical  fits. We use Eggleton's functional form, which has the advantage of being a continuous function of $q$, as a template in our work.  In this work, we carry out calculations of Roche lobes beyond the Newtonian case. We employ the Arnowitt-Deser-Misner (ADM) form of post-Newtonian expansion and use the corresponding Lagrangian at the second order (2PN) level wherein terms up to $(M/r)^2$, where $M=m_1+m_2$ and $r$ is the distance, are retained.  The same procedure as used in the Newtonian case for finding the Roche lobes is utilized.  Our strategy is to (i) construct the effective potential for the point particle in the vicinity of two stars (the 3--body problem) in the co--rotating frame; (ii) evaluate equipotential surfaces and calculate the corresponding effective Roche volume and radius for this potential; and (iii) provide new fitting formulae as Eggleton did for applications involving mass transfer.  The organization of this work is as follows. In \\S \\ref{SEC:2PNPOTENTIAL}, we calculate the effective potential for three bodies at the 2PN level. We establish the Lagrangian in \\S \\ref{SEC:L2PN}. The transverse-traceless part of the Lagrangian is evaluated explictly in \\S\\ref{SEC:UTT} through the introduction of an approximation valid for regions near $m_1$ for test particles. In \\S \\ref{SEC:2PNROCHE}, we evaluate the Roche lobes and their effective radii as a function of $q$ and a relativity parameter for this potential, and provide a simple analytical fit. In this section, we also show the impact of post--Newtonian corrections on the positions of the Lagrange points and on the position of the center of mass. Our conclusions are contained in \\S \\ref{sec:conclusion}.   \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=0.9\\textwidth]{3body_figure.eps} \\caption{ The notation used in the evaluation of the Roche lobes in the 2PN approximation.  Stellar masses are denoted by $m_1$ and $m_2$ and the point-particle mass is taken to be $m_0$. Vectors ${\\bf R}_A$ ($A=0,1,2$) denote positions of the three bodies with respect to the origin $O$, ${\\bf r}$ is the position of a generic point $P$, and ${\\bf r}_A$ is the position of this point with respect to the mass $m_A$ (we show only ${\\bf r}_0$). The vectors ${\\bf R}_{AB}$ indicate positions of the three bodies with respect to each other. \\label{FIG:3BODY}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:conclusion}  We have utilized the second order post--Newtonian approximation in the Arnowitt--Deser--Misner gauge to calculate Roche lobe volumes. These results are an improvement over the Newtonian case in that post--Newtonian gravity introduces corrections in the case of moderately strong gravitational field.  In the course of our calculations, we have derived an approximate three--body Lagrangian that is valid in the case when one of the bodies is a point particle. This calculation requires an evaluation of the transverse--traceless term $U_{TT}$ of the Lagrangian for which an exact result is not available.  However, as shown in this work, utilization of an approximation valid in the vicinity of the less massive star enables this problem to be circumvented.  Using these results, we calculated Roche lobes in the 2PN effective potential in the co--rotating frame and computed effective Roche lobe radii that can be used to model mass transfer through Roche lobe overflow. In addition, we computed changes to the positions of the Lagrange points and to the center of mass due to post--Newtonian effects. We find that corrections to Newtonian results for Roche lobe radii can be as significant as 20--30\\% at low mass ratio $q\\lesssim 0.1$. Whereas for $q\\gtrsim 0.7$ the Roche lobe radius increases ($\\approx 15\\%$ for $q=1.0$), for low $q$'s the Roche lobe is smaller than in the Newtonian case. We have provided our results in the form of a simple fitting formula that depends on two physical parameters: the mass ratio and the ratio of the total mass and the separation.  Research support of the U.S. Department of Energy under grant number DOE/DE-FG02-87ER-40317 is gratefully acknowledged."
    },
    "0512/astro-ph0512168_arXiv.txt": {
        "abstract": "We present a $\\sim$19 ks \\textit{Chandra} ACIS-S observation of the globular cluster Terzan 1.  Fourteen sources are detected within 1\\farcm4 of the cluster center with 2 of these sources predicted to be not associated with the cluster (background AGN or foreground objects).  The neutron star X-ray transient, X1732$-$304, has previously been observed in outburst within this globular cluster with the outburst seen to last for at least 12 years. Here we find 4 sources that are consistent with the \\textit{ROSAT} position for this transient, but none of the sources are fully consistent with the position of a radio source detected with the VLA that is likely associated with the transient. The most likely candidate for the quiescent counterpart of the transient has a relatively soft spectrum and an unabsorbed 0.5-10 keV luminosity of $2.6\\times10^{32}$ ergs s$^{-1}$, quite typical of other quiescent neutron stars.   Assuming standard core cooling, from the quiescent flux of this source we predict long ($>400$ yr) quiescent episodes to allow the neutron star to cool. Alternatively, enhanced core cooling processes are needed to cool down the core.  However, if we do not detect the quiescent counterpart of the transient this gives an unabsorbed 0.5-10 keV luminosity upper limit of $8\\times10^{31}$ ergs s$^{-1}$.  We also discuss other X-ray sources within the globular cluster.  From the estimated stellar encounter rate of this cluster we find that the number of sources we detect is significantly higher than expected by the relationship of \\citet{pooley2003}. ",
        "introduction": "Neutron star X-ray transients form a subgroup of low-mass X-ray binaries. Although usually in a quiescent state with typical X-ray luminosities of $10^{32}$ - $10^{34}$ ergs s$^{-1}$, occasionally these systems go into outburst where the luminosity increases to around $10^{36}$ - $10^{38}$ ergs s$^{-1}$.  These outbursts are attributed to a large increase in the mass accretion rate onto the neutron star. The X-ray spectra from these quiescent neutron star transients are usually dominated by a soft component at around 1 keV, and in some cases an additional power-law component, dominating at energies above a few keV, is present. The most widely accepted model used to explain the soft, thermal X-ray emission of these transients in their quiescent states is that in which the emission is due to the cooling of the neutron star, which has been heated during the outbursts. In this case, the quiescent luminosity should depend on the time-averaged accretion rate \\citep{BBR98,campanaetal98a}.  However, the X-ray emission from the cooling of the neutron star cannot directly explain the hard power-law tail and its origin is not well understood.  Suggestions include residual accretion down to the magnetospheric radius, or pulsar shock emission \\citep[e.g.,][]{stellaetal94,campanaetal98a,campanastella00,menoumcclintock01}.  Many quiescent neutron stars and other quiescent low-mass X-ray binaries have been found in several globular clusters using the \\textit{Chandra} and \\textit{XMM-Newton} X-ray observations \\citep[see][and references therein]{heinkeetal03,pooley2003}. Galactic globular clusters provide an ideal location to study these types of sources - the distance to host clusters can usually be determined more accurately than for the Galactic quiescent X-ray binaries.  The known distance and reddening allows accurate luminosities to be derived and removes the distance uncertainty from the quiescent properties.  The high incidence of compact binaries in globular clusters is likely explained by the formation of such binaries via exchange encounters in the very dense environments present \\citep*[e.g.][]{verbunt87,hut91}, though the ultracompact systems may be formed via direct collisions followed by orbital decay \\citep{ivanova05}.  Terzan 1 \\citep{terzan66} is a globular cluster at a distance of $5.2 \\pm 0.5$ kpc and a reddening of $E(B-V) = 2.48 \\pm 0.1$ \\citep{ortolanietal99}.  In 1980, the \\textit{Hakucho} satellite detected X-ray bursts from a source located in Terzan 1 \\citep{makishimaetal81,inoueetal81}.  Several years later, the persistent source X1732$-$304 was detected within Terzan 1 with SL2-XRT onboard Spacelab 2 \\citep{skinneretal87} and \\textit{EXOSAT} \\citep{warwicketal88,parmaretal89}. This is likely to be the same source as the bursting source detected previously by \\textit{Hakucho}.  Subsequent X-ray observations with \\textit{ROSAT} \\citep{JVH95,verbuntetal95} detected the source with a similar luminosity, between $2.0\\times10^{35}$ ergs s$^{-1}$ and $1.3\\times10^{36}$ ergs s$^{-1}$ \\citep[see Fig.~3 of][]{GPO99}.  It is also assumed that X1732$-$304 is the source of hard X-rays detected with the SIGMA and ART-P telescopes \\citep{borreletal96a,borreletal96b,pavlinskyetal95}.  A possible radio counterpart of X1732$-$304 was observed with the VLA within the \\textit{ROSAT} error circles \\citep{martietal98}.  A \\textit{BeppoSAX} observation of X1732$-$304 in April 1999 discovered the source in a particularly low state with the X-ray intensity more than a factor of 300 lower than previous measurements \\citep{GPO99}.  A more recent short ($\\sim$ 3.6 ks) \\textit{Chandra} HRC-I observation of Terzan 1 did not conclusively detect X1732$-$304 with a 0.5-10 keV luminosity upper limit of $(0.5-1)\\times10^{33}$ ergs s$^{-1}$ depending on the assumed spectral model \\citep*{WHG02}.  However, an additional X-ray source, CXOGLB J173545.6-302900, was detected by this observation.  In this paper we study the X-ray sources detected in a recent $\\sim$19 ks \\textit{Chandra} ACIS-S observation of Terzan 1 and discuss possible quiescent counterparts of the neutron-star X-ray transient X1732$-$304. ",
        "conclusions": "We have presented a $\\sim$19 ks \\textit{Chandra} ACIS-S observation of the globular cluster Terzan 1.  Within 1.4 arcmin of the cluster centre we detect 14 X-ray sources and predict that 2 of these are not associated with the cluster (background AGN or foreground objects).  The brightest of these sources, CX1, is consistent with the position of CXOGLB J173545.6-302900 first observed during a short \\textit{Chandra} HRC-I observation of Terzan 1 \\citep{WHG02}.  This source has a particulary hard spectrum for a globular cluster source, with a simple power-law fit giving a photon index of $0.2\\pm0.2$.  Such a hard spectrum suggests that this could be an intermediate polar \\citep{muno2004}.  The position of the second brightest source, CXOGLB J173547.2-302855 (CX2), is the only source to have a position that is consistent with all 3 of the previous {\\it ROSAT} pointings that observed the neutron star transient X1732$-$304 \\citep{JVH95} as well as the {\\it EXOSAT} observation.  This source's X-ray color and spectrum suggest that it could be a quiescent neutron star and therefore we find that this source is the most likely candidate for the quiescent counterpart of X1732$-$304.  CX2 has a position closest to the position of the radio source detected in Terzan 1 by the VLA \\citep{martietal98}, though the positions are not fully consistent.  Assuming standard core cooling, from the quiescent flux of CX2 we predict extremely long ($>$400 yr) quiescent episodes of X1732$-$304, or enhanced core cooling, to allow the neutron star to cool.  Having estimated the stellar encounter rate of this cluster we find significantly more sources than expected by the relationship of \\citet{pooley2003} perhaps because the cluster was previously much larger and that most of the stars have been lost due to passages through the Galactic disk.  \\subsection*{Acknowledgements}  The authors wish to thank the referee, Frank Verbunt, for careful and helpful comments which have improved this paper.  EMC gratefully acknowledges the support of a PPARC Studentship at the University of St Andrews.  DP gratefully acknowledges support provided by NASA through Chandra Postdoctoral Fellowship grant PF4-50035 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060.  The authors wish to thank P. Broos for his support with ACIS Extract.  This work makes use of the Digitized Sky Surveys which were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This research has made use of the USNOFS Image and Catalogue Archive operated by the United States Naval Observatory, Flagstaff Station (http://www.nofs.navy.mil/data/fchpix/).  This work has also made use of the VizieR online database of astronomical catalogues \\citep{ochsenbein00}."
    },
    "0512/astro-ph0512397_arXiv.txt": {
        "abstract": "{Several items on the diagnostics and interpretation of coronal loop observations are under debate.} {In this work, we analyze a well-defined loop system detected in a time-resolved observation in several spectral bands to study how far one can go in characterizing the loop structure and evolution.  } {The dataset includes simultaneous sequences of images in the 171 \\AA, 195 \\AA~and 284 \\AA~filter bands of TRACE, and in one filter of Yohkoh/SXT, with a time coverage of about 2.5 hours, and two rasters taken with SoHO/CDS in twelve relevant lines, forming between $\\log T \\approx 5.4$ (O V 629 \\AA) and $\\log T \\approx 6.4$ (Fe XVI 360 \\AA). The loop is initially best visible in the TRACE 195 \\AA~filter band, with some correspondence with the simultaneous SXT images, and later in the 171 \\AA~filter band, with good correspondence with the CDS raster images in the lines with formation temperature around $\\log T \\approx 6.0-6.1$.  We have taken as pixel-by-pixel background the latest TRACE, Yohkoh and CDS images where the loop has faded out. We examine the loop morphology evolution, the light curves, the TRACE filter ratio distribution and evolution, the images and emission measure from the CDS spectral lines.} {Our analysis detects that, after background subtraction, the emission along the loop and its evolution are non-uniform, especially in the 171 \\AA~filter band, and that the TRACE 195/171 filter ratio has a moderately non-uniform distribution along the loop and evolves in time. Both the light curves and the filter ratio evolution indicate a globally cooling loop. Relatively hot plasma may be present at the beginning while, during the first CDS raster, the data indicate a rather moderate thermal structuring of the loop.} {Our data analysis supports a coherent scenario across the different bands and instruments, points out difficulties in diagnostic methods and puts quantitative basis for detailed forward modeling.} ",
        "introduction": "As pointed out by the first X-ray images at high resolution  (e.g. Vaiana et al. 1973), coronal loops are the building blocks of the X-ray luminous solar corona. Hot coronal loops are known to be steady and stable over time scales longer than the characteristic plasma cooling times, and equilibrium scaling laws ruling some physical conditions (i.e. the maximum temperature, the pressure and the heating rate) generally hold (Rosner et al. 1978).  The high space and time resolution of the {\\it Transition Region And Coronal Explorer} (TRACE, Handy et al. 1999) telescope has brought new insight into the structure of the solar corona, and allowed to address new issues driven by the detection, for instance, of the filamentary structure of the coronal loops (Schrijver et al. 1999, Reale \\& Peres 2000) and of oscillations and propagating waves (e.g. De Moortel et al. 2000, Nakariakov \\& Ofman 2001).  In parallel to such achievements, the interpretation of ``conventional\" loop structures as observed with TRACE is under debate.  To summarize, already the first observations with TRACE detected steady loops which showed to have a ratio of the fluxes in the 195 \\AA~and 171 \\AA~filter passbands almost constant along the loop (Lenz et al. 1999). The ratio of the emission in different passband filters is typically used as a temperature indicator.  Unfortunately the temperature diagnostics using TRACE filters are particularly difficult, because the functions linking the ratios to the temperature are multi-valued. Nevertheless, if one assumes that most of the emission comes from plasma in the temperature range of maximum filter sensitivity, the function can be inverted and what is typically obtained for TRACE loops is that they are almost isothermal, much more than predicted by standard static loop models, and also overdense with respect to static loops at $\\sim 1$ MK.  Was it a new class of loops?  Soon after, it was shown that an alternative interpretation is possible: bundles of thin strands, each behaving as a single ``standard\" loop, convolved with the TRACE temperature response could appear as a single almost isothermal loop (Reale \\& Peres 2000). Next to this, another possibility has been invoked that long loops result to be mostly isothermal if heated at their footpoints (Aschwanden et al.  2001). This model suffers from the problem that footpoint-heated loops have been proven to be thermally unstable (e.g. Serio et al. 1981)  and therefore cannot be long-lived, as instead observed. A further alternative is to explain observations with steady non-static loops, i.e. with significant flows inside (Winebarger et al. 2001, 2002). Also this hypothesis does not seem to answer the question (Patsourakos \\& Klimchuk 2004). One of the reasons why the situation is so unclear is inherent in the data: although sometimes bright and well defined, the loops under analysis are always surrounded by other bright structures, which often intersect them along the line of sight.  Moreover, a uniform diffuse background emission also affects the temperature diagnostics, by adding systematic offsets which alter the filter ratio values. This problem emerged dramatically when the analysis of the same loop structure observed with the {\\it Soft X-ray Telescope} (SXT, Tsuneta et al. 1991) on board Yohkoh (Ogawara et al. 1991) led to three different results depending mostly on the different ways to treat the background (Priest et al. 2000, Aschwanden 2001, Reale 2002). One obvious way to check for the validity of the data analysis and of the loop diagnostics is to compare data from imaging instruments to simultaneous and cospatial data obtained from spectroheliographic instruments like the {\\it Coronal Diagnostic Spectrometer} (CDS, Harrison et al. 1995) on board SoHO (Domingo et al. 1995).  A loop observed on the solar limb with SoHO/CDS was analyzed by Schmelz et al. (2001), who found that whereas single line ratios tend to yield flat temperature distributions along the loop, a careful reconstruction of the emission measure distribution vs temperature (DEM) at selected points along the loop shows that this may not be a realistic result. A whole line of works started from this study reconsidering and questioning the basic validity of the temperature diagnostics with TRACE and emphasizing once again the importance of the background subtraction, but also the need to obtain accurate spectral data (Schmelz 2002, Martens et al. 2002, Aschwanden 2002, Schmelz et al. 2003). Similar results but different conclusions are reached by Landi \\& Landini (2004), and Landi \\& Feldman (2004) who analyze a loop observed with SoHO and, finding it nearly isothermal, consider this evidence as real and invoke a non-constant cross-section to explain it. From their analysis of SoHO/CDS data compared to other similar analyses made by other authors, Schmelz et al. (2005) propose that  there may be two different classes of loops, multi-thermal and isothermal, while Aschwanden \\& Nightingale (2005) analyze the thinnest loop structures detected with TRACE and find that a few are isothermal along the line of sight and may therefore be elementary loop components.  Another puzzling issue, certainly linked to the loop isothermal appearance, is the loop overdensity. In order to explain both these pieces of evidence, several authors claim that the loop cannot be at equilibrium and it must be filamented and cooling from a hotter state, probably continuously subject to heating episodes (nanoflares, Warren et al. 2002, Warren et al. 2003, Cargill \\& Klimchuk 2004).  The presence of nanoflares might explain the presence of coronal loops, stable although heated at the footpoints and with a peaked distribution of emission measure, as observed in active stars (Testa et al. 2004). An anticoincidence between hot and cooler loops has been found from the comparison of simultaneous Yohkoh and TRACE data (Nagata et al. 2003, Schmieder et al. 2004), who, however, find that the DEM of loops have a moderate but finite width.  Time-dependent modeling of one coronal loop observed with TRACE pointed out that the detailed description of the evolution of this loop requires a heating located at intermediate position between the apex and the footpoints, probably initially high and then slowly decaying (Reale et al. 2000).  The current debate in the interpretation of coronal loop observations points out the presence of intrinsic limitations in the  information that one can derive from present-day data.  Here, we take the analysis of a multi-wavelength observation of a time-evolving coronal loop as a guide to study how deep one can go in the  diagnostics and characterization of the loop, and puts  the basis for further analysis through detailed forward modeling which we leave for a future work. To this purpose we have searched for the observation of a loop in particularly good conditions for analysis: a simple and well-defined system, as isolated as possible, imaged in more than one TRACE filter band, in several SoHO/CDS spectral lines and with Yohkoh/SXT, and for a time period of more than one hour. Its eventual disappearance allows us to use the last images as point-to-point background to be subtracted. We try to use the coherence of the structure, its evolution and the spectral data to extract the maximum possible information from the data.  The selection of the loop observation, its description and the methods of the data analysis are illustrated in Section~\\ref{sec:data}, the results of the analysis are shown in Section~\\ref{sec:res} and they are discussed in Section~\\ref{sec:disc}; we draw our conclusion in Section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this work we use a good combination of multi-wavelength data, time/space/spectral resolution and signal-to-noise ratio to investigate the how deep one can go in the direct diagnostics and interpretation of a coronal observation. The observation of a specific structure, a coronal loop, and of its evolution helps us in the attempt to find coherent results and limitations to the information that it is possible to derive.  The analysis of the collected information shows that it is overall possible to obtain a coherent scenario and to  detect several details of the emission and evolution of a coronal loop. It confirms temperature diagnostics with TRACE to be difficult, and a proper subtraction of the high background critical for it, but also indicates that, in particular conditions, some sound information can be obtained. Spectroscopic data from SoHO/CDS provide useful complementary information, constraining the loop thermal structure, although in the limit of lower temporal and spatial resolution.  The coherent scenario that we obtain across bands and instruments appears to confirm the overall evolution of the coronal loops entirely visible with TRACE (Warren et al. 2002, 2003) and the presence of thermal structuring (Schmelz et al. 2005), but also adds several qualitative and quantitative details and puts several constraints to be matched coherently through detailed loop modeling."
    },
    "0512/astro-ph0512504_arXiv.txt": {
        "abstract": "We review some modern applications of the theory of few-body encounters between binaries and single stars.  In particular we focus on the treatment of adiabatic encounters, in a regime which is of importance in encounters between a star and a planetary system in a star cluster. ",
        "introduction": "Roughly speaking, gravitational scattering will be defined as the study of few-body encounters in which particles interact by Newtonian gravity, with certain types of initial conditions; namely, a few (normally two or three) bound subsystems, such as single stars or binaries, approach from infinity.  Then the problem is to characterise the outcome, usually in a statistical sense.  The new book by Valtonen \\& Karttunen (2006) will very quickly become the standard reference on this problem. While approximate analytical methods yield useful results in some limiting situations (see, for example, Sec.\\ref{sec:adiabatic}), computer simulation is an essential tool.  And while efficient codes exist for computing individual scattering encounters (e.g. {\\sl triple}\\footnote {\\tt http://www.ast.cam.ac.uk/$\\sim$sverre/web/pages/nbody.htm}, {\\sl fewbody}\\footnote{\\tt http://www.astro.northwestern.edu/$\\sim$fregeau/code/fewbody/}) the scattering packages in {\\sl starlab}\\footnote{\\tt http://www.ids.ias.edu/$\\sim$starlab/} provide the additional, extremely valuable functionality of efficiently sampling parameter space.   Many possibilities exist for graphic rendering of individual encounters, such as {\\sl GLanim}\\footnote{\\tt http://grape.astron.s.u-tokyo.ac.jp/$\\sim$makino/softwares/GLanim/}.  \\subsection{Examples and applications}  Though the topic was developed throughout the twentieth century, it remains topical, because of fresh applications, such as the following.  \\subsubsection*{Scattering of normal stars by a binary black hole}  As a result of mergers, central binary black holes are expected in many galaxies.  Observationally, binary black holes are studied at high energies (e.g. NGC 6240, studied with Chandra: Komossa et al 2003) and at visual wavelengths (e.g. the famous Tuorla object OJ287: Sillanpaa et al 1988).  A problem of long standing is the evolution of the relative orbit of the black holes, as they scatter stars from the surrounding galaxy (Hills 1983, Gould 1991b, Mikkola \\& Valtonen 1992, Fukushige et al 1992, Quinlan 1996, Zier \\& Biermann 2001, Merritt 2001, 2002, Yu \\& Tremaine 2003, Milosavljevi\\'c \\& Merritt 2003, Chatterjee et al 2003, Gualandris et al 2005).   \\subsubsection{Evolution of planetary systems in star clusters} \\label{subsec:clusters} For the solar system, though it is not now in a star cluster, the question of stellar perturbations has been considered for a long time (e.g. Lyttleton \\& Yabushita 1965).  In recent years this question has arisen because of the initially surprising absence of planets (searched for photometrically) in star clusters (e.g. the globular star cluster 47 Tuc; see Gilliland et al 2000, Weldrake et al 2005).  From the theoretical point of view the study of such scattering is simplified by the fact that one component of the participating binary is effectively massless (Davies \\& Sigurdsson 2001, Spurzem et al 2003, Fregeau et al 2005).  We shall say more about this in Sec.\\ref{sec:adiabatic}  \\subsubsection{Capture of exotic particles by multiple systems} The idea here is that the flux caused by interactions with normal matter may be enhanced in situations where the particles can be trapped gravitationally, which allows them multiple opportunities of interacting.  This is a more speculative problem, but has still attracted considerable interest (Press \\& Spergel 1985, Gould 1987, 1988, 1991a, Damour \\& Krauss 1998, 1999,  Gould \\& Alam 2001, Lundberg \\& Edsj\\\"o 2004).  Theoretically it is the case where the incoming particle is nearly massless.  The complexity of the problem is illustrated by a comparable problem of solar system dynamics: what fraction of comets and asteroids are destroyed by colliding with the sun?  Several planetary resonances are involved in the fact that almost half of a sample of near-Earth asteroids end by colliding with the sun (Farinella et al 1994).  It is also a common fate of short-period comets (Levison \\& Duncan 1994).  \\subsubsection{The M4 triple}  The nearby globular cluster M4 contains a triple system in which a distant companion, of mass comparable with or somewhat larger than Jupiter's mass, is in orbit about a binary consisting of a white dwarf and a millisecond pulsar.  It is likely to have formed in a four-body scattering encounter between two binaries (Rasio et al 1995, Ford et al 2000, Fregeau et al 2005). ",
        "conclusions": ""
    },
    "0512/astro-ph0512218_arXiv.txt": {
        "abstract": "We describe the effects of weak gravitational lensing by cosmological large scale structure on the diffuse emission of 21 centimeter radiation from neutral hydrogen at high redshifts during the era of reionization. The ability to observe radial information through the frequency, and thus three-dimensional regions of the background radiation at different redshifts, suggests that 21 cm studies may provide a useful context for studying weak lensing effects. We focus on the gravitational lensing effects on both the angular power spectra and the intrinsic, three-dimensional power spectra. We present a new approach for calculating the weak lensing signature based on integrating differential Fourier-space shells of the deflection field and approximating the magnification matrix.  This reduces the problem of calculating the effect and higher order corrections to solving coupled systems of differential equations.  This method is applied to reionization models of the 21 cm background spectra up to small angular scales over a range in redshift. The effect on the angular power spectrum is typically less than one percent on small angular scales, and very small on scales corresponding to the feature imprinted by reionization bubbles, due to the near-scale invariance of the angular power spectrum of the 21 cm signal on these scales. We describe the expected effect of weak lensing on three-dimensional 21 cm power spectra, and show that lensing creates aspherical perturbations to the intrinsic power spectrum which depend on the polar angle of the wavevector.  The effect on the three-dimensional power spectrum is less than a percent on scales $k \\lesssim 0.1 \\text{ h Mpc}^{-1}$, but can be $\\gtrsim 1\\%$ for highly inclined modes for $k \\gtrsim 1 \\text{ h Mpc}^{-1}$.  The angular variation of the lensing effect on these scales is well described by a quartic polynomial in the cosine of the polar angle. The detection of the gravitational lensing effects on 21 cm power spectra will require very sensitive, high resolution observations by future low-frequency radio arrays. ",
        "introduction": "After recombination at the surface of last scattering of the cosmic microwave background (CMB) at $z = 1100$, most of the baryons in the universe are in the form of neutral hydrogen.  Soon after the first galaxies and stars formed, they began to reionize the intergalactic medium.  Understanding the astrophysical processes through which this happened is a major challenge in modern cosmology.  Observations of the Gunn-Peterson trough in high-redshift quasars suggest that the universe was reionized by redshift $z \\sim 6$ \\citep[e.g.][]{fan02}.  Meanwhile, the large optical depth to Thomson scattering of CMB photons observed by \\textit{WMAP} provides an integral constraint on reionization, indicating that reionization began before $z \\gtrsim 14$ \\citep{spergel03}.  At face value, the evidence suggests that reionization was a long and complicated process; it is a topic of intense current research.  Since neutral hydrogen radiates in the rest frame at a wavelength of 21 cm due to the ``spin-flip'' hyperfine transition, 21 cm tomography has been suggested as a very promising probe of the history of reionization.  This radiation should be observable at earth at low radio frequencies after being cosmologically redshifted.  This 21 cm radiation from the reionization epoch should carry immensely rich information about the reionization process and the formation of structure. Furthermore, since by scanning through frequency one effectively is scanning through different redshifts of 21 cm emission, one has the remarkable ability to probe the evolution of the neutral hydrogen density with time.  Early work studying the potential of 21 cm observations to probe the high-redshift universe was done by, e.g. \\citet{scottrees90,madaumeiksinrees97} and \\citet{tozzi00}.  Recent research has focused on the expected statistical signal of the 21 cm intensity from the epoch of reionization (EoR) \\citep{zfh04, moraleshewitt04, bharadwajali05}.  It was proposed by \\citet{zfh04} that one may study the angular power spectra of brightness temperature fluctuations at different frequencies (and cosmic epochs) in order to probe the details of the reionization history.  An alternative approach emphasizes the full three-dimensional nature of the fluctuations \\citep{moraleshewitt04}.  Angular observations of the 21 cm sky stepped through different frequencies can be combined to form an image cube by mapping frequency to line-of-sight distance, and such three-dimensional image cubes can be used to probe the statistical properties of the full three-dimensional distribution of HI fluctuations at a particular epoch during reionization. The simple three-dimensional structure of the intrinsic fluctuations can be harnessed to remove foregrounds with distinct symmetry properties. Data analysis techniques applying this approach have been developed by \\citet{moraleshewitt04, morales05, bowmanmoraleshewitt05} and \\citet{ moralesbowmanhewitt05}.  Detection of the 21 cm signal from high redshift is among the major goals of planned future low frequency radio arrays, including the Mileura Widefield Array (MWA), the Primeval Structure Telescope (PAST), the Low Frequnecy Array (LOFAR), and the Square Kilometer Array (SKA).  The expected statistical signal detected by future 21 cm observations will be of a complicated nature, resulting from a combination of cosmological sources and astrophysical processes.  The intrinsic 21 cm signal is determined by fluctuations in density, neutral fraction, and the spin temperature describing the relative populations of HI spin states, which is affected by the background radiation field. \\citet{fzh04a} and \\citet{fzh04b} have shown how inhomogenous reionization will imprint a shoulder feature on the power spectra; hence the power spectra can be an effective discriminant between possible reionization scenarios. \\citet{bharadwajali04} have pointed out the importance of the effect of peculiar velocities, which introduce redshift space distortions and are caused by the infall towards (or outflow from) density perturbations.  \\citet{barkanaloeb05} have described how these peculiar velocity fluctuations lead to anisotropies in the intrinsic three-dimensional power spectrum $P(\\bm{k})$ with characteristic angular dependences on the orientation of $\\bm{k}$.  These anisotropies may be exploited to separate out the effects of the density fluctuations and astrophysical processes.  The effects of radiative processes through the spin temperature have recently been considered by \\citet{pritchardfurlanetto05} and \\citet{hirata05}.  These sources of fluctuations together determine the intrinsic power spectrum, and hence the autocorrelation function,  at any given epoch (or redshift).  As the 21 cm photons emitted by HI stream through the universe, their paths are perturbed stochastically by density inhomogeneities, and their observed directions in the sky are altered.  Hence, gravitational lensing by large scale structure presents another source of fluctuations and modifies the intrinsic power spectra and correlation functions.  These weak lensing modifications are best thought of as secondary fluctuations; they are determined by density fluctuations along the photon path between emission and observation.  Furthermore, they act upon existing fluctuations: without intrinsic brightness fluctuations characterized by the physics of the EoR, lensing would have no effect.  The effect of weak gravitational lensing on diffuse backgrounds has been studied extensively in the context of the CMB.  The effects on the temperature and polarization power spectra have been discussed in \\citet{seljak96,zaldarriagaseljak98,hu00,cooray04} and \\citet{challinorlewis05}.  Furthermore, the methods of using weak lensing to reconstruct intervening mass distributions has been studied in, e.g. \\citet{zaldarriagaseljak99,hu2001a,hu2001b} and \\citet{hirataseljak03}.  Future observations of 21 cm backgrounds will present a new source of diffuse emission that will be subject to weak lensing by large scale structure.   In contrast to the CMB, which presents a single, narrow surface at one redshift, the 21 cm emission from different redshifts provides multiple source planes that may be observed in order to study lensing effects.  Indeed, the three-dimensional nature of 21 cm tomography provides a new and interesting context for lensing sudies.  Some initial steps have been taken in this direction.  \\citet{pen04} has suggested a method to extract information on the projected lensing convergence from future 21 cm surveys.  In this paper we focus on the expected weak gravitational lensing effects on the angular and three-dimensional power spectra of 21 cm emission from the era of reionization.  In \\S \\ref{diffeqsection} we motivate and derive a new differential method for calculating the lensing modifications.  Additionally the lensing effects in certain limiting cases are explored.  In \\S \\ref{clresults}, we apply these methods to models of the 21 cm angular power spectra. In \\S \\ref{3dlensingsection}, we discuss the lensing distortions to the three-dimensional power spectra of 21 cm fluctuations, and show that they produce asymmetric perturbations.  These effects have not been considered in previous work.  In \\S \\ref{detectsection} we compare the magnitude of these effects to expected sensitivites of a future SKA experiment. Finally, we conclude in \\S \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper, we have addressed the effects of gravitational lensing of large-scale structure on the brightness temperature power spectrum of high redshift 21 cm.  Gravitational lensing presents a secondary source of fluctuations caused by density perturbations between emission and observation, and reprocesses intrinsic fluctuations caused by density, neutral fraction, and spin-temperature fluctuations in the EoR.  We have presented a new formulation for calculating the gravitational lens effect on temperature power spectra of diffuse backgrounds.  This harmonic-space approach leads to a hierarchical system of equations that describe the lensing effect on the power spectrum as one integrates differential shells of lensing modes regulated by a single cutoff scale $\\Lambda$.  We have described the general structure of the system of equations, which can be closed at the chosen level by approximating the magnification matrix.  This method is an alternative to previous series solutions, which require evaluation of high-dimensional integrals or multiple transforms between position space and harmonic space.  We have applied this method to the case of 21 cm fluctuations from the era of reionization.  The effect on two-dimensional spectra $C_l$ at a single frequency (redshift) is typically less than $1 \\%$ on scales $100 < l < 10^5$.  In particular we find that higher-order corrections are negligible on all scales despite the dominance of small scale power.  At late times when the spectra are dominated by the ionized bubbles, the flattening of the power spectrum on small scales suppresses the lensing effect due to near scale-invariance of the spectrum.  For small surveys, the three-dimensional power spectrum $P(\\bm{k})$ can be directly probed by mapping frequency to line-of-sight distance.  We have shown how the expected lensing effect on the three-dimensional power spectrum can be reduced to a two-dimensional calculation for each set of temperature modes with the same line-of-sight component wavenumber $k_\\parallel$.  Since the relevant quantity is the transverse gradient power spectrum, which is different for each $k_\\parallel$, gravitational lensing will affect transverse wavevectors differently from line-of-sight wavevectors.  The result is a dependence of the lensed fraction on the angle between $\\bm{k}$ and the line-of-sight that is well described by a quartic polynomial in the cosine.  Thus, weak lensing generates angular components of the power spectrum with amplitudes of order one percent or more on small scales, $k \\gtrsim 1 \\text{ h Mpc}^{-1}$.  Finally, we have compared the predicted lensing signatures to the potential sensitivity limits of SKA observations of the 21 cm emission.  Future experiments will likely need greater sensitivity and resolution in angular and frequency space in order to detect the effects of weak lensing.  The gravitational lensing effects on power spectra will become important for futuristic high-precision 21 cm experiments that can measure the power spectrum to better than one percent and can probe scales $k \\gtrsim 1 \\text{ h Mpc}^{-1}$."
    },
    "0512/gr-qc0512022.txt": {
        "abstract": "We apply the method of matched asymptotic expansions to analyse whether cosmological variations in physical `constants' and scalar fields are detectable, locally, on the surface of local gravitationally bound systems such as planets and stars, or inside virialised systems like galaxies and clusters. We assume spherical symmetry and derive a sufficient condition for the local time variation of the scalar fields that drive varying constants to track the cosmological one. We calculate a number of specific examples in detail by matching the Schwarzschild spacetime to spherically symmetric inhomogeneous Tolman-Bondi metrics in an intermediate region by rigorously construction matched asymptotic expansions on cosmological and local astronomical scales which overlap in an intermediate domain. We conclude that, independent of the details of the scalar-field theory describing the varying `constant', the condition for cosmological variations to be measured locally is almost always satisfied in physically realistic situations. The proof of this statement provides a rigorous justification for using terrestrial experiments and solar system observations to constrain or detect any cosmological time variations in the traditional `constants' of Nature.  PACS Nos: 98.80.Es, 98.80.Bp, 98.80.Cq $\\ $ ",
        "introduction": "In recent years there has been growing interest in the observational and theoretical consequences of time variations in the values of the traditional constants of Nature, notably of the fine structure constant, $\\alpha $, \\cite% {webb, chand, sdss, qu, lev, lev2, rocha, darl, oh, drink}, the electron-proton mass ratio, $\\mu =m_{e}/m_{pr}$, \\cite{ubachs, petit, tz}, and Newton's `constant' of gravitation, $G$, \\cite{bert}. In all cases the experimental evidence that can be brought to bear on the problem is a combination of local (laboratory, terrestrial, and solar system) and global (astronomical and cosmological) observations \\cite{uzan, olive, jdb, jdbroysoc, posp}. The first theoretical challenge is to develop self-consistent extensions of general relativity which incorporate varying `constants' rigorously by including the gravitational effects of the variations and ensuring that energy and momentum are totally conserved by the variations that replace the former constants. This is achieved by regarding the `constant' as a scalar field with particular couplings and self-interaction. In the case of varying $G$, the Brans-Dicke theory \\cite% {bd, bkm} provides the paradigm for a scalar-tensor theory of this type. Recently, the same philosophy has been applied to produce simple extensions of general relativity which self-consistently describe the spacetime variation of $\\alpha $, \\cite{bek, bsbm}, and $\\mu $, \\cite{bmmu}. It is also possible to extend these studies to include simultaneous variations of several constants, and to include the weak coupling by a generalisation of the Weinberg-Salam theory to include spacetime varying coupling `constants' \\cite{bs, sb}. In these theories the analysis of the behaviour of their solutions is simplified because we know that the allowed variations in constants like $\\alpha $ are constrained already to be `small' and will not have any significant effect on the expansion dynamics of the universe in recent times. The Brans-Dicke theory is different. Small variations in $G$ will always have direct consequences for the expansion dynamics of the universe. Typically, a power-law time variation of $G\\propto t^{-n}$ creates a variation of the expansion scale factor that goes as $a(t)\\propto t^{(2-n)/3}$ in a dust-dominated Friedmann universe \\cite{newtgrav}.  These theories are confronted with a variety of laboratory, geochemical, and astronomical observations. In the case of variations of $\\alpha $ we have laboratory constraints on atomic lines, indirect bounds from the Oklo natural reactor operation 1.8 billion years ago, radioactive decay products in meteoritic data back to 4.6 billion years ago, and quasar spectra out to redshifts $z\\lesssim O(4)$, as the prime sources of observational evidence against which to test theories which permit spacetime variations \\cite{uzan, olive, jdb, jdbroysoc, posp}. Generally, the data from all these diverse physical scales are lumped together and used to test time variations of $% \\alpha $. Thus laboratory or solar system evidence is compared directly with quasar data and used to constrain the allowed cosmological variations of $% \\alpha $. Similar tactics are used to constrain the allowed variations of other constants, like $G$ or $\\mu $.  This simultaneous use of terrestrial and astronomical bounds on constants assumes implicitly the unproven requirement that any variation of a constant on cosmological scales is `seen' locally inside virialised structures like galaxies or solar systems, and has a measurable effect in laboratory experiments on Earth. It is not obvious a priori that this need be the case: we would not expect to test the expansion of the universe by measuring the expansion of the Earth. The central question that this paper addresses is the extent to which global variations of `constants' on cosmological scales that take part in the Hubble expansion of the universe are seen locally on the surface of gravitationally-bound structures, like planets, or inside bound systems of stars like galaxies \\cite{bmot}. Only if we can show that cosmological variations have calculable local effects will it be legitimate to use laboratory and solar system observations to constrain theories of varying constants in the way that is habitually done, without proof, in the literature. So far, detailed analyses of spatial variations of constants have been made only for small variations, where the isotropy of the microwave background places very strong limits on spatial variations because of the associated Sachs-Wolfe effects created by the gravitational potential perturbations that accompany spatial fluctuations in `constants' via their associated scalar fields because of the coupling of the latter to matter \\cite{jbspace1, jbspace}.  We are concerned with the dynamics of spacetime scalar fields that are weakly coupled to gravity and matter. We will not consider theories where there are two or more scalar fields interacting amongst themselves, although the method we use here could also be easily extended to that scenario. Theories which introduce varying constants self-consistently into Einstein's conception of a gravitation theory do so by associating the `constant' $\\mathbb{C}$ to be varied with a new scalar field, so $\\mathbb{C} \\rightarrow \\mathbb{C}(\\varphi )$. The variations of this scalar field gravitate and contribute to the curvature of spacetime like any other form of mass-energy. They must also conserve energy and momentum and so their forms are constrained by a covariant conservation equation for the scalar field. Typically, this results in a wave equation of the form \\begin{equation} \\square \\varphi =\\lambda f(\\varphi )L(\\rho ,p),  \\label{gen} \\end{equation}% where $\\varphi $ is a scalar field associated with the variation of some `constant' $\\mathbb{C}$ via a relation $\\mathbb{C}=f(\\varphi )$,$\\lambda $ is a dimensionless measure of the strength of the space-time variation of $\\mathbb{C}$, $f(\\varphi )$ is a function determined by the definition of $\\varphi ,$ and $L(\\rho ,p)$ is some linear combination of the density, $\\rho $, and pressure, $p$, of the matter that is coupled to the field $\\varphi $ and $% f(\\varphi )\\simeq 1$ for small variations in $\\varphi $ and $\\mathbb{C}$. This form includes all the standard theories for varying constants, like $% G$, $\\alpha $, and $\\mu $, of refs. \\cite{bd, bek, bsbm, bmmu}.  We shall refer to our scalar field as the `dilaton', denote it by $\\phi $, and analyse the form of the standard equation (\\ref{gen}) a little further by assuming that $\\phi (\\vec{x},t)$ satisfies a conservation equation that can be decomposed into the form \\begin{equation} \\square \\phi =B_{,\\phi }(\\phi )\\kappa T-V_{,\\phi }\\left( \\phi \\right) \\label{cons} \\end{equation}% where $T$ is the trace of the energy momentum tensor, $T=T_{\\mu }^{\\mu }$ (with the contribution from any cosmological constant neglected). We absorb any dilaton-to-cosmological constant coupling into the definition of $V(\\phi )$. The dilaton to matter coupling $B(\\phi )$ and the self-interaction potential, $V(\\phi )$, are arbitrary functions of $\\phi $ and $\\kappa =8\\pi G $ and $c=\\hslash =1$. This covers a wide range of theories which describe the spacetime variation of `constants' of Nature, and includes Einstein-frame Brans-Dicke (BD) and all other, single field, scalar-tensor theories of gravity. In cosmologies that are composed of dust, cosmological constant and radiation it will also contain the Bekenstein-Sandvik-Barrow-Magueijo (BSBM) of varying $\\alpha $, \\cite{bsbm}, and other (single dilaton) theories which describe the variation of standard model couplings, \\cite{posp}. Note that one could, for example, generalise the form of this conservation equation, (\\ref{cons}), whilst maintaining relativistic invariance, by adding a coupling to $\\sqrt{T^{\\alpha \\beta }T_{\\alpha \\beta }}$, or we could also break local Lorentz invariance by adding extra couplings to the pressure, $P$, defined w.r.t. to some preferred coordinate system. In this paper we will mostly be considering spacetimes in which the pressure of the matter vanishes, $P=0$, and so all of the potential extra couplings mentioned above will reduce to a form that is included in the conservation equation (\\ref{cons}) that we assume.  We will assume that the background cosmology is isotropic and homogeneous and work with a Friedmann-Robertson-Walker (FRW) background metric: \\begin{equation} \\mathrm{d}s^{2}=\\mathrm{d}t^{2}-a^{2}(t)\\left( \\frac{\\mathrm{d}r^{2}}{% 1-kr^{2}}+r^{2}\\{\\mathrm{d}\\theta ^{2}+\\sin ^{2}\\theta \\mathrm{d}\\phi ^{2}\\}\\right) ,  \\label{frw} \\end{equation}% where $k$ is the spatial curvature parameter. In the background universe the dilaton field is also assumed to be homogeneous (so $\\phi =\\phi _{c}(t)$) and therefore satisfies the ordinary differential equation (ODE): \\begin{equation} \\frac{1}{a^{3}}\\left( a^{3}\\dot{\\phi}_{c}\\right) ^{\\cdot }=B_{,\\phi }(\\phi _{c})\\kappa \\left( \\epsilon _{c}-3P_{c}\\right) -V_{,\\phi }\\left( \\phi _{c}\\right) .  \\label{dilcons} \\end{equation}  \\ \\ The dilaton conservation equation also reduces to a similar ODE if we are only interested in its time-independent mode $\\phi (r$) in a static spherically symmetric background, like the Schwarzschild metric. In the actual universe, however, spacetimes that look static in some locality must still match smoothly on to the cosmological background spacetime on large scales. Hence, if we are to model the evolution of the dilaton field accurately in some inhomogeneous region, embedded in a homogeneous background, we must demand that at large distances $\\phi \\rightarrow \\phi _{c} $ in some appropriate way. Even if spherical symmetry is assumed, the conservation equation is generically a non-linear, second-order, partial differential equation (PDE) and its solution is far from straightforward to find, either exactly or approximately. Even numerical models are technically difficult to set-up, see ref. \\cite{harada}, and only allow us to consider one particular choice of $B(\\phi )$, $V(\\phi ),$ and the spacetime background at a time. However, as we have mentioned above, if we are to bring all of our experimental evidence to bear on these models, we need to know how to interpret local observations in the light of our cosmological ones. Perhaps the most important piece of knowledge we would like to have is the correlation between the local and global (or cosmological) time variation of the dilaton. In particular, we would like to find the criteria under which it is true that \\begin{equation} \\dot{\\phi}(\\mathbf{x},t)\\approx \\dot{\\phi}_{c}(t),  \\label{wettcond} \\end{equation}% that is, when does the local time-variation of $\\phi $ track the cosmological one?  In this paper we will try to answer this question by applying asymptotic methods commonly used in fluid dynamics in order to construct asymptotic approximations to the behaviour of $\\phi $ close to a spherical static mass, that match to the cosmological solution, $\\phi _{c}$, at large distances. From this analysis, we can derive a sufficient condition for eq. (\\ref% {wettcond}) to hold. We limit ourselves in this paper to considering spacetime backgrounds that are a spherical symmetric. In a subsequent work we will generalise our results to deal with more general non-spherical backgrounds, \\cite{shawbarrow2}.  Throughout our analysis we will refer to the spherical, static mass as a `star', however it could be taken to be a planet (e.g. the Earth), a black-hole, or even a galaxy or cluster of galaxies. We are mostly interested in the (realistic) case where the surface of our `star' lies far-outside its own Schwarzschild-radius. By applying our results to the black-hole case, however, we will comment on the problem of `gravitational memory'; that is, is the cosmological background value of $\\phi $ on the horizon at the time when a black hole forms, frozen-in, or `remembered', when the black-hole forms, or does it continue to track the background cosmological evolution?  This paper is organised as follows: in section II we review some of the previous studies into the problem of local vs. global dilaton evolution and note where our work extends and improves these studies. In section III we will introduce the method of matched asymptotic expansions which we will use to carry out our study, and provide some simple examples of its application. In section IV we define our geometrical set-up of a star in a background cosmology, and detail our particular choices of possible spacetime backgrounds. Then, in section V, we construct overlapping asymptotic expansions to study the constraints required if the local and global evolution is to match together. In section VI we derive the conditions that must be satisfied for our method to be valid. Taking these into account, we interpret and generalise our results in section VII. In this section we shall also make a conjecture about a general condition that is sufficient for eq. (\\ref{wettcond}) to hold, which will apply to more general spacetime backgrounds that those explicitly considered in this paper. Finally, in section VIII we use our sufficient condition to show that we \\emph{do} expect eq. (\\ref{wettcond}) to hold here on Earth. ",
        "conclusions": "In this paper we have considered the extent to which a cosmological time variation in a scalar field (dilaton) would be detectable on the surface of gravitational-bound systems that are otherwise disconnected from changes that occur over cosmological scales. This problem is of particular relevance when the dilaton defines the local value of one of the traditional `constants' of Nature, and when the system is the Earth or our solar system. Several scalar-tensor theories have already been developed which self-consistently describe the variations of traditional `constants' of Nature, like $\\alpha $ and $G$. By matching rigorously constructed asymptotic expansions of the associated scalar field found in different limits, close to the Earth and far away from it at cosmological scales, we have been able to derive approximate, analytical expressions for the scalar field near the surface of such bound systems. This result was found under two major assumptions:  the physically realistic condition that the scalar field should be weakly coupled to matter and gravity(in effect the variations of `constants' on large scales occur more slowly than the universe is expanding )  and thus have a negligible back-reaction on the cosmological background, and the less realistic one of spherical symmetry. We do not expect the relaxation of the spherical symmetry condition to greatly alter the qualitative nature of our conclusions, and we shall present a rigorous treatment of the problem when no symmetry is present in a subsequent work.  Finally we have extracted from our analysis a sufficient condition for the local time-variation of a scalar field, or varying physical `constant', to track the cosmological one, and we have proposed a generalisation of this condition that is applicable to scenarios more general than those explicitly considered here. By evaluating the condition for an astronomical scenario similar to the one appropriate for our solar system, we have concluded that almost irrespective of the form of the dilaton-to-matter, and the form of the dilaton self-interaction, its time variation in the solar system will track the cosmological one. We have therefore provided a general proof of what was previously merely assumed: that \\emph{terrestrial} and \\emph{solar system} based observations can legitimately, be used to constrain the \\emph{% cosmological} time variation of supposed `constants' of Nature. \\newline  \\noindent \\textbf{Acknowledgements} DS is supported by a PPARC studentship. We would like to thank P.D. D'Eath and T. Clifton for helpful discussions. \\appendix"
    },
    "0512/astro-ph0512054_arXiv.txt": {
        "abstract": "{} { The distribution of RR Lyrae stars (RRLS) in the inner Large Magellanic Cloud (LMC), and the structure of the halo of the LMC delineated by these stars are studied here. } {RRLS identified by the OGLE II survey are used to estimate their number density distribution in the bar region of the LMC. To find their location, I estimated the scale-height of their distribution in the LMC using extinction corrected average magnitudes of ab type stars. } {The density is found to vary differently along and across the bar of the LMC, and the difference is found to be statistically significant. The density distribution is found to be elongated like the LMC bar and the position angle (PA) of the elongation is estimated to be 112.$^o$5 $\\pm$ 15.$^o$3. This value of PA is found to be same as the PA$_{maj}$ of the bar, within the errors, estimated using red clump stars and giants. The ellipticity of their density distribution is estimated to be $\\sim$ 0.5, very similar to the ellipticity of the bar, estimated from giants. The above results show that majority of the population of RRLS in the central region of the LMC are found to have the signature of the bar. This result could mean that most of these stars are located in the disk, considering the bar as a disk feature. On the other hand, their scale-height was found to be 3.0$\\pm 0.9$ kpc. This indicates that RRLS are located in the halo and not in the disk. } {Thus these stars in the inner LMC have halo-like location and a disk-like density distribution. I discuss some possible formation scenarios for this puzzling combination.} ",
        "introduction": "The Large  Magellanic Cloud (LMC) is known to be a disk galaxy with or without a halo. There have been many efforts to find evidence for the presence of a halo in the LMC. Such evidence is looked for in the old stellar population like the globular clusters (GCs) and the field RR Lyrae stars (RRLS). The oldest GCs appear to lie in a flat rotating disk whose velocity dispersion is 24 km s$^{-1}$ (Kinman et al. 1991 and Schommer et al. 1992). van den Bergh (2004) re-estimated the velocity dispersion of the GCs and concluded that they could still have formed in the halo. Since the number of GCs in the LMC is small (13 clusters), it is difficult to infer the signature of the halo from this sample. The other tracer is the RRLS, which are almost as old as the oldest GCs. Recent surveys of the LMC like the OGLE and MACHO have identified a large number of RRLS. Among the follow up studies, Minniti et al. (2003) found a kinematic evidence for the LMC halo, by estimating the velocity dispersion in the population of RRLS. Alves (2004) estimated the mass of the LMC based on their surface density and remarked that their exponential scale length is very similar to that of the young LMC disk. Freeman (1999) remarked that the similarity of their scale length with the scale length of the young stellar disk suggested that the RRLS in the LMC were disk objects, like the old clusters, and supported the view that the LMC may not have a metal-poor halo.  In this study, I used the RRLS catalogue of Soszynski et al. (2003) in the LMC, consisting of 7400 stars (after correcting for multiple entries), to obtain their number density distribution. The aim of the present study is to identify any feature in the distribution, that could be described as the signature of a halo or a disk distribution. Indications for a distribution which is not circularly symmetric in the inner regions can be seen in figure 12 of Alcock et al. (1996), where a large fluctuation in the RRLS density is found in the inner 3 Kpc. The above fluctuations could arise due to averaging in position angles (PAs), when there is a non-circular distribution. The density map of RRLS shown in figure 1 of Soszynski et al. (2003) indicated that the stars have a preferred distribution along the major axis of the bar, with a steep gradient across the bar. This is a strong indication that these stars may be located in an elongated bar like distribution. The RRLS as detected in the OGLE survey are projected on the bar, since the OGLE survey essentially traces the bar. If they belonged to the halo, then their distribution in the inner LMC should show a symmetric distribution. On the other hand, if their distribution showed differences along and across the bar and also any other property of the bar as found in the younger population, then the RRLS population in the inner LMC may belong to the bar or disk, since the bar is considered as a disk feature. This point is further explored in this study by comparing the distribution of RRLS with the distribution of red clump (RC) stars and red giant stars, where RC stars and giants belong to a younger population and also trace the feature of the LMC bar.  If the distribution of RRLS mimics that of the bar/disk, then the obvious conclusion will be that the stars are located in the disk of the LMC and not in the halo. This conclusion needs to be verified by finding the actual locations RRLS. This can be done either by estimating the individual distances to the stars or by estimating  their scale-height. The first option is quite tricky as it requires information on metallicity of all stars. The second option is doable with the existing data and the result obtained can discriminate their disk/halo degeneracy. Since the LMC is an almost face-on galaxy, it is a very tricky problem as the stars get projected on the disk even if they are distributed in a spheroid. In such cases, indirect methods are used to estimate the scale-height and this is such an attempt based on the magnitude dispersion of the ab type RRLS. The method was used by Clementini et al. (2003) to estimate the line of sight depth in their sample. I have attempted this method with a much larger sample, distributed over a larger area. The ab type RRLS from the above catalogue is used and their average magnitude and dispersion as a function of location over the surveyed region, after correcting for the interstellar extinction was estimated. The depth and the scale-height of their distribution can be estimated from the observed dispersion, after correcting for dispersion due to other factors. We also compared the estimated scale-height with that found for our Galaxy. The results obtained are discussed in the light of possible formation scenarios.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\caption{ Number density distribution of RR Lyrae stars in the LMC. The unit for the number density, d is stars/sq. degree. The colour code is explained in the figure, where magenta denotes locations of highest density and blue denotes locations of lowest density. \\label{fig1}} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig2.eps}} \\caption{ The normalised number density distribution of RRLS along (filled circles) the PA = 120$^o$ (along the bar, filled circle), 30$^o$ ( across the bar, open square), 75$^o$ (asterix) and  165$^o$ (open circle). The errors include statistical error and a 20\\% more due to any improper incompleteness correction. \\label{fig2}} \\end{figure} ",
        "conclusions": "The analysis presented here showed that the number density distribution of the RRLS differs along and across the bar, indicating a preferred elongation along the major axis of the bar. The bar feature as shown by RRLS is very similar to that shown by the younger population, in terms of PA$_{maj}$, ellipticity and variation of density along the major axis. This indicates that the bar feature as seen in the disk is also found in the RRLS. One of the possible interpretation of this result is that these stars are located in the disk and not in the halo and hence they mimic the bar feature. This interpretation is against the general belief that the majority of RRLS belongs to the halo population. We estimated the scale-height of the ab type RRLS and was found to be 3.1$\\pm $0.9 kpc, indicating that the majority are not located in the disk, but in the halo, which may be flattened. Therefore, the RRLS in the LMC are in fact located in the halo. Thus RRLS in the inner LMC have a strange combination of a disk-like density distribution and halo-like location. We also find marginal evidence for stretching of the LMC halo towards our Galaxy.  The scenario described by RRLS in the inner LMC seems to show the signatures of both disk as well as halo. This result is quite puzzling. This is similar to a situation where a spheroidal system, mimics the feature of a flattened disk-like system. At this stage it is not very clear as to what could be the origin of such a system. Detailed study of RRLS including their kinematics and metallicity is required to solve this puzzle.  Here we discuss some possibilities for obtaining such systems, in general. If RRLS were formed in the halo, then they will most probably have a symmetric density distribution with a non-disk like luminosity function. After it is formed, it is not possible to create the above imprints of the disk on an already existing spheroidal distribution. Thus this scenario is unlikely. If we assume that RRLS were initially formed in the disk, and then the distribution is quite likely to have the disk properties as obtained here. Such a distribution can get dislocated into a puffed up distribution. This would relocate the stars into the halo (away from the disk) but still maintaining the disk properties in terms of density and luminosity. Thus this scenario can explain the observed properties. This kind of situation can arise in mergers, where the disk gets heated up and results in a puffed up distribution. Jog \\& Chitre (2002) found some galaxies which have exponential light profile (like the disk) and kinematics indicating pressure support rather than rotation (puffed up location). Bournaud et al. (2004) found that mergers of intermediate mass galaxies (4:1 - 10:1) results in similar such peculiar systems. These type of mergers are more common  at high red shift as shown in the hierarchical merging models (Steinmetz \\& Navarro 2002) for galaxy formation. Combining all the above information, it is likely that the LMC went through mergers in the mass range as indicated above in its early formation. Recent interactions of the LMC with our Galaxy can induce tidal effects on the LMC, which has been already noticed in the disk by van der Marel (2001). The halo of the LMC can also get stretched towards our Galaxy due to tidal effects. Thus interaction and merger events may have resulted in the present distribution of RRLS."
    },
    "0512/astro-ph0512262_arXiv.txt": {
        "abstract": "Constraints on cosmological parameters from upcoming measurements with the Mileura Widefield Array---Low Frequency Demonstrator (MWA-LFD) of the redshifted 21 cm power spectrum are forecast assuming a flat $\\Lambda$CDM cosmology and that the reionization of neutral hydrogen in the intergalactic medium occurs below a redshift of $z=8$. We find that observations with the MWA-LFD cannot constrain the underlying cosmology in this scenario.  In principle, a similar experiment with a 10-fold increase in collecting area could provide useful constraints on the slope of the inflationary power spectrum, $n_{s}$, and the running of the spectral index, $\\alpha_{s}$, but these constraints are subject to the caveat that even a small reionization contribution could be confused with the cosmological signal.  In addition to the redshifted 21 cm signal, we include two nuisance components in our analysis related to the systematics and astrophysical foregrounds present in low-frequency radio observations. These components are found to be well separated from the signal and contribute little uncertainty ($<$30\\%) to the measured values of the cosmological model parameters. ",
        "introduction": "Hydrogen gas in the intergalactic medium (IGM) is a promising tool for studying some of the most interesting topics in astrophysics, including the nature of the first luminous objects, the processes of structure formation, and cosmology. Prior to the epoch of reionization (EOR), when radiation from the first luminous sources reionized the IGM, the 21 cm hyperfine line from neutral hydrogen is predicted to be visible against the cosmic microwave background (CMB) with relative brightness temperature of order $\\pm$10 mK. At high redshifts, the spin temperature of HI is expected to be less than the CMB temperature, making the line visible in absorption, whereas at lower redshifts, after radiation from the first luminous objects has begun to heat the IGM but before it has reionized it, the line should be evident as emission.  As the primordial hydrogen cools following recombination and later reheats, density contrasts are revealed as fluctuations in the brightness temperature of the redshifted 21 cm line \\citep{1972A&A....20..189S, 1979MNRAS.188..791H, 1990MNRAS.247..510S, 2002ApJ...572L.123I, 2003MNRAS.341...81I, 2004PhRvL..92u1301L, 2005ApJ...626....1B}. At high redshifts, the power spectrum of these fluctuations is expected to follow closely the dark matter power spectrum and measurements from this epoch could, independently from CMB and galaxy clustering experiments, constrain cosmological models. At lower redshifts, as the first luminous objects ionize their surroundings, voids are expected to appear in the diffuse emission \\citep{1997ApJ...475..429M, 2000ApJ...528..597T, 2003ApJ...596....1C, 2004ApJ...608..622Z, 2004ApJ...613...16F}. Measurements of the power spectrum during this epoch would help reveal the characteristics of the first luminous objects.  Recent efforts have distinguished theoretical aspects of redshifted 21 cm observations that are particularly relevant for constraining cosmological models. \\citet{2005MNRAS.363L..36B} discuss the ability of redshifted 21 cm observations to constrain the geometry of the high redshift universe between recombination and reionization; \\citet{2005MNRAS.363..251A} and \\citet{2005ApJ...624L..65B} have shown that differences in the line-of-sight versus angular components of a redshifted 21 cm power spectrum measurement can be used to separate primordial density perturbations from features caused by the radiative processes responsible for reionization; and \\citet{Barkana_AP} has discussed the application of the Alcock-Paczynski (AP) test to redshifted 21 cm measurements. Additionally, \\citet{2005ApJ...626....1B} consider the effects of the earliest galaxies on the redshifted 21 cm fluctuations and \\citet{2005MNRAS.362.1047N} discuss using redshifted 21 cm observations to study the thermal history of hydrogen gas by detecting a cutoff in the power spectrum due to thermal broadening of the line.  The experimental challenges of detecting and characterizing neutral hydrogen in the IGM at high redshift are significant.  While reaching the sensitivity limit needed to detect the redshifted 21 cm radio background is readily achievable by very small experiments, statistical measurements of the fluctuation power spectrum require the collecting area of the much larger arrays---MWA-LFD, LOFAR, and PAST---currently under development \\citep{2006ApJ...638...20B}. Even after the required collecting areas have been reached, anticipated challenges remain. Astrophysical foreground contaminants are five orders of magnitude brighter than the redshifted 21 cm emission. Removing the foreground signatures from observations requires extremely precise observational calibration.  In this paper, we expand on earlier calculations of the sensitivity of the first generation arrays to forecast their ability to place constraints on cosmological models from measurements of the redshifted 21 cm power spectrum.  In order to treat the best-case scenario and to simplify the analysis, we have chosen to assume that reionization has not yet begun by the target redshift range of $8 < z < 10$.  This assumption represents the most optimistic scenario generally consistent with the existing evidence about when reionization began from quasar absorption line measurements \\citep{2001ApJ...560L...5D, 2001AJ....122.2850B, 2003AJ....125.1649F, 2004Natur.427..815W} and from Thomson scattering measurements by the WMAP satellite \\citep{2006astro.ph..3449S}.  These measurements suggest reionization occurred at redshifts $z\\gtrsim6$ and $z\\sim10$, respectively.  We begin in Section 2 by reviewing the observational process of measuring the highly redshifted 21 cm power spectrum and describing the fiducial experiments.  In Section 3, the method to forecast constraints on cosmological parameters is described in terms of statistical errors using a Fisher matrix treatment of the full three-dimensional power spectrum. The calculation marginalizes over several cosmological parameters and two anticipated contributions related to the astrophysical foreground contaminants. The results are discussed in Section 4 and encapsulated in Figure \\ref{f2}.  A similar study has been performed in parallel by \\citet{McQuinn_Cosmology}.  While we analyze the characteristics relevant to constraining cosmological models with the initial generation of experiments and in the presence of two nuisance contributions, their analysis explores the benefit of combining redshifted 21 cm measurements with information from other experiments, such as WMAP and Planck. Together, the two efforts provide a thorough overview of the potential of future redshifted 21 cm power spectrum measurements to contribute to cosmology. ",
        "conclusions": "The first generation of redshifted 21 cm experiments have been shown previously to have the potential to characterize the processes and history of reionization at redshifts $6<z<12$. In this effort, we have considered a scenario in which reionization is not the dominant contribution to the measured power spectrum during this period.  Our findings demonstrate that the initial experiments would not contribute substantially to knowledge of the underlying cosmology under these circumstances.  While they may provide some use to cosmological studies, the primary science return of the first generation of redshifted 21 cm experiments will be regarding the astrophysics of reionization.  Under the optimistic assumptions used in this paper, the second-generation experiments, similar to the MWA5000 reference experiment, will be better suited to constrain $\\Omega_M$ and the primordial power spectrum through $n_s$ and $\\alpha_s$.  Even for the spectral index properties it is difficult to envision a scenario where redshifted 21 cm observations alone could provide an unambiguous constraint on cosmological information, however, since the signal could easily be confused if the universe had a small amount of reionization structure on small scales or spin temperature fluctuations in the IGM.  We have also considered the covariance between the redshifted 21 cm power spectrum due to a fully neutral IGM and the primary residual foregrounds of low-frequency radio observations. The signal and contamination components were easily distinguished in the full three-dimensional power spectrum analysis."
    },
    "0512/astro-ph0512112_arXiv.txt": {
        "abstract": "We propose a fast and efficient bispectrum statistic for Cosmic Microwave Background (CMB) temperature anisotropies to constrain the amplitude of the primordial non-Gaussian signal measured in terms of the non-linear coupling parameter $f_\\mathrm{NL}$. We show how the method can achieve a remarkable computational advantage by focussing on subsets of the multipole configurations, where the non-Gaussian signal is more concentrated. The detection power of the test, increases roughly linearly with the maximum multipole, as shown in the ideal case of an experiment without noise and gaps. The CPU-time scales as $\\ell_\\mathrm{max}^3$ instead of $\\ell_\\mathrm{max}^5$ for the full bispectrum which for Planck resolution $\\ell_\\mathrm{max}\\sim3000$ means an improvement in speed of a factor $10^7$ compared to the full bispectrum analysis with minor loss in precision. We find that the introduction of a galactic cut partially destroys the optimality of the configuration, which will then need to be dealt with in the future. We find for an ideal experiment with $\\ell_\\mathrm{max}=2000$ that upper limits of $f_\\mathrm{NL}<8$ can be obtained at $1\\sigma$. For the case of the WMAP experiment, we would be able to put limits of $|f_\\mathrm{NL}|<40$ if no galactic cut were present. Using the real data with galactic cut, we obtain an estimate of $-80<f_\\mathrm{NL}<80$ and $-160<f_\\mathrm{NL}<160$ at 1 and $2\\sigma$ respectively. ",
        "introduction": "In recent years a number of papers have focussed on the statistical nature of the fluctuations in the Cosmic Microwave Background radiation (CMB). In the standard inflationary scenario the quantum fluctuations of the inflaton scalar field follow a nearly Gaussian distribution, with small deviations arise by considering second-order terms of the equations \\cite{acqua,malda}: slow-roll conditions necessarily entail that deviations from Gaussianity is very low. Nonetheless, the subsequent gravitational evolution unavoidably enhances the primordial non-Gaussian signal up to the largest scales where the primordial seeds were produced, giving rise to a non-linearity parameter (see below) $f_{\\rm NL}\\sim {\\cal O}(1)$. The situation changes in the presence of a second scalar field during inflation; under these circumstances it has been shown by (Bartolo Matarrese and Riotto 2002), that non-Gaussianity can be transferred from the isocurvature to the adiabatic mode, leading to non-negligible values for $f_{\\rm NL}$ (see also \\cite{uzan,rigopoulos,seery}). Alternative scenarios for the production of the primordial seeds, such as the {\\it curvaton}(see \\cite{mollerach,enqvist,lyth,moroi,lyth2}) and the {\\it inhomogeneous reheating} mechanisms (see Dvali, Gruzinov and Zaldarriaga 2004) may also lead to higher values of $f_{\\rm NL}$. A general review for the primordial non-Gaussian scenarios can be found in (Bartolo et al. 2004). Testing for non-Gaussianity has then become the basic tool to discriminate among different models for the production of energy-density perturbations. It has become common practice to quantify the amount of non-Gaussianity by means of the dimensionless {\\it non-linearity parameter} $f_\\mathrm{NL}$ (see, {\\it e.g.} \\cite{ks}), setting the strength of quadratic non-linearities in an expansion of the large-scale gravitational potential $\\Phi$ (conventionally defined so that the temperature anisotropy is $\\Delta T/T \\equiv - \\frac{1}{3}\\Phi$ in the Sachs-Wolfe limit) in terms of a Gaussian random field $\\Phi_{\\rm G}$, namely \\begin{equation} \\Phi({\\bf x}) = \\Phi_{\\rm G}({\\bf x}) + f_{\\rm NL} \\Phi_{\\rm G}^2({\\bf x}) \\; \\end{equation} (up to a constant offset, which only affects the monopole contribution). Detailed calculations of the non-linearity parameter $f_{\\rm NL}$ during and after inflation (see \\cite{bmrPRL,bkmr04}) have shown that it unavoidably contains an angle-dependent part, whose role could be extremely important to look for specific signatures of inflationary non-Gaussianity as recently shown in \\cite{michele2}. In what follows, however, we will follow the common practice of taking $f_{\\rm NL}$ as a constant parameter. The upper limits on the estimated value of $f_\\mathrm{NL}$ have become more and more stringent as the sensitivity of CMB experiments has improved. With MAXIMA data, \\cite{santos} put a limit of  $|f_\\mathrm{NL}|< 950$, at 1 $\\sigma$ level. Using  COBE data, \\cite{komatsua} (with the bispectrum) and \\cite{cayona} (with Spherical Mexican Hat Wavelets (SMHW)) found $|f_\\mathrm{NL}|<1500$ and $|f_\\mathrm{NL}|< 1100$ respectively, at 1 $\\sigma$ level; these intervals were shrinked to $-58<f_\\mathrm{NL}<134$ at 2$\\sigma$ level in the first release of WMAP see \\cite{komatsu}. Using SMHW, \\cite{wang} found $f_\\mathrm{NL}=50\\pm80$ at 1 sigma and $ f_\\mathrm{NL}<220 $ at 2$ \\sigma$ level; in \\cite{cabella1} we constrained $f_\\mathrm{NL}=-5\\pm 175$ at 2$\\sigma$ level, combining the local curvature and spherical wavelets. Very recently the constraints on $ f_\\mathrm{NL}$ have been improved by \\cite{creminelli}, who find $-27<f_\\mathrm{NL}< 121$ at 2$\\sigma$ level. In \\cite{gaztanaga} very stringent limits are found but a direct comparison with the other methods is unfeasible because of discrepancies in the non-Gaussian models adopted.\\\\ \\noindent In this paper we focus on functionals of the normalized bispectrum; the latter has been considered by many authors in the literature, including \\cite{ks}; more recently, \\cite{Babich} has discussed conditions under which the bispectrum is the optimal estimator of primordial non-Gaussianity and \\cite{babzalda} showed that tighter constraints are expected from a joint analysis of temperature and polarization data.\\\\ \\noindent Here, we implement some procedures which were proposed in the statistical literature \\cite{marinucci}. In that paper, a full analytic derivation is provided for the test behaviour in the presence of an ideal experiment; the behaviour in the presence of non-Gaussianity is also discussed. Here, we investigate the properties of these tests under a realistic experimental setting, using both simulations and WMAP\\ data. The plan of this paper is as follows: in Section \\ref{sect:test} we describe the proposed procedure; Section \\ref{sect:impl} and \\ref{sect:data} describe the implementation, simulations and datasets used; Section \\ref{sect:MC} discusses the results of the procedure applied to the simulations; in Section \\ref{sect:concl} we draw some conclusions and discuss directions for future research. ",
        "conclusions": "\\label{sect:concl}  The bispectrum is one of the most common statistics to test for non-Gaussianity in CMB data. In particular, for estimating the non-linear coupling constant in primordial non-Gaussian models, the bispectrum has proven to produce very stringent limits. The drawback of the bispectrum is that the number of elements scale with the maximum multipole $L$ as $L^3$ and computational time scales as $L^5$.  In this paper we have introduced a new procedure to test for non-Gaussianity on CMB data; our approach is based on an integrated form of the bispectrum, where a phase factor is introduced and the focus is narrowed on nearly-collapsed configurations. Our approach is computationally very convenient as it scales only as $L^3$ and presents the added bonus to allow for explicit analytic results under idealized circumstances, in both Gaussian and non-Gaussian settings. The power properties are shown to be encouraging by Monte Carlo experiments: indeed some simple calculations suggest that the non-Gaussian signal grows linearly with the experiment resolution.  By comparing to limits on $f_\\mathrm{NL}$ published in the literature based on all elements of the bispectrum, we have shown that the collapsed configuration bispectrum does indeed seem to produce comparable near optimal results. For the ideal experiment we can constrain $f_\\mathrm{NL}<8$ using $L=2000$ and for a WMAP like experiment we find $f_\\mathrm{NL}<40$ when a galactic cut is not introduced. However, it turns out that a galactic cut does destroy the nice properties of the collapsed configuration. The multipoles get strongly coupled and the constrains on $f_\\mathrm{NL}$ becomes bigger that what one would expect from a pure increase in sampling variance. For this reason, when the galactic cut is introduced, we obtain $-80<f_\\mathrm{NL}<80$ at $1\\sigma$ applied to the WMAP data. Clearly, data from future CMB experiments like Planck will still need to be analysed with a galactic cut as it will be impossible to completely eliminate the galactic plane. For that reason, a way of dealing with the non-optimality of the collapsed configuration bispectrum will be necessary. One possible solution to this would be by some refilling procedure (or one could use the already existing Gibbs sampling technique \\cite{eriksen2}) which could restore orthogonality of the spherical harmonics. In this work, we did not make any prediction of the constrains on $f_\\mathrm{NL}$ which can be achieved by the Planck experiment. As we are still unable to deal optimally with cut sky, this would not produce a precise limit and is therefore postponed to future work.  Finally, we note that for an ideal experiment it seems feasible to achieve the bounds predicted in \\cite{ks}, despite the fact that we are using here only $L$, rather than $L^3$, bispectrum configurations. On the other hand, the presence of gaps greatly deteriorated the performance of our procedure. As the collapsed bispectrum seems potentially a very promising technique for high resolution data, we view the gap handling issue as a research priority in our future work."
    },
    "0512/astro-ph0512438_arXiv.txt": {
        "abstract": "We briefly discuss some open problems and recent developments in the investigation of the origin and propagation of ultra high energy cosmic rays (UHECRs). ",
        "introduction": "The detection of cosmic rays with energy in excess of $10^{20}$ eV puzzles the scientific community, mainly for two reasons: 1) most acceleration mechanisms need to be stretched to their extreme to allow for the generation of particles with this energy; 2) Energy losses of protons with energy high enough to suffer photopion production are such that, if the source distribution is homogeneous on scales larger than the loss length, a flux suppression should be expected in the spectrum, the so-called {\\it GZK feature} \\cite{gzk}. It is not clear yet if this feature  has been observed.  The importance of these two issues is here addressed in the context of the most recent findings in the investigation of particle acceleration at shock fronts, and in the light of the up-to-date measurements of the cosmic ray flux, anisotropy and chemical composition. A more extended discussion of these topics and their connection to the physics of lower energy cosmic rays can be found in two recent reviews \\cite{hillas,blasirev}.  This review is organized as follows: we discuss some developments in the theory of particle acceleration at shock waves in Sec. \\ref{sec:shocks}). In Sec. \\ref{sec:transition} we address the important issue of identifying the energy region in which the transition from a galactic to an extragalactic origin of cosmic rays takes place. In Sec. \\ref{sec:gzk} we discuss the current measurements of the spectrum and anisotropies of UHECRs. We conclude in Sec. \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We briefly discussed three issues related to the origin and propagation of ultra-high energy cosmic rays. In particular we summarized some recent developments in the investigation of particle acceleration at non-relativistic shock fronts, including the non linear dynamical reaction of the accelerated particles on the shock and the generation of waves by the accelerated particles. We also addressed the important issue of the transition from galactic to extragalactic origin of the observed cosmic rays. Finally we discussed the status of the measurements of spectrum and anisotropies of UHECRs."
    },
    "0512/astro-ph0512106_arXiv.txt": {
        "abstract": "Low-mass X-ray binaries (LMXBs), which occur in old stellar populations, have velocities exceeding those of their parent distribution by at least 20\\,km\\,s$^{-1}$. This makes them ideal probes for dark matter, in particular in dwarf spheroidals (dSph), where the LMXBs should penetrate well outside the visible galaxy. We argue that the most likely explanation of the observation of LMXBs in the Sculptor dSph by \\cite{MaccaroneEtal2005} is the presence of a dark matter halo of $\\ga 10^9\\,\\msun$, corresponding to a total-mass to light ratio of $\\ga600\\mlsun$. In this case there should be an extended halo of LMXBs which may be observable. ",
        "introduction": "A recent study by \\cite{MaccaroneEtal2005} reports the detection of five X--ray binaries with $L_X > 6\\times 10^{33}$~erg s$^{-1}$ in the Sculptor dwarf spheroidal (dSph) galaxy. Their membership of Sculptor is secure, as they are optically identified with counterparts in the catalogue of \\cite{SchweitzerEtal1995} which have appropriate proper motions and photometry. Given the old stellar population of Sculptor (see below) these must be low--mass X--ray binaries (LMXBs). In the Milky Way, LMXBs are observed to have relatively high space velocities $v_{\\rm sp} \\sim 20-100\\,$km\\,s$^{-1}$ (e.g. \\citealt{BrandtPodsiadlowski1995} and references therein; see \\citealt{PodsiadlowskiPfahlRappaport2005} for a recent review) resulting from the recoil of the system after the supernova creating the compact component. These velocities are far above the stellar velocity dispersion $\\sim11$\\,km\\,s$^{-1}$ in Sculptor, so it is interesting to ask how Sculptor can retain any LMXBs at all. ",
        "conclusions": "We have argued that the most likely explanation for the presence of LMXBs in Sculptor is a dark matter halo of $\\ga 10^9\\,\\msun$. This is $\\ga 100$ times the observed stellar mass and corresponds to a total-mass to light ratio of $\\ga600\\mlsun$. This value is comparable to that proposed for the Draco dSph as a result of simulations of tidal stripping by \\cite{ReadEtal2005, ReadEtal2006}. We argued further that in this case there should be an extended halo of LMXBs which might be observable. An estimate of the total number of LMXBs in this halo would constrain the transient duty cycle still further.  The obvious test of our ideas would come from measuring at least the radial velocities of the observed LMXBs. This should be possible given that \\cite{MaccaroneEtal2005} were able to make optical identifications of some of these sources, and would be extremely interesting whatever the result. If the deduced space velocities are high, this would confirm the presence of a very massive dark matter halo. If not, this would have important implications for understanding transient duty cycles and LMXB formation in general."
    },
    "0512/astro-ph0512040_arXiv.txt": {
        "abstract": "Higher order cumulants of point processes, such as skew and kurtosis, require significant computational effort to calculate. The traditional counts-in-cells method implicitly requires a large amount of computation since, for each sampling sphere, a count of particles is necessary. Although alternative methods based on tree algorithms can reduce execution time considerably, such methods still suffer from shot noise when measuring moments on low amplitude signals. We present a novel method for calculating higher order moments that is based upon first top-hat filtering the point process data on to a grid. After correcting for the smoothing process, we are able to sample this grid using an interpolation technique to calculate the statistics of interest. The filtering technique also suppresses noise and allows us to calculate skew and kurtosis when the point process is highly homogeneous. The algorithm can be implemented efficiently in a shared memory parallel environment provided a data-local random sampling technique is used. The local sampling technique allows us to obtain close to optimal speed-up for the sampling process on the Alphaserver GS320 NUMA architecture. ",
        "introduction": "\\label{sect:introduction} \\noindent Modern cosmology, the study of the large scale structure and evolution of our universe \\citep{pea}, has advanced to the point where we can now answer some very fundamental questions about the distribution of matter within our universe. Ever since Einstein postulated the theory of General Relativity and, together with De Sitter \\citep{pais}, showed how it could be applied to the universe as a whole, generations of physicists have pondered on the question of what is the overall geometry of our universe. Within the past few years observations of the relic microwave radiation from the ``Big Bang'' \\citep{wmap} have shown that the universe exhibits a geometry quite unlike that expected from theoretical prejudices alone.  Although on the largest scales the distribution of matter within our universe is both homogeneous and isotropic, on smaller scales---less than 1/20th the size of our visible universe---it is highly inhomogeneous. Even though the matter distribution of the universe was exceptionally smooth 300,000 years after the creation event \\citep{KT}, over billions of years the ubiquitous attraction of the gravitational force amplifies the minute fluctuations in the early matter distribution into the structure we see today. Moreover, the current best theories of structure formation suggest that the matter distribution we observe is formed in a `hierarchical clustering' manner with the small structures merging to form larger ones and so forth \\citep{pea}. This growth of structure is accelerated by an unseen massive `dark matter' component in our universe.  Although dark matter cannot be observed directly, there is sufficient evidence within observations to conclusively infer its existence. Modifications to Newton's equations, to change gravitational accelerations on large scales, have had limited success, and cannot presently be cast in a form compatible with General Relativity \\citep{stm}.  Understanding the distribution of matter within our local universe can tell us much about the cosmic structure formation process. While on the very largest scales gravity is the dominant force, on smaller scales gas pressure forces, from the gaseous inter-galactic (IGM) and inter-stellar mediums (ISM), can play a significant role.  In clusters of galaxies, for example, hydrodynamic forces produced by the IGM lead to a distribution of gas that is held close to hydrostatic equilibrium. Indeed, understanding the interaction between the ISM and the stars that condense out of it, is currently one of the hottest research areas in cosmology \\citep{rt01}. Since if we can understand this process we are much closer to being able to infer how the galaxies we observe relate to the underlying distribution of dark matter that dominates the evolution of structure.  Although we are yet to absolutely determine the relation between galaxies and dark matter, measuring the distribution of galaxies is the only way of infering the distribution of all matter (visible or not). Measurements of the speed of recession of local galaxies, led \\cite{ep} to form the distance-redshift relation now know as `Hubble's Law', which has become a bedrock for the development of cosmological theory. Although modern surveys of galaxies use an updated, and more accurate, form of the distance-redshift relation to uncover the spatial distribution of galaxies, the principles involved remain the same as those used by Hubble.  Aided by highly automated observing and computer driven data analysis, a new generation of high quality galaxy redshift surveys is mapping our local Universe with exquisite precision. The 2 degree field \\citep{2df}  and Sloan Digital Sky Survey \\citep{SDSS} provide astronomers with a survey of the local universe out to a redshift of $z\\simeq 0.3$, and contain over 200,000 and one million (when complete) redshifts respectively. In figure \\ref{2df} we show the distribution of galaxies for the 2dF survey to give an visual impression of the type of inhomogeneity observed.  \\begin{figure*}[!t] \\includegraphics[scale=0.875]{2dFcone2.eps} \\caption{Distribution of galaxies in the two main slices from the 2dF galaxy redshift survey. Each point represents a galaxy, and they combine to trace filament and wall structures in three dimensions. The geometry of the distribution is directly related to the statistical properties of the conditions in our universe following the ``Big Bang''.} \\label{2df} \\end{figure*}   Traditionally, one of the primary goals of analysis of redshift surveys is the calculation of the two point auto-correlation function (2-pt CF). The large sample volumes provided by 2dF and the SDSS have allowed the 2-pt CF to be calculated with great accuracy.  While the initial conditions produced by the ``Big Bang'' are widely believed to exhibit Gaussian statistics (\\eg Kolb and Turner, 1990), the formation of structure by gravitational instability introduces non-Gaussian features into the statistics of the matter distribution. Hence, the 2-pt CF cannot be a complete descriptor of the underlying matter distribution at late times. Astronomers were aware of this issue comparatively early in the development of the field, and the theoretical basis for calculating higher order statistics was developed through the 1970's (see \\cite{P80} for a detailed summary).  Early attempts to measure higher order moments of the mass distribution, via the counts-in-cells method (again see \\cite{P80}), suffered from inadequate sample size. Because higher order moments tend to be progressively dominated by the most dense regions in a given sample, ensuring that adequate sampling has been performed is of utmost importance. Ensuring low sample variance is also necessary, and given one sample the only way to check this is to analyse sub-samples, which rapidly depletes the available information.  From a theoretical perspective, higher order statistics are interesting in relation to gravitational perturbation theory and the evolution of non-linear gravitational clustering. Analyses examining the accuracy of numerical simulation methods often rely upon higher order statistics. This is especially important in the study of gravitational clustering in `scale free' universes \\citep{PC90}. The development of fast, parallel, statistical algorithms is vital to progress in this arena. While the development of parallel simulation algorithms has advanced forward rapidly (\\eg Thacker et al., 2003) development of parallel analysis tools has lagged behind. This is partially due to the fact that the benefits of developing a parallel analysis code can be shorted lived because the required analyses can change rapidly (much faster than the simulation algorithms themselves). The rapid development times available on shared memory parallel machines make them an ideal complement to large distributed memory machines which most simulations are now run on.  Although throughout this paper we discuss the application of our new method to cosmology, it can be applied equally well to the statistics of any point process. Indeed the terms `particle' and `point' are often used interchangeably. The method can also be modified to apply to different dimensions, although in 2 dimensions the gains are expected to be less significant due to the reduced amount of work in the counts-in-cells method.  The layout of this paper is as follows: in section \\ref{sect:stats}, we quickly review the statistics we wish to calculate. This is followed by an explicit description of our new algorithm, and an examination of its performance.  Next we present a brief case study on applying our algorithm to cosmology and conclude with a brief summary. ",
        "conclusions": "We have presented a new fast algorithm for rapid calculation of one point cumulants for point processes. Our algorithm is based upon a smoothed field approach, which reproduces the underlying statistical properties of the point processes field from which it is derived. The method is significantly faster \\linebreak than counts-in-cells methods because the overhead of evaluating the number of particles in a given sphere has been removed. We are able to calculate $10^9$ sample points on a $512^3$ data set in less than 2 CPU hours, which is over 4x faster than the results reported for tree-optimized counts-in-cells methods \\citep{S99}. We also note that while tree methods also lead to very large speed ups, they are still subject to noise from the point process for low amplitude signals.  We are currently applying this new technique to examine the evolution of high order moments in cosmological density fields at low amplitude levels and will present our findings elsewhere (Thacker, Couchman and Scoccimarro in prep). We also anticipate making the codes described in this paper publically available in the near future."
    },
    "0512/astro-ph0512089_arXiv.txt": {
        "abstract": "We present exact solutions of the massless Klein-Gordon equation in a spacetime in which an infinite straight cosmic string resides. The first solution represents a plane wave entering perpendicular to the string direction. We also present and analyze a solution with a static point-like source. In the short wavelength limit these solutions approach the results obtained by using the geometrical optics approximation: magnification occurs if the observer lies in front of the string within a strip of angular width $8\\pi G\\mu$, where $\\mu$ is the string tension. We find that when the distance from the observer to the string is less than $ 10^{-3} {(G \\mu)}^{-2}\\lambda \\sim 150 {\\rm Mpc} (\\lambda/{\\rm AU}) (G\\mu/10^{-8})^{-2}$, where $\\lambda$ is the wave length, the magnification is significantly reduced compared with the estimate based on the geometrical optics due to the diffraction effect. For gravitational waves from neutron star(NS)-NS mergers the several lensing events per year may be detected by DECIGO/BBO. ",
        "introduction": "Typical wavelength of gravitational waves from astrophysical compact objects such as BH(black hole)-BH binaries is in some cases very long so that wave optics must be used instead of geometrical optics when we discuss gravitational lensing. More precisely, if the wavelength becomes comparable or longer than the Schwarzschild radius of the lens object, the diffraction effect becomes important and as a result the magnification factor approaches unity \\cite{Ohanian,Bliokh,Bontz,Thorne,Deguchi}. Mainly due to the possibility that the wave effects could be observed by future gravitational wave observations, several authors \\cite{Takahashi:2003ix,Seto:2003iw,Nakamura:1997sw,Yamamoto:2003cd,Baraldo:1999ny,T.Nakamura:1999,Yamamoto:2003wg,Suyama:2005mx,Takahashi:2005sx,Takahashi:2004mc} have studied wave effects in gravitational lensing in recent years.  In most of the works which studied gravitational lensing phenomenon in the framework of wave optics, isolated and normal astronomical objects such as galaxies are concerned as lens objects. Recently Yamamoto and Tsunoda\\cite{Yamamoto:2003wg} studied wave effects in gravitational lensing by an infinite straight cosmic string. The metric around a cosmic string is completely different from that around a usual massive object.  Cosmic strings generically arise as solitons in a grand unified theory and could be produced in the early universe as a result of symmetry breaking phase transition\\cite{Hindmarsh:1994re,vilenkin}. If symmetry breaking occurred after inflation, the strings might survive until the present universe. Recently, cosmic strings attract a renewed interest partly because a variant of their formation mechanism was proposed in the context of the brane inflation scenario\\cite{Dvali:1998pa,Dvali:1999tq,Burgess:2001fx,Alexander:2001ks,Dvali:2001fw,Jones:2002cv,Shiu:2001sy}. In this scenario inflation is driven by the attractive force between parallel D-branes and parallel anti D-branes in a higher dimensional spacetime. When those brane-anti-brane pairs collide and annihilate at the end of inflation, lower-dimensional D-branes, which behave like monopoles, cosmic strings or domain walls from the view point of four-dimensional observers, are formed generically \\cite{Majumdar:2002hy,Dvali:2002fi,Jones:2003da,Dvali:2003zj,Copeland:2003bj}.  For some time, cosmic string was a candidate for the seed of structure formation of our universe, but this possibility was ruled out by the measurements of the spectrum of cosmic microwave background (CMB) anisotropies\\cite{Spergel:2003cb,Percival:2002gq}. The current upper bound on the dimensionless string tension $G\\mu$ is around $10^{-7} \\sim 10^{-6}$, which comes from the observations of CMB\\cite{Jeong:2004ut,Pogosian:2003mz,Pogosian:2004ny,Wyman:2005tu} and/or the pulsar timing \\cite{Kaspi:1994hp,Thorsett:1996dr,McHugh:1996hd,Lommen:2002je}. Although cosmic string cannot occupy dominant fraction of the energy density of the universe, its non-negligible population is still allowed observationally\\cite{Bouchet:2000hd,Rocher:2004my}. In fact, Sazhin et al.\\cite{Sazhin:2005fd,Sazhin:2003cp} reported that CSL-1, which is a double image of elliptical galaxies with angular separation $ 1.9~{\\rm arcsec} $, could be the first case of the gravitational lensing by a cosmic string with $ G\\mu \\approx 4\\times 10^{-7} $.  We study in detail wave effects in the gravitational lensing by an infinite straight cosmic string. In Ref.~\\cite{Yamamoto:2003wg}, wave propagation around a cosmic string was studied but they put the waveform around the string by hand. \\footnote{After submitting this paper, we have noticed a paper \\cite{Linet:1986db} in which the solutions of the wave equations around the cosmic string are given, though the apparent expressions are different from those given in this paper. In \\cite{Linet:1986db} the author estimated the amplitude of the diffracted wave to be suppressed by ${\\cal O}(G\\mu)$ compared with that corresponding to the geometrical optics. We show that the importance of the diffraction effects are determined by the combination of three parameters, $G\\mu$, the distance from the string to the observer and the wavelength and that the relative amplitude of the diffracted wave can be ${\\cal O}(1)$ for realistic astrophysical situations.} Their prescription is correct only in the limit of geometrical optics, which breaks down when the wavelength becomes longer than a certain characteristic length. In this paper, we present exact solutions of the (scalar) wave equation in a spacetime with a cosmic string. We analytically show that our solutions reduce to the results of the geometrical optics in the short wavelength limit. We derive a simple analytic formula of the leading order corrections to the geometrical optics due to the finite wavelength effects and also an expression for the long wavelength limit. Interference caused by the lensing remains due to the diffraction effects even when only a single image can be seen in the geometrical optics. This fact increases the lensing probability by cosmic strings.  This paper is organized as follows. In section II, we construct a solution of the wave equation on a background spacetime with an infinite straight cosmic string in the case that a source of the wave is located infinitely far. An extension to the case in which a point source is located at a finite distance is given in Appendix B. In section III, we study properties of the solution obtained in sec. II in detail. In section IV, we focus on compact binaries as the sources of gravitational waves and discuss the possible effects due to finiteness of the lifetime and the frequency evolution of the binaries on the detection of the gravitational waves which pass near a cosmic string. We also give a rough estimate for the event rate of the lensing of gravitational waves from NS-NS mergers assuming DECIGO/BBO. Section V is devoted to summary. ",
        "conclusions": "We have constructed a solution of the Klein-Gordon equation for a massless scalar field in the flat spacetime with a deficit angle $2\\pi\\Delta\\approx 8\\pi G\\mu$ caused by an infinite straight cosmic string. We showed analytically that the solution in the short wavelength limit reduces to the geometrical optics limit. We have also derived the correction to the amplification factor obtained in the geometrical optics approximation due to the finite wavelength effect and the expression in the long wavelength limit.  The waveform is characterized by a ratio of two different length scales. One length scale $r_s$ is defined as the separation between the image position on the lens plane in the geometrical optics and the string. We have two $r_s$ since there are two images corresponding to which side of the string the ray travels. (When the image cannot be seen directly, we assign a negative number to $r_s$.) The other length scale $r_F$, which is called Fresnel radius, is the geometrical mean of the wavelength and the typical separation among the source, the lens and the observer. The waveform is characterized by the ratios between $r_s$ and $r_F$. If $r_F>r_s$, the diffraction effect becomes important and the interference patterns are formed. Even when the image in the geometrical optics is not directly seen by the observer, the interference patterns remain. In contrast, in the geometrical optics magnification and interference occur only when the observer can see two images which travel both sides of the string. Namely, the angular range where lensing signals exist is broadened by the diffraction effect. This broadening may increase the lensing probability by an order of magnitude compared with that estimated by using the geometrical optics when the distance to the source is around the critical distance $D_c$ given in Eq.~(\\ref{es4}).   We finally estimated the rate of lensing events which can be detected by LISA and DECIGO/BBO assuming BH-BH or NS-NS mergers as a source of gravitational waves. For possible values of the parameters that determines the event rate such as string reconnection rate, string tension and the event rate of the unlensed mergers, the lensing event rate could reach several per yr."
    },
    "0512/hep-th0512257_arXiv.txt": {
        "abstract": "We present here a detailed study of the quasi-normal spectrum of brane-localised Standard Model fields in the vicinity of $D$-dimensional black-holes. A variety of such backgrounds (Schwarzschild, Reis\\-sner-Nordstr\\\"om and Schwarzszchild-(Anti) de Sitter) are investigated. The dependence of the quasi-normal spectra on the dimensionality $D$, spin of the field $s$, and multipole number $\\ell$  is analyzed. Analytical formulae are obtained for a number of limiting cases: in the limit of large multipole number for Schwarzschild, Schwarzschild-de Sitter and Reissner-Nordstr\\\"om black holes,  in the extremal limit of the Schwarzschild-de Sitter black hole, and in the limit of small horizon radius in the case of Schwarzschild-Anti de Sitter black holes. We show that an increase in the number of hidden, extra dimensions results in the faster damping of all fields living on the brane, and that the localization of fields on a brane affects the QN spectrum in a number of additional ways, both direct and indirect. ",
        "introduction": "Upon an external perturbation of a black hole background, realized either through the addition of a field or by perturbing the metric itself, the gravitational system enters a phase of damping oscillations, or quasi-normal ringing \\cite{Nollert1, KS} as it is alternatively called. During this phase, the frequency of the field consists of a real part $\\omega_{\\rm Re}$, that drives the field oscillations, and of an imaginary part $\\omega_{\\rm Im}$, that cau\\-ses the simultaneous damping of these oscil\\-lations. The smaller  $\\omega_{\\rm Im}$ is, the longer the damping time, therefore certain quasi-normal modes (QNMs) can dominate the spectrum at very late times after the initial perturbation, thus governing the dynamical evolution of the black hole.  The spectrum of quasi-normal modes has been the subject of an intensive study over the years, not only for their theoretical interest, but also due to the fact that their potential experimental detection could lead to the discovery of black holes. In a 4-dimensional context, they have been both numerically and analytically studied for a variety of black-hole backgrounds \\cite{quasi-4D}. Among the different species of fields, the quasi-normal modes associated to gravitons, generated by metric perturbations, were particularly looked at, for the simple reason that gravitational QNMs originating from astrophysical black holes could potentially be observed with the help of gravitational wave detectors \\cite{Nollert1, KS}. Unfortunately, such an experimental confirmation has not been obtained up to now.  A few years ago, the landscape in gravitational physics changed with the formulation of theories postulating the existence of additional spacelike dimensions in nature \\cite{RS, ADD}. According to these theories, all Standard Model (SM) particles (scalars, fermions and gauge bosons) are restricted to live on a (3+1)-dimensional hypersurface -- a brane -- embedded in the higher-dimensional `bulk'. Gravitons can progagate both on and off the brane, with the same being true for other particles, like scalars, that carry no charges under the Standard Model gauge group. This geometrical set-up protects the accurately observed properties of SM fields from being altered due to the presence of extra dimensions while opening the way for the study of new gravitational backgrounds.   The theory with Large Extra Dimensions \\cite{ADD} predicts the existence of $d$ additional spacelike compact dimensions, all having -- in the simplest case -- the same size $L$. Black holes produced by the gravitational collapse of matter on the brane would naturally extend off the brane, being gravitational objects. Whereas macroscopic astrophysical black holes extend mainly along the usual three, non-compact spatial dimensions, thus being effectively 4-dimensional objects, microscopic black holes with a horizon radius $r_H \\ll L$ would virtually live in a non-compact spacetime with $D=4+d$ dimensions in total. Higher-dimensional generalizations of 4-dimensional black hole solutions \\cite{Tangherlini, Myers}, derived previously, came back in the spotlight, and the study of the QNMs associated with these higher-dimensional backgrounds soon became the subject of a renewed research activity \\cite{quasi-D}.  One of the most exciting predictions of the theory with Large Extra Dimensions \\cite{ADD} is that such microscopic black holes may be created on ground-based accelerators during the collision of highly energetic particles with center-of-mass energy $\\sqrt{s}> M_*$ \\cite{creation}. The energy scale $M_*$ denotes the fundamental Planck scale of the higher-dimensional gravitational theory, that becomes effective at distances $r < L$. This scale can be much lower than the 4-dimensional Planck scale $M_{Pl}$, even as low as 1 TeV, therefore trans-planckian collisions can be easily realised at next-generation particle colliders. What is of the outmost importance is the fact that these tiny black holes, contrary to what happens in the case of macroscopic black holes, will be created in our neighbourhood in a controlled experiment inside a laboratory; therefore, their detection, either through the emission of Hawking radiation or the detection of their QNMs spectrum, will be substantially more favoured. In the latter case, the spectrum of QN modes associated with SM fields living in our brane will be the most important of all due to the well-developed techniques for the detection of fermions and gauge bosons, compared to the ones for the up-to-now elusive gravitons.  In this work, we attempt to fill the gap in the existing literature by presenting a comprehensive study of the QN modes associated with the brane-localized SM fields. We will examine a variety of higher-dimensional black hole backgrounds, namely the Schwarzschild, Reis\\-sner-Nordstr\\\"om and Schwarzschild-(Anti) de Sitter, all described by the same static, spherically-symmetric line-element with a single metric function. The spectrum of the QN modes for scalars, fermions and gauge bosons will be derived in each case, as a function of the dimensionality of spacetime, the multipole number and additional fundamental parameters such as the bulk cosmological constant and charge of the black hole. In what follows, we will be assuming that the black hole is characterized by a mass $M_{BH}$ that is at least a few orders of magnitude larger than the fundamental scale of gravity $M_*$, so that quantum corrections can be safely ignored. Also, in order to avoid a hierarchy problem, the brane self-energy can be naturally assumed to be of the order of the fundamental Planck scale $M_*$ and thus much smaller than $M_{BH}$, therefore its effect on the gravitational background can also be ignored.   In section II, we present the theoretical framework for our analysis and the equations of motion for the SM fields propagating in the brane background. Section III focuses on the associated QNM spectrum in the case of an induced-on-the-brane $D$-dimensional Schwarzschild black hole, and investigates the dependence of the spectrum on the dimensionality of spacetime, spin of the particle and multipole number. In Section IV we proceed to consider the QN modes of SM particles propagating on a brane embedded in a charged $D$-dimensional Reis\\-sner-Nordstr\\\"om black hole, and the effect of the black-hole charge on the QN spectrum is examined in detail. The spectra of QNMs for brane-localised SM fields in the background of a $D$-dimensional Schwarzschild-de Sitter and Schwarzschild-Anti de Sitter black hole are derived in sections V and VI, respectively, and the role of the bulk cosmological constant is investigated. We present our conclusions in section VII. ",
        "conclusions": "The theories postulating the existence of extra, compact spacelike dimensions have opened the way for low-scale gravitational theories and the observation of strong gravitational phenomena such as the creation of tiny, higher-dimensional black holes in the controlled environment of a ground-based particle accelerator. Standard Model fields, that are restricted to live on the 4-dimensional brane and happen to propagate in the vicinity of such a black hole, will feel only the projected-on-the brane black-hole background. In such a situation, the observation of their QN frequencies may be highly more likely to take place than the one for the diligently pursued, but up to now elusive, gravitons. Such a detection will provide not only evidence for the existence of the black holes themselves, but also information on the fundamental parameters of the higher-dimensional theory.   In this work, we have investigated a variety of sphe\\-ri\\-cally-symmetric $D$-dimensional black-hole backgrounds, that are then projected onto the 4-dimensional brane. The induced-on-the-brane backgrounds depend on a single metric function $h(r)$, that carries a signature of the parameters of the $D$-dimensional theory, such as the dimensionality $D$, charge $Q$ and cosmological constant $\\Lambda$. The same function characterizes the form of the effective potentials felt by the scalars, fermions and gauge bosons propagating in the brane background. In 3 of the cases studied (Schwarzschild, Reissner-Nordstr\\\"om and Schwarzschild-de Sitter), the effective potentials had the form of positive-definite barriers, a result that allowed us to use the WKB method to derive the QN spectrum in the 6th order beyond the eikonal approximation. In the 4th case (Schwarzschild-Anti de Sitter), the effective potential diverged at infinity, and the Horowitz-Hubeny method was used instead.   In the case of a Schwarzschild black-hole background, the QN spectrum of all SM fields was computed for various values of the dimensionality of spacetime $D$ and multipole number $\\ell$. It was found that, as the number of hidden, extra dimensions increased, the imaginary part of the QN frequency for all species increased as well, thus making the ring-down phase on the brane shorter. The real part of the QN frequency was also found to be $D$-dependent and to predominanlty increase with $D$, although particular modes with $\\ell \\simeq n$ may deviate from this behaviour. In respect to the dependence on the spin, fields with higher spin were found to damp with a slower rate and thus to survive longer.   If a charge $Q$ is present in the bulk background, the QN frequency spectrum is found to significantly deviate from the one of purely 4-dimensional ones, and to resemble more the one in the vicinity of a $D$-dimensional Gauss-Bonnet black hole: an increase in the charge of the black hole was found to lead to a monotonically decreasing behaviour for the imaginary part of the QN frequency -- and thus to a longer ring-down phase -- and to a monotonically increasing behaviour for the real part, rendering all SM fields much better oscillators than in the neutral case.   When a positive cosmological constant is turned on in the bulk, the resulting QN spectrum of the SM fields on the brane does not yield any surprises. As $\\Lambda$ increases, both the real and imaginary part of the QN frequency are suppressed, and for $\\Lambda=\\Lambda_{ext}$, the suppression reaches the magnitude of 90\\%. The number of dimensions enhances again the individual values of $\\omega_{\\rm Re}$ and $\\omega_{\\rm Im}$, however the relative suppression as $\\Lambda$ varies comes out to be only mildly $D$-dependent. In the extreme limit, the feature of fields with different spin decaying with almost the same rate is observed also in the present case of brane-localised fields in an effective SdS background.   In the case of a negative cosmological constant being present in the bulk, the spectrum of QN frequencies was shown to depend on the ratio of the black-hole horizon to the AdS radius, similarly to the case of purely 4-dimensional or $D$-dimensional SAdS backgrounds.  For large black holes, the QN spectrum comes out to be proportional to $r_H$, and to become equidistant for the higher overtones. In the opposite limit of a very small black hole, the QN frequencies approach the normal modes of the projected-on-the-brane AdS spacetime.   Summarizing our results, we may say that the existence of extra dimensions affects the QNM spectrum of fields living on the brane both in a direct and an indirect way. In all the cases studied here, an increase in the number of transverse-to-the-brane dimensions causes a significant increase in the imaginary part of all fields, thus directly affecting the damping rate of the field perturbations on the brane. Additional features, such as the distance in the frequency spacing between successive quasi-normal modes in the SAdS spacetime, or asymptotic values of QN frequencies in various limits and backgrounds, also depend explicitly on the total number of dimensions $D$. In an indirect way, the localization of fields on a brane embedded in a higher-dimensional black-hole background leads to a deviation from the behaviour observed either in the purely 4-dimensional or in a purely $D$-dimensional case: the monotonic behaviour of the QN spectrum as a function of the charge $Q$, for the majority of the modes, is an indicative example of this.   In this work, we have restricted our analysis to the study of spherically-symmetric $D$-dimensional black-hole backgrounds projected on the brane. The study of axi\\-ally-symmetric, rotating black-hole backgrounds has already been initiated and we hope to report our results soon in a follow-up article. In addition, here we were limited by the consideration of quasi-normal spectra for only massless fields, yet, as was shown in \\cite{massivescalar} for scalars, and in \\cite{massivevector} for vector fields, the mass term can change the lower modes of the spectrum considerably, leaving unaffected the high damping limit of the spectrum -- the study of the effect of the mass term on the QN spectrum of brane-localized fields in various backgrounds is also among our future plans. \\vspace{3mm}"
    },
    "0512/astro-ph0512510_arXiv.txt": {
        "abstract": "We present a detailed chemical abundance analysis of the parent star of the transiting extrasolar planet TrES-1. Based on high-resolution Keck/HIRES and HET/HRS spectra, we have determined abundances relative to the Sun for 16 elements (Na, Mg, Al, Si, Ca, Sc, Ti, V, Cr, Mn, Co, Ni, Cu, Zn, Y, and Ba). The resulting average abundance of $<[$$X$/H$]>$ $= -0.02\\pm0.06$ is in good agreement with initial estimates of solar metallicity based on iron. We compare the elemental abundances of TrES-1 with those of the sample of stars with planets, searching for possible chemical abundance anomalies. TrES-1 appears not to be chemically peculiar in any measurable way.  We investigate possible signs of selective accretion of refractory elements in TrES-1 and other stars with planets, and find no statistically significant trends of metallicity [$X$/H] with condensation temperature $T_c$. We use published abundances and kinematic information for the sample of planet-hosting stars (including TrES-1) and several statistical indicators to provide an updated classification in terms of their likelihood to belong to either the thin disk or the thick disk of the Milky Way Galaxy. TrES-1 is found to be a very likely member of the thin disk population. By comparing $\\alpha$-element abundances of planet hosts and a large control sample of field stars, we also find that metal-rich ([Fe/H]$\\gtrsim 0.0$) stars with planets appear to be systematically underabundant in [$\\alpha$/Fe] by $\\approx 0.1$ dex with respect to comparison field stars. The reason for this signature is unclear, but systematic differences in the analysis procedures adopted by different groups cannot be ruled out. ",
        "introduction": "The possibility that super-solar metallicity could imply a higher likelihood of a given star to harbor a planet was investigated since the first detections by precision radial-velocity surveys (\\citeauthor{gonzalez97}~\\citeyear{gonzalez97}, ~\\citeyear{gonzalez98a},~\\citeyear{gonzalez98b}; \\citeauthor{fuhrmann97}~\\citeyear{fuhrmann97},~\\citeyear{fuhrmann98}; \\citeauthor{laughlin97}~\\citeyear{laughlin97}). A number of studies have been performed throughout these years, with increasingly larger sample sizes, employing both spectroscopic and photometric techniques for metallicity determination (using iron as the primary reference element), and adopting control samples of field stars without detected planets (\\citeauthor{santos00}~\\citeyear{santos00},~\\citeyear{santos01}, ~\\citeyear{santos03},~\\citeyear{santos04a},~\\citeyear{santos05}; \\citeauthor{reid02}~\\citeyear{reid02}; \\citeauthor{laughlin00}~\\citeyear{laughlin00}; \\citeauthor{gonzalez00}~\\citeyear{gonzalez00}; \\citeauthor{gonzalez01}~\\citeyear{gonzalez01}; \\citeauthor{israelian01}~\\citeyear{israelian01}; \\citeauthor{queloz00a}~\\citeyear{queloz00a}; \\citeauthor{smith01}~\\citeyear{smith01}; \\citeauthor{gimenez00}~\\citeyear{gimenez00}; \\citeauthor{martell02}~\\citeyear{martell02}; \\citeauthor{heiter03}~\\citeyear{heiter03}; \\citeauthor{sadakane02}~\\citeyear{sadakane02}; \\citeauthor{pinsonneault01}~\\citeyear{pinsonneault01}; \\citeauthor{murray02}~\\citeyear{murray02}; \\citeauthor{laws03}~\\citeyear{laws03}; \\citeauthor{fischer05}~\\citeyear{fischer05}).  The global trend is that planet-harboring stars are indeed more metal rich than stars without known planets. Based on observationally unbiased stellar samples, the strong dependence of planetary frequency on the host star metallicity has been clearly demonstrated by e.g., Santos et al. (\\citeyear{santos01}, \\citeyear{santos04a}), and Fischer \\& Valenti (\\citeyear{fischer05}). Furthermore, the metallicity enhancement is likely to be ``primordial'' in nature, i.e. due to the intrinsically high metal content of the protoplanetary cloud from which the planetary systems formed, as opposed to the possibility of ``self-enrichment'', caused by accretion of rocky planetesimal material onto the parent star (see Gonzalez \\citeyear{gonzalez03} for a review of the subject). This conclusion is primarily based upon the evidence of no dependence of the iron-abundance enhancement on the stellar effective temperature, as theoretical calculations would predict (e.g., ~\\citealp{dotter03}; \\citealp{cody05}, and references therein, but see also \\citealp{vauclair04} for somewhat different arguments), and it bears important consequences for the proposed models of giant planet formation by core accretion (e.g., \\citealp{ida05}; \\citealp{kornet05}) and disk instability (\\citealp{boss02}).  Based on detailed chemical abundance analyses of metals other than iron, several attempts have been made in the recent past to confirm the observed trend and to put on firmer grounds (or refute) the idea that stars with planets are primordially metal-rich, and have not been polluted. Many authors have determined the abundances of over a dozen other elements for planet hosts, including light elements such as Li and the isotopic ratio $^6$Li/$^7$Li (\\citeauthor{gonzalez00}~\\citeyear{gonzalez00}; \\citeauthor{ryan00}~\\citeyear{ryan00}; \\citeauthor{israelian01}~\\citeyear{israelian01},~\\citeyear{israelian03}, ~\\citeyear{israelian04}; \\citeauthor{reddy02}~\\citeyear{reddy02}; \\citeauthor{mandell04}~\\citeyear{mandell04}) and Be (\\citeauthor{garcia98}~\\citeyear{garcia98}; \\citeauthor{deliyannis00}~\\citeyear{deliyannis00}; \\citeauthor{santos02}~\\citeyear{santos02},~\\citeyear{santos04b}), refractories such as the $\\alpha$-elements Si, Mg, Ca, Ti, and the iron-group elements Cr, Ni, and Co, and volatiles such as C, N, O, S, and Zn (\\citeauthor{santos00}~\\citeyear{santos00}; \\citeauthor{gonzalez01}~\\citeyear{gonzalez01}; \\citeauthor{smith01}~\\citeyear{smith01}; \\citeauthor{takeda01}~\\citeyear{takeda01}; \\citeauthor{sadakane02}~\\citeyear{sadakane02}; \\citeauthor{zhao02}~\\citeyear{zhao02}; \\citeauthor{bodaghee03}~\\citeyear{bodaghee03}; \\citeauthor{ecuvillon04a}~\\citeyear{ecuvillon04a},~\\citeyear{ecuvillon04b}, ~\\citeyear{ecuvillon05a}; \\citeauthor{beirao05}~\\citeyear{beirao05}; \\citeauthor{gilli05}~\\citeyear{gilli05}).  For instance, detection of anomalous light-element abundances in the atmosphere of a star could be indicative of recent planetary accretion events. While evidence for Li excesses in some planet-harboring stars has been reported in the literature (\\citeauthor{israelian01}~\\citeyear{israelian01},~\\citeyear{israelian03}; \\citeauthor{laws01}~\\citeyear{laws01}), clearly suggesting that accretion of planetary material can actually take place in some stars, as implied by theoretical arguments (\\citeauthor{montalban02}~\\citeyear{montalban02}; \\citeauthor{boesgaard02}~\\citeyear{boesgaard02}; \\citealp{sandquist02}), in general stars with planets have normal light-element abundances, typical of field stars (e.g., \\citeauthor{ryan00}~\\citeyear{ryan00}; \\citeauthor{israelian04}~\\citeyear{israelian04}).  Arguments in favor of the ``self-enrichment'' hypothesis could also be substantiated if volatile elements were to exhibit different abundance trends with respect to refractory elements. One way to approach the problem is to make use of the condensation temperatures $T_c$ of the elements, a typical diagnostic employed for investigating chemical fractionation patterns in many areas of planetary science and astronomy (e.g., \\citeauthor{lodders03}~\\citeyear{lodders03}, and references therein). In this particular case, volatiles, having low $T_c$-values, are expected to show a deficiency in accreted material with respect to refractories. However, the most recent evidence (e.g., \\citeauthor{bodaghee03}~\\citeyear{bodaghee03}; \\citeauthor{ecuvillon04a}~\\citeyear{ecuvillon04a},~\\citeyear{ecuvillon04b}, ~\\citeyear{ecuvillon05a}; \\citeauthor{gilli05}~\\citeyear{gilli05}) is that the abundance distributions of other elements in stars with planets are simply the extension of the observed behavior for [Fe/H], a result quantified by trends of decreasing [$X$/Fe] with increasing [Fe/H], for both refractories and volatiles. It thus seems unlikely that pollution effects can be responsible for the overall metallicity enhancement of the planet host stellar sample.  The primary goal of this work is to present a detailed study of the chemical composition of the parent star of the recently discovered transiting extrasolar planet TrES-1 (GSC 02652-01324; \\citealp{alonso04}). We have done so by undertaking a detailed chemical abundance analysis using our Keck and Hobby Eberly Telescope (HET) spectra of TrES-1. Secondly, we have compared the elemental abundances of TrES-1 with those of the sample of stars with planets, in order to search for possible chemical abundance anomalies in the former. To this end, we have utilized results from uniform studies of elemental abundances of large sets of planet hosts available in the literature. Third, in an attempt to find circumstantial evidence of possible selective accretion of planetary material, we have further investigated the sample of planet hosts and TrES-1, searching for statistically significant trends of [$X$/H] with condensation temperature. Finally, we have utilized the chemical composition information for TrES-1 and a large sample of planet-hosting stars along with their kinematic properties in order to classify them, based on a number of diagnostic indicators, in terms of their likelihood of being members of the thin or thick disk populations of the Milky Way Galaxy (e.g., \\citealp{gilmore83}; \\citealp{carney89}. See Majewski \\citeyear{majewski93}, and references therein, for a comprehensive review and discussion of formation scenarios). This analysis has the purpose of revisiting and updating the results of a few past studies (\\citealp{gonzalez99}; \\citeauthor{reid02}~\\citeyear{reid02}; \\citeauthor{barbgratt02} \\citeyear{barbgratt02}; \\citeauthor{santos03} \\citeyear{santos03}) which, using limited sample sizes, confirmed the strong similarity between the kinematic properties of stars with planets and that of control samples of stars without known planets.  This paper is organized as follows. In Section 2 we present our chemical abundance analysis for the planet-hosting star TrES-1. All elemental abundances are compared in Section 3 with those of selected, uniformly studied samples of planet hosts. Section 4 is dedicated to an updated classification of planet-harboring stars in terms of different stellar populations in our Galaxy. Finally, Section 5 contains a summary of the main results and concluding remarks. ",
        "conclusions": "In this work we have carried out an abundance analysis of 16 elements for the parent star of the transiting extrasolar planet TrES-1. The resulting average abundance of $<[X$/H$]>= -0.02\\pm0.06$ is in good agreement with initial estimates of solar metallicity based on iron (\\citeauthor{sozzetti04} \\citeyear{sozzetti04}; \\citeauthor{laughlin05} \\citeyear{laughlin05}). TrES-1 appears not to be chemically peculiar in any measurable way when compared to a large sample of known stars with planets. No convincing evidence for statistically significant trends of metallicity [$X$/H] with condensation temperature $T_c$ can be found for TrES-1 or other planet hosts, a further indication that selective accretion of planetary material is not likely to be responsible for the observed high metal content of stars with detected planets, a conclusion similar to those drawn by others (e.g., \\citeauthor{santos01} \\citeyear{santos01},~\\citeyear{santos04a}; \\citeauthor{fischer05} \\citeyear{fischer05}). Using its abundance and kinematic information, we have classified TrES-1 as a likely member of the thin disk population, and provided updated membership probabilities for a large set of planet hosts, based on the relative agreement between different statistical indicators (purely kinematic in nature, solely based on chemistry, and a combination of the two). Finally, we have highlighted an apparent systematic underabundance in [$\\alpha$/Fe] of stars with planets compared to a large comparison sample of field stars. The more likely cause for this signature resides in unknown systematics in the details of the abundance analysis procedures adopted by different authors. However, we have found hints for differences between the [$\\alpha$/Fe] abundances of planet hosts and control samples analyzed in exactly the same way. In this respect, we stress the importance of continuously updating and expanding uniform, detailed studies of the chemical composition of planet hosts (including both refractory and volatile elements spanning a wide range of condensation temperatures) as new objects are added to the list, as well as statistically significant control samples of stars without any detected planets, following those recently undertaken by e.g., Bodaghee et al. (\\citeyear{bodaghee03}), Santos et al. (\\citeyear{santos04a}, \\citeyear{santos04b}, \\citeyear{santos05}), Israelian et al. (\\citeyear{israelian04}), Ecuvillon et al. (\\citeyear{ecuvillon04a}, \\citeyear{ecuvillon04b}, \\citeyear{ecuvillon05a}, \\citeyear{ecuvillon05b},  \\citeyear{ecuvillon05c}), Beir\\~ao et al. (\\citeyear{beirao05}), Gilli et al. (\\citeyear{gilli05}), Fischer \\& Valenti (\\citeyear{fischer05}), Valenti \\& Fischer (\\citeyear{valenti05}), and Gonzalez (\\citeyear{gonzalez05}). Such investigations would also help to disentangle possible signatures induced by the presence of planets from effects related instead to Galactic chemical evolution."
    },
    "0512/astro-ph0512383_arXiv.txt": {
        "abstract": "We present new optical observations of young massive star clusters in Arp 220, the nearest ultraluminous infrared galaxy, taken in $UBVI$ with the Hubble Space Telescope ACS/HRC camera. We find a total of 206 probable clusters whose spatial distribution is centrally concentrated toward the nucleus of Arp 220. We use model star cluster tracks to determine ages, luminosities, and masses for 14 clusters with complete $UBVI$ indices or previously published near-infrared data. We estimate rough masses for 24 additional clusters with $I < 24$ mag from $BVI$ indices alone. The clusters with useful ages fall into two distinct groups: a ``young'' population ($< 10$ Myr) and an intermediate-age population ($\\simeq 300$ Myr). There are many clusters with masses clearly above $10^6 M_{\\odot}$ and possibly even above $10^7 M_{\\odot}$ in the most extreme instances. These masses are high enough that the clusters being formed in the Arp 220 starburst can be considered as genuine young globular clusters. In addition, this study allows us to extend the observed correlation between global star formation rate and maximum cluster luminosity by more than an order of magnitude in star formation rate. ",
        "introduction": "Young massive star clusters \\citep[or YMCs; see][]{larsen02} are an intriguing mode of star formation in the present-day universe. While their older and usually more massive counterparts, the classic globular clusters, are found around almost every type of galaxy \\citep{h01}, rich populations of luminous blue star clusters are found predominantly in starburst and interacting systems \\citep[][among many others]{h92,w93,ws95,bastian05}. However, the fact that individual young clusters have been found in several nearby dwarf galaxies \\citep[e.g.][]{c94,o94,bhe02} and small populations are found in several nearby spiral galaxies \\citep{m96,l00,larsen01,larsen02} suggests that massive star cluster formation is a relatively wide-spread phenomenon, although it seems to occur with high efficiency only in the most active star-forming systems.  Many questions remain about the properties and ultimate fate of these young massive clusters. The combination of high and variable reddening and uncertain ages has often made it difficult to determine accurate masses for them. Dynamical masses are the most reliable, but these are only available for a few systems \\citep{ho96,mengel02}. Recently, intermediate-age clusters with dynamical masses greater than 10$^7$ M$_\\odot$ have been identified in two merger-remnant galaxies, \\object{NGC 7252} and \\object{NGC 1316} \\citep{maraston04,bastian06}. Among merger and merger-remnant galaxies, only in the nearest system, the Antennae (\\object{NGC 4038}\\object[NGC 4039]{/39}), have accurate photometric masses and ages been determined for large numbers of clusters \\citep{w99,zf99,wz02}. In the Antennae, both the youngest clusters ($< 6$ Myr) and a slightly older population of clusters (25-160 Myr) reach masses as large as $3-4\\times 10^6$ M$_\\odot$ \\citep{zf99,wz02}. A single YMC in \\object{NGC 6946}, a much more modest starburst system, has a mass near $10^6 M_{\\odot}$ as well \\citep{larsen01}. In comparison, the most massive globular clusters range from $5\\times 10^6$ M$_\\odot$ for $\\omega$ Cen in the Milky Way \\citep{meylan95} to above $\\sim 10^7$ M$_\\odot$ for the most extreme known cases such as the cluster G1 in \\object{M31} \\citep{meylan01}, the most massive clusters in \\object{NGC 5128} \\citep{mar04}, and the most luminous globular clusters in supergiant elliptical galaxies \\citep{h05}. Thus, an intriguing question is whether we can identify very young clusters ($< 10$ Myr) as massive as $10^7$ M$_\\odot$ in galaxies in the local universe. Since stars and star clusters form in molecular gas, the best place to search for the most massive young star clusters is in the most gas-rich galaxies.   \\object{Arp 220} is the closest example of an ultra-luminous infrared galaxy \\citep{so87}. At a distance of 77 Mpc, it is only four times more distant than the Antennae system and only slightly more distant than the merger remnant NGC 7252, and represents our best chance to identify and study very young massive star clusters in an ultra-luminous infrared galaxy. Arp 220 has faint tidal tails and distortions seen in both optical and HI emission in the outer parts of the galaxy \\citep{a66,j85,h00} and twin nuclei separated by only 300 pc \\citep{s98,s99,so99}. By comparison with the models of \\citet{mh96} and assuming the two progenitor galaxies to be similar to the Milky Way, \\citet{m01} estimate the time since the beginning of the interaction to be $\\sim 700$ Myr, with the recent burst of star formation that powers the galactic superwind and bubbles likely to have started 10-100 Myr ago. Arp 220 contains a molecular gas mass of $9\\times 10^9$ M$_\\odot$ \\citep{s97}.  Although this is roughly half the total mass of molecular gas in the Antennae \\citep{g01}, its molecular gas is concentrated to the inner 750 pc radius of the galaxy, so that its average molecular gas surface density reaches an astounding $5\\times 10^4$ M$_\\odot$ pc$^{-2}$ \\citep{s97}. This surface density corresponds to $A_v = 3300$ mag for a standard gas-to-dust ratio and is comparable to the average surface density in a dense star-forming core inside a giant molecular cloud \\citep{motte01}. In short, if any nearby galaxy has the fuel and the conditions to be forming extremely massive young star clusters, it should be Arp 220.  The first observations of Arp220 with the Hubble Space Telescope identified eight compact objects, of which the two brightest were suggested to be massive associations of young stars \\citep{s94}. Near-infrared observations by \\citet{s98} identified eight bright star cluster candidates in Arp 220. \\citet{s01} combined these two sets of data to estimate ages for three of these star clusters in the range of 10-100 Myr. In this paper, we present new $UBVI$ observations obtained with the Hubble Space Telescope (HST) to search for additional young massive star clusters in Arp 220 and determine more detailed properties for them. The observations and data reduction are presented in \\S~\\ref{s2}, and we estimate masses, ages, and reddenings for the cluster candidates by comparison to the \\citet{bc03} models in \\S~\\ref{s3}. We discuss the implications of our results for the formation of young massive clusters in \\S~\\ref{s4}. ",
        "conclusions": "We have used new $UBVI$ optical imaging with the ACS/HRC camera on the Hubble Space Telescope to identify 206 star cluster candidates in the ultraluminous infrared galaxy Arp 220. These cluster candidates show a radial gradient in their surface density with distance from the center of Arp 220, which suggests that most of them are star clusters associated with the galaxy. One of the star clusters is spatially resolved and may have a half-light diameter of roughly 20 pc, which would be twice the size of the massive Galactic globular cluster $\\omega$ Cen.  Due to high and variable reddening, only seven clusters are detected in our deep $U$ image. We have been able to derive accurate masses and ages for these seven clusters, as well as for seven additional clusters with previously published 1.6 $\\mu$m data from the NICMOS camera. These clusters divide into two distinct age groups: young clusters with ages $< 10$ Myr, and intermediate age clusters with ages of 70 to 500 Myr. Most of the younger clusters are more massive than $10^6$ M$_\\odot$, with the most massive being perhaps as much as $10^7$ M$_\\odot$ depending on its precise age. The intermediate mass clusters are somewhat less massive on average, ranging from $2\\times 10^5$ to $2\\times 10^6$ M$_\\odot$. Rough mass estimates for 24 clusters with $I < 24$ mag suggest most of these clusters have masses in the range $10^5 - 10^6$ M$_\\odot$.  The identification of a very young, massive star cluster in Arp 220 allows us to extend the correlation between the global star formation rate and the most luminous cluster seen by \\citet{bhe02} by an order of magnitude. This result implies that very high star formation rates are required to form clusters more massive than $10^7$ M$_\\odot$, which suggests that the merger remnants NGC 7252 and NGC 1316 should have experienced peak star formation rates greater than 100 M$_\\odot$ yr$^{-1}$ at some point in the merging process."
    },
    "0512/astro-ph0512456_arXiv.txt": {
        "abstract": "We present an analysis of the rotation measures (RMs) of polarized extragalactic point sources in the Southern Galactic Plane Survey. This work demonstrates that the statistics of fluctuations in RM differ for the spiral arms and the interarm regions. Structure functions of RM are flat in the spiral arms, while they increase in the interarms. This indicates that there are no correlated RM fluctuations in the magneto-ionized interstellar medium in the spiral arms on scales larger than $ \\sim 0^{\\circ}.5$, corresponding to $\\sim 17$~pc in the nearest spiral arm probed. The non-zero slopes in interarm regions imply a much larger scale of RM fluctuations. We conclude that fluctuations in the magneto-ionic medium in the Milky Way spiral arms are not dominated by the mainly supernova-driven turbulent cascade in the global ISM but are probably due to a different source, most likely H~{\\sc ii} regions. ",
        "introduction": "Structure in the neutral and ionized interstellar gas of the Milky Way is ubiquitous and present on many scales. There have been many observational, theoretical, and computational papers concerning turbulence in the neutral gas and molecular clouds, but relatively little is known about the structure of the warm ionized gas component (see Elmegreen \\& Scalo 2004, Scalo \\& Elmegreen 2004 for an overview). Turbulence in the ionized gas is suggested by observations of non-thermal linewidths in H$\\alpha$ \\citep{r85, trh99}, and modeled in numerical simulations of the multiphase ISM (see V\\'azquez-Semadeni et al. (2003) and references therein). Furthermore, observed power spectra or structure functions of electron density show power law behavior indicative of incompressible hydrodynamical turbulence (Cordes et al.\\ 1985, Armstrong et al.\\ 1995). However, care should be taken in interpretation of these density spectra, as the connection between density and velocity structure is not unambiguous, and fluctuations in electron density can be created by a number of different processes such as small-amplitude plasma waves or differences in ionization fraction.  The observed density fluctuations in the ionized ISM in the Galactic plane differ from those in the halo. The plane shows enhanced scattering of extragalactic radio sources (Spangler \\& Reynolds 1990), higher rotation measures (RMs) from extragalactic sources (e.g.\\ Clegg et al.\\ 1992), and increased scintillation of pulsars and angular broadening of extragalactic sources (e.g.\\ Cordes et al.\\ 1985). Although fluctuations on scales of hundreds of parsecs have been observed out of the  Galactic plane (e.g.\\ Armstrong et al.\\ 1995), the gas in the plane may be partly dominated by structures on much smaller scales (Haverkorn et al.\\ 2004). Internal structure in individual H~{\\sc ii} regions in the Galactic plane may be responsible for these enhanced fluctuations in the ionized ISM (Spangler \\& Reynolds 1990, Haverkorn et al.\\ 2004).  These previous studies have considered how the properties of fluctuations in the ISM vary with Galactic latitude. In this {\\em Letter}, we examine structure in the thermal electron density and magnetic field in the warm ionized ISM as a function of longitude, to detect any change in characteristics between spiral arms and interarm regions. ",
        "conclusions": "We have shown that in the Galactic spiral arms no correlated fluctuations exist in the magnetized interstellar plasma on scales larger than $\\sim0^{\\circ}.5$, corresponding to $\\sim 17$~pc in the nearest spiral arm probed, i.e.\\ the Carina arm which starts at $\\sim2$~kpc distance in this direction \\citep{gg76, r03}. In the interarm regions, however, correlated magnetoionic fluctuations {\\it are} present on large scales. Since the RM is a line of sight average through the entire Galaxy, it is not possible to associate a spatial scale to this angular scale. But assuming that the largest angular scales represent nearby structure at a fairly arbitrary distance of $\\sim1$~kpc away, an outer scale of $4-5^{\\circ}$ in the interarm regions would correspond to a spatial scale of about 100~pc.  The measured slopes of RM SFs in the interarm regions roughly agree with slopes of {\\it velocity} structure functions in simulations of incompressible (magneto-) hydrodynamical turbulence \\citep{k41, mg01, clv02}. This suggests that the density and velocity spectra may be coupled, which is only the case for subsonic or mildly supersonic turbulence \\citep{kr05, blc05}. Therefore, if the RM fluctuations are connected to velocity structure, the turbulence in the interarm ionized gas must be subsonic or only mildly supersonic, in agreement with observational results from H$\\alpha$.  On the other hand, the RM structure observed in spiral arms is probably {\\it not} a part of the turbulent cascade in the diffuse ionized medium. If the density spectrum traces the velocity spectrum, this would mean that the outer scale of turbulence would be $\\la 17$~pc. However, the dominant source of turbulent driving is believed to be supernova remnants and superbubbles, injecting energy on much larger scales \\citep{m03, db05}.  A shallow RM SF can also be caused by highly supersonic compressible turbulence, which will flatten the density SF \\citep{kr05, blc05}. However, as mentioned before, it is doubtful whether the ionized gas in the spiral arms is very supersonic.  Therefore, the fluctuations in RM in the spiral arms are probably not connected to the turbulent cascade in the diffuse ionized ISM at all. Instead, a probable source of these fluctuations is \\ion{H}{2} regions, which are of the correct size, sufficiently abundant \\citep{hgm04} and concentrated in spiral arms. Widespread interstellar turbulence injected by \\ion{H}{2} regions is not expected to be significant \\citep{m03}, but the structure could be caused by the \\ion{H}{2} regions themselves, or by turbulence inside them."
    },
    "0512/physics0512231_arXiv.txt": {
        "abstract": "The kinetic theory of collisionless electrostatic shocks resulting from the collision of plasma slabs with different temperatures and densities is presented. The theoretical results are confirmed by self-consistent particle-in-cell simulations, revealing the formation and stable propagation of electrostatic shocks with very high Mach numbers ($M \\gg 10$), well above the predictions of the classical theories for electrostatic shocks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512366_arXiv.txt": {
        "abstract": "{We study the red giant populations of two dE galaxies, AM~1339-445 and AM~1343-452, with the aim of investigating the number and luminosity of any upper asymptotic giant branch (AGB) stars present. The galaxies are members of the Centaurus~A group (D $\\approx$ 3.8 Mpc) and are classified as outlying (R $\\approx$ 350 kpc) satellites of Cen~A\\@. The analysis is based on near-IR photometry for individual red giant stars, derived from images obtained with ISAAC on the VLT. The photometry, along with optical data derived from WFPC2 images retrieved from the HST science archive, enable us to investigate the stellar populations of the dEs in the vicinity of the red giant branch (RGB) tip.  In both systems we find stars above the RGB tip, which we interpret as intermediate-age upper-AGB stars. The presence of such stars is indicative of extended star formation in these dEs similar to that seen in many, but not all, dEs in the Local Group. For AM~1339-445, the brightest of the upper-AGB stars have M$_{bol}$ $\\approx$ --4.5 while those in AM~1343-452 have M$_{bol}$ $\\approx$ --4.8 mag.  These luminosities suggest ages of approximately 6.5 $\\pm$ 1 and 4 $\\pm$ 1 Gyr as estimates for the epoch of the last episode of significant star formation in these systems. In both cases the number of upper-AGB stars suggests that $\\sim$15\\% of the total stellar population is in the form of intermediate-age stars, considerably less than is the case for outlying dE satellites of the Milky Way such as Fornax and Leo~I. ",
        "introduction": "\\label{sect:Intro}  Dwarf elliptical and dwarf spheroidal galaxies (we will refer to both classes as dwarf elliptical (dE) galaxies from now on) are often assumed to have simple star formation histories because, at the present epoch, they generally lack neutral hydrogen and show no current or very recent star formation.  These properties separate them, morphologically, from the gas-rich, star-forming dwarf irregular (dIrr) galaxies.  However, detailed studies of resolved stellar populations in Local Group (LG) dEs have revealed a surprising diversity of star formation histories, with some LG dEs containing stars as young as $\\sim$1 Gyr, or perhaps even younger \\citep[e.g.\\ Fornax and Leo~I; see reviews by][]{dacosta97,mateo98,vandenbergh99,vandenbergh00,grebel00}.  The dwarf galaxies in the LG follow a morphology-density relation in that the dEs are preferentially close to the Milky Way (MW) or M31, the two most massive galaxies in the LG, while almost all of the dIrr galaxies are located in the outskirts of the LG in more isolated regions.  Further, among the MW and M31 dE satellite galaxies, there is a tendency for the systems that lie at larger distances to have larger intermediate-age (i.e.\\ age $\\approx$ 1--10 Gyr) populations \\citep[e.g.][]{vandenbergh94}.  This has led to the conjecture that proximity to a luminous galaxy, and the type of that galaxy, can play an important role in defining dwarf galaxy evolution.  In particular, externally driven processes such as ram-pressure stripping by the hot gaseous halo of a massive galaxy, and/or tidal effects, may control the rate at which gas is lost from a dwarf system, and thus its star formation history. This approach has support from numerical simulations \\citep[e.g.][]{mayer+01b}. \\citet{grebel+03} also argue that external gas removal mechanisms are required to generate low-luminosity dEs, though they suggest that the transition-type dwarfs, i.e.\\ the dwarfs classified as dIrr/dE, are the best model for low luminosity dEs that have not lost their gas, rather than dIrrs.  However, the model in which the evolution of a dwarf system is halted by externally induced gas loss, does not easily explain the existence of LG dEs like Tucana and Cetus.  These dwarfs show very little, if any, evidence for extended star formation despite having isolated locations far from the MW or M31 \\citep[e.g.][]{dacosta98,sarajedini+02}.  If these dEs have always remained isolated, then an internal process must have been responsible for the apparent complete gas loss at early times.  In order to explore further the importance of external vs.\\ internal processes in driving the star formation in these small systems, it is necessary to study in detail dE galaxies in environments different from that of the LG\\@. The relatively nearby Centaurus~A Group (D $\\approx$ 3.8 Mpc) is one such environment.  This group, which has the unusual giant E~galaxy Cen~A as its single dominant member, is more compact than the LG and it probably contains perhaps twice as many galaxies.  For instance, \\citet{k+02} list 13 galaxies that are most likely within a 600 kpc radius of Cen~A, and which are brighter than M$_{B}$ $\\approx$ --12. Besides Cen~A itself, these include the giant galaxies NGC~4945 and NGC~5102 as well as 10 dwarf galaxies, 6 of which are early-type and 4 late-type.  This catalogue is likely to be significantly incomplete. In contrast, in the Local Group, for an approximately equivalent volume centered on the LG barycentre and with an equivalent magnitude cutoff, the compilation of \\citet{vandenbergh00} yields a (complete) sample of 13 galaxies: 4 large galaxies (MW, M31, M33 and the LMC), 6 early-type dwarfs (M32, NGC~205, NGC~185, NGC~147, Fornax and Sagittarius) and 3 late-type dwarfs (SMC, IC10 and IC1613).  Thus it is likely that the Cen~A group has provided a significantly different environment for its dwarf members than has the LG\\@.  We have begun a program to study the red giant populations of the dE galaxies in the Cen~A group, with the ultimate aim of investigating the extent to which star formation history indicators correlate with distance from the dominant galaxy of the group, Cen~A\\@.  Specifically, we will investigate the number and luminosity of upper-AGB stars in the dE galaxies. Upper-AGB stars are stars with sufficient mass to evolve to luminosities above the RGB tip, and their presence, provided the system is relatively metal-poor ($\\langle$[Fe/H]$\\rangle$ $\\leq$ --1.0), as is the case for all but the most luminous dEs, is an unambiguous indicator of the existence of an intermediate-age ($\\sim$1--10 Gyr) population.  The luminosity of the brightest upper-AGB stars is also a measure of the age of the youngest intermediate population of significance.  Because of their cool effective temperatures, these upper-AGB stars are best studied in the near-infrared. Indeed it was near-IR observations that provided the first evidence for the diversity of star formation histories among the MW dSph companions \\citep[cf.][]{aaronson+mould80}.  Using the ISAAC near-IR array at the ESO Very Large Telescope (VLT) we have obtained $J_s$ and $K_s$-band images of 14 Cen~A group dwarf galaxies. The full data set, together with description of the reduction and analysis techniques, will be presented in a forthcoming paper (Rejkuba et al., in preparation). We present here the first results of our program -- analysis of the resolved red giant stellar populations of two Cen~A group dE galaxies, AM~1339-445 (KK 211) and AM~1343-452 (KK 217). For these two galaxies it is possible to supplement the near-IR data with $V$ and $I$ band WFPC2 images from the HST science archive, permitting estimates of the distances and of the average metallicity of their stars.  In Table~\\ref{tab:target_character} we summarize the fundamental parameters of the two dwarf ellipticals. The last column gives the literature reference for the tabulated data.  Using the integrated magnitudes of \\citet{JBF00} and the reddenings and moduli adopted here (see Table~\\ref{tab:target_character} and following sections) the absolute blue magnitudes of these two galaxies are M$_{B}=-11.9$ and $-10.8$, respectively, with colours $(B-R)_0=1.38$ and $1.35$, typical for dE galaxies \\citep[e.g.][]{evans+90}. Neither galaxy is detected in the HIPASS survey: the 3-$\\sigma$ upper limits on their H{\\small I} contents correspond to $2.8 \\times 10^6 M_\\odot$  and $3.1 \\times 10^6 M_\\odot$, respectively, for an assumed 10 kms$^{-1}$ line width and detection limits of the survey from \\citet{barnes+01}. \\citet{k+02} list line-of-sight distances for the two galaxies based on the $I$ magnitude of the red giant branch tip (see also Sect.\\ \\ref{sect:results_optical}). With those distances, the angular separations from Cen~A, and a distance for Cen~A of 3.84 Mpc \\citep{rejkuba04}, the true separation of AM~1339-445 from Cen~A is $\\sim$390 kpc, while that of AM~1343-452 is approximately 320 kpc. \\citet{k+02} label both dwarfs as companions of Cen~A\\@. The distances of both these dEs from Cen~A are therefore somewhat larger than those of the outer dE satellites Leo~I and Leo~II from the MW, and those of the outer dE satellites And~II, And~VI (Peg) and And~VII (Cas) from M31.  \\begin{table} \\caption{Fundamental parameters of the two target dE galaxies in Cen A group.} \\label{tab:target_character} \\centering \\begin{tabular}{lccl} \\hline \\hline Name               & AM~1339-445  & AM~1343-452 &\\\\ \\hline $\\alpha_{J2000.0}$ & 13:42:05.8   &  13:46:18.8 &\\\\ $\\delta_{J2000.0}$ & $-$45:12:21  & $-$45:41:05 &\\\\ Type               & dE           & dE          & 1\\\\ $B_T$ (mag)        & $16.32$        & $17.57$       & 1\\\\ $R_T$ (mag)        & $14.76$        & $16.02$       & 1\\\\ $(B-R)_0^T$ (mag)  & $1.38$           & $1.35$   & 1 \\\\ $r_{\\mathrm{eff},R}$ (arcsec)& $23.8$   & $14.7$    & 1\\\\ $\\langle \\mu \\rangle_{\\mathrm{eff},R}$& $23.63$  & $23.85$ & 1\\\\ E($B-V$) (mag)     & $0.111$        & $0.121$    & 2 \\\\ $(m-M)_0$ (mag)   & $27.87 \\pm 0.27$ & $27.99 \\pm 0.37$ & 1 (SBF)\\\\ & $27.77 \\pm 0.21$ & $27.92 \\pm 0.25$ & 3 (TRGB)\\\\ & $27.74 \\pm 0.20$ & $27.86 \\pm 0.20$ & 4 (TRGB)\\\\ $\\langle \\mathrm{[Fe/H]} \\rangle $ &$-1.4 \\pm 0.2 $ & $-1.6 \\pm 0.2$& 4\\\\ \\hline References: &\\multicolumn{3}{l}{1: \\citet{JFB00}; 2: \\citet{schlegel+98};}\\\\ &\\multicolumn{3}{l}{3: \\citet{k+02}; 4: this work} \\end{tabular} \\end{table}   The paper is organised as follows.  In the following section the observations and reductions are described, first for the near-IR data and then for the WFPC2 data.  The third section outlines the analysis of the resulting colour-magnitude diagrams (CMDs) for the two dwarfs.  The fourth section discusses the results, which are summarised in the final section. A preliminary description of this work has appeared in \\citet{dacosta05} and \\citet{rejkuba+05}. ",
        "conclusions": "We have presented an analysis of the red giant populations of two dE galaxies in the Cen~A group, AM~1339-445 and AM~1343-452, using a combination of near-IR and optical data.  Both dEs are distant companions of Cen~A, the dominant galaxy of the group.  Using the luminosity of the tip of the RGB we have measured distance moduli of both galaxies, $(\\mathrm{m}-\\mathrm{M})_0(\\mathrm{AM1339-445})=27.74 \\pm 0.20$, and $(\\mathrm{m}-\\mathrm{M})_0(\\mathrm{AM1343-452})=27.86 \\pm 0.20$, which are in good agreement with previously published values \\citep{k+02,JFB00}. The mean colour of the upper RGB stars is used to determine the mean metallicities of $\\langle \\mathrm{[Fe/H]}\\rangle =-1.4 \\pm 0.2$ for AM~1339-445 and $\\langle \\mathrm{[Fe/H]}\\rangle =-1.6 \\pm 0.2$ for AM~1343-452. The integrated colours of these two dEs are similar to those of other dEs and they follow the same luminosity-metallicity relation of the LG dEs.  We find evidence for the presence of intermediate-age upper-AGB stars in both galaxies, with the most luminous of these stars being 0.2-0.3 mag brighter in AM~1343-452, than in AM~1339-445.  The luminosities of these stars indicate that significant star formation continued in AM~1343-452 until an age of $\\sim$4 Gyr as against $\\sim$6.5 Gyr in AM~1339-445.  In this respect these Cen~A group dEs are similar to the outlying dE satellites of the Milky Way.  However, we estimate that the fraction of the total population that is of intermediate-age is perhaps $\\sim$15\\%, which is significantly less than the dominant intermediate-age populations found in outlying Milky Way dE satellites such as Fornax and Leo~I.  With only two galaxies it is premature to draw any definite conclusions regarding our long term goal of investigating the role of environment on the evolution of dE galaxies in the Cen~A group (cf.~Sect.~\\ref{sect:Intro}). Nevertheless, it is interesting that despite the rather large distance of both dEs from Cen~A, their intermediate-age populations are small and relatively old, particularly when compared to the outer dE satellites of the Milky Way.  We must, however, await similar analyses for additional dEs in this group before drawing any inferences from this result."
    },
    "0512/astro-ph0512150_arXiv.txt": {
        "abstract": "We consider the general problem of the tidal capture or circularisation from large eccentricity of a uniformly rotating object. We extend the self-adjoint formalism introduced in our recent paper(Papaloizou \\& Ivanov 2005 hereafter PI))  to derive  general expressions for the energy and angular momentum transfered when  the  planet or a star passes through  periastron in a parabolic or highly eccentric orbit around a central mass. These can be used  without making a low frequency approximation as was done in PI. We show how these can be adapted  to the low frequency limit in which only inertial modes contribute for baratropic planet models. In order to make quantitative estimates, we  calculate the inertial mode eigenspectrum  for planet models of  one and five  Jupiter masses $M_J,$  without a solid core, with different radii corresponding to different ages. The spectra are found in general to be more complex than of a polytrope with index  $n=1.5,$ considered in PI, because of the existence  global modes associated with the transition from molecular to metallic Hydrogen. Nonetheless the main tidal response is still found to be determined  by two global modes which have polytropic counterparts.  These also determine the uniform angular velocity in a state of pseudo synchronisation, for which the angular momentum transfered during an encounter is zero. This is found to be close to $1.55$  times the circular orbit angular velocity at periastron for all models considered. This is in contrast to the situation when only the $f$ mode is considered (Ivanov \\& Papaloizou 2004, hereafter IP) and the equilibrium angular velocity is found to be much larger.  We consider the multi-passage problem when there is no dissipation finding that stochastic instability  resulting in the  stochastic gain of inertial mode energy over many periastron passages occurs under similar conditions to those already found by IP for the $f$ modes. We find that this requires  circularisation to  start with a  semi-major axis exceeding $\\sim 30 AU,$ for final periods of $\\sim 3$ days. reducing to $\\sim 1-2 AU$  for final periods $\\sim 1.2$ days.  Finally we  apply our  calculations of the energy transfer during a periastron passage to  the problem of the  tidal circularisation of the orbits of the extrasolar planets in a state of pseudo synchronisation, expected because of the relatively small inertia of the planet,  and  find that inertial mode excitation dominates the tidal interaction for $1 M_J$ planets that start with semi- major axes less than $10 AU$  and   end up on circular orbits with final period in the $4-6$ day range.  It is potentially able to account for initial circularisation up to a final $6$ day  period within a few $Gyr$  But in the case of $5M_J$ oscillation modes excited in the star are more important. ",
        "introduction": "The process of  tidal capture, whereby a body in  initially unbound orbit has a close approach to another gravitating mass, and through the excitation of internal modes of oscillation, becomes subsequently bound is of general importance in astrophysics. It is believed to play a role in binary formation in globular clusters (eg. Press \\& Teukolsky 1977, hereafter PT,  and references therein) and to play a role in the tidal interactions of stars in galactic centres. A  related problem is the circularisation of an orbit from large eccentricity which may occur after a tidal capture. In this  situation, each periastron passage is very similar to a  close encounter in a weakly unbound orbit  that causes a change in the energy contained within internal modes of oscillation at the expense of that in the orbit.  In order to estimate capture probabilities and circularisation time scales the excitation of normal modes of oscillation by a perturbing potential must be considered. In general modes associated with spherical stars, or small perturbations of them,  have been considered (eg. PT, IP and references therein). However, if the tidally perturbed body has multiple encounters, angular momentum transfer should lead it to rotate at an angular velocity related to that at periastron. In this case  modes of oscillation with periods comparable to the rotation period need to be considered (eg. Papaloizou \\& Pringle 1978).  An approach for dealing with the problem of calculating tidal phenomena while taking into account the part of the oscillation spectrum with eigen frequencies comparable to the rotation frequency was recently proposed by PI. They adopted a self-adjoint formalism for calculating the tidal response that used the density perturbation to describe the motions within the star and they made a low frequency approximation that the rotation and mode frequencies were very much less than the inverse sound crossing time. They also evaluated the inertial mode spectrum for a polytrope with index $n=3/2$ and showed that inertial modes dominated the response at low frequencies or large periastron distances.  Given the importance of rotation, the purpose of this paper is to develop further the work begun in PI and to use the results to  obtain time scales for tidal phenomena such as circularisation of the orbit starting from high eccentricity.  We begin by noting that the self-adjoint formalism can be extended to general linear displacements without making a low frequency approximation. In this form, which uses the Lagrangian displacement to describe the perturbations it has wider applicability than the formalism presented in PI which made a low frequency approximation. We also wish to evaluate the expressions obtained for the energy and angular momentum exchanged during close periastron passages for realistic planet models. Accordingly we calculate sequences of  models  of one Jupiter mass $M_J$ and $5M_J.$ Each sequence  has a range of radii in order to take account of the fact that the planet radius varies with age. For simplicity we do not include a solid core noting that according to Saumon \\& Guillot (2004) observations of Jupiter are not inconsistent with this. We calculate the oscillation  spectra together with required overlap integrals for the models using a set of basis or trial functions.  We then use the determined eigenmode properties for these models to estimate tidal phenomena such as the attainment of pseudo synchronism for which the angular momentum transfer at periastron passage is zero. Such a state is expected to be attained, on account of the small inertia of the planet compared to that of the orbit, with subsequent energy exchanges occurring with the planet rotating at an equilibrium angular velocity. The inclusion of inertial modes is important for determining this angular velocity which turns out to be  close to that of the orbit at periastron. Calculations in IP incorporating only the  $f$ modes gave a much higher value.  Having calculated the equilibrium angular velocity we are able to consider multiple encounters and determine conditions for the onset of stochastic instability, that  results in the  stochastic gain of inertial mode energy over many periastron passages when there is no dissipation (Kochanek 1992, Mardling 1994a,b). Then we make estimates of circularisation time scales for the two different planet masses taking into account the evolutionary variation of the planet radius with age.  The plan of the paper is as follows. In section 2 we formulate the new self-adjoint formalism for the problem of linear perturbations proposed in our recent paper(PI) using the Lagrangian displacement to describe the motions and then use it to derive  general expressions for the energy and angular momentum transfered when a rotating planet or a star passes through periastron in a parabolic or highly eccentric orbit around a central mass without making a low frequency approximation as in PI. It can thus be applied to general situations in which $p,$ $f,$ $g$ and inertial modes are significant. In section 3  we show that the general formalism can be reduced to that of PI in an appropriate low frequency limit.  In this case only the inertial modes contribute to the tidal response for the baratropic stellar models we consider. In section 4 we discuss the tidal energy and angular momentum transfered through inertial modes giving very simple expressions that apply when the planet or star rotates in a state of pseudo synchronisation for which the angular momentum transfered in an encounter is essentially zero. We apply these to  the multi-passage problem when there is no dissipation finding conditions for the occurrence of the  stochastic instability. We find that this requires the circularisation process to start with a rather large semi-major axis $\\sim 30 AU,$ for final periods of $\\sim 3$ days. However, this could be reduced to $\\sim 1-2 AU$ for the shortest observed final periods $\\sim 1.2$ days.  In section 5  we  calculate the eigenspectrum of the inertial modes for the planet models which have been obtained using a realistic equation of state.  We find that the counterparts of the two global modes found in a polytrope with $n=1.5$ exist in these cases but that there are additional global modes associated with the transition from molecular to metallic Hydrogen. We go on to apply the results to  the problem of tidal circularisation of the orbits of the extrasolar planets in sections 6 and 7 finding that inertial modes dominate the tidal interaction for $M_J$ planets ending up on circular orbits with final period in the $4-6$ day range and are potentially  able to account for the initial circularisation at a $6$ day  period within a few $Gyr$ in that case. However, in the case of $5M_J$ oscillation modes excited in the star become more important. Finally in section 8 we discuss our results. ",
        "conclusions": "In this Paper we have developed and extended a new self-adjoint formalism for the problem of small stellar oscillations proposed in our recent paper(PI). We went on to apply this to the calculation of the tidal response of uniformly rotating fully convective planets and stars which undergo a single encounter in a parabolic orbit or a sequence of multiple encounters in a weakly bound orbit. This was then applied to the problem of the tidal circularisation of the orbits of the extrasolar planets.  In section 2 we  began by  showing that the general formalism, presented in PI when $W=\\rho'c^2/\\rho+\\Psi^{int}$ is used as variable characterising linear eigenmodes, can also be applied when  the Lagrangian displacement is used. The advantage of this approach is that, as no low frequency approximation is needed, it allows us to consider pulsations of  arbitrary frequency. We derived  general expressions for the energy and angular momentum transfer that occurs when a rotating planet or a star passes through  periastron in a parabolic or highly eccentric orbit around a central mass The amplitudes and phases of the eigenmodes excited in  the planet or a star as a result of the tidal encounter were also determined.   We showed that these expressions can be represented in terms of a spectral decomposition over the normal modes of the self-adjoint operator $(\\ref {eq p15})$. We also obtained a self-consistent expression for the rate of dissipation of  an excited mode to leading order in the small parameter defined by the ratio of pulsation period to a  characteristic  dissipation time scale.  In section 3  we  focused on the case of low frequency inertial modes, showing how the more general formalism in terms of the Lagrangian displacement can be reduced to that in PI in the low frequency limit. This procedure allows one to separate out the contribution associated with the inertial modes and simplify the equations governing  free pulsation of a rotating planet or a star. In that case the corresponding self-adjoint problem is formulated through equations $(\\ref {eq p47})$ and $(\\ref {eq p48})$. Note that the tidal response of the inertial modes to  impulsive tidal forcing has not been described using other methods.  We showed  in section 4 that when the rotating planet or a star approaches periastron in an unperturbed state the tidal energy and angular momentum transfered through the excitation of inertial modes are given by very simple expressions $(\\ref {eq p65})$ and $(\\ref {eq p68})$. We considered the multi-passage problem assuming that there was no significant dissipation of  mode energy between successive periastron passages, and found a simple condition $(\\ref {eq p78})$ for the  occurrence of  stochastic instability that  results in the  stochastic gain of inertial mode energy over many periastron passages. This was quantitatively similar to that given by IP for $f$ modes. We found that  for  stochastic instability  the circularisation process has to start with $a >  \\sim 30 AU,$ for final periods of $\\sim 3$ days. or $\\sim 1-2 AU$ for final orbital periods $\\sim 1.2$ days.   In section 5  we applied our formalism to  fully convective rotating giant planets. In order to evaluate the expressions for the energy and angular momentum transfer the eigenspectrum of  $(\\ref {eq p48})$ must be found numerically, for a given model. We calculated the eigenspectrum of the inertial modes for several planet models with a realistic equation of state. The details and stability of the numerical method we used are discussed in section 5.2. We considered planet models with two different masses equal to $1M_{J}$ and $5M_{J}$, with  radii in the range $1-2R_{J}$. We found that, as in the case of $n=1.5$ polytrope considered by PI,  the tidal response was determined by a few 'global' eigenmodes with a large scale distribution of  perturbed quantities. Two 'main' or 'standard' global modes have eigenfrequencies close to those corresponding to the $n = 1.5$ polytrope. These modes have the  largest overlap integrals characterising coupling between the tidal field and the eigenmodes. However, we found that for realistic planet models, there could be  'non-standard' global modes possibly related to a sharp change of structure occurring  near to the point of ionisation of hydrogen. The structure of the 'non-standard' modes as well as their stability with respect to a change of the planet mass and radius was  discussed.  The results obtained for the eigen spectra of realistic planet models were applied to the problem of tidal circularisation of the orbits of the extrasolar planets in sections 6 and 7. In section 6 we calculated the rotational angular velocity for which the net angular momentum transferred was zero. Because of the relative low inertia of the planet compared to the orbit it is expected to rapidly attain this angular velocity and achieve a state of pseudo synchronisation. The angular velocity associated with pseudo synchronisation was found to be always close to $1.55$ times the angular velocity  of a circular orbit at periastron.  Orbital evolution arises  through  the transfer of  orbital energy to the modes of oscillation. Transfers at successive periastron passages lead to orbital circularisation. This is the case, either when the mode energy is dissipated directly between encounters, or when the system is in the regime of  stochastic instability. We compared the contribution   associated with the inertial modes with the contribution associated with the fundamental modes, in a state of pseudo synchronisation, to the circularisation time scale in section 7. We found  that the inertial modes led to effective circularisation of a giant planet with mass $\\sim 1M_{J}$ on a time scale smaller than or of the order of a few $Gyr$, when  the initial semi major axis was less than $10 AU$ and the final orbital period after circularisation $P_{orb} < 6days.$ In this case the inertial modes  play the most important role at the longer periods. However, the inertial waves are less important for the planets of  larger mass $5M_{J}$ that we considered. In that case the dynamic tides exerted on the star central star, which we assumed to be solar like,  play a major role.  In our opinion, the most important unresolved question related to the theory discussed above is concerned with internal  dissipation. Of particular concern is the internal location of modal dissipation and the related issue of whether the energy deposited can be  radiated away between successive periastron passages. There are several potentially interesting channels of dissipation which may operate in our case. The non-linear parametric instability discussed by Kumar and Goodman (1996) for the case of $g$ modes in convectively stable stars may be effective for the inertial modes as well. This instability operates when there are eigenmodes in the stars with eigenfrequencies approximately half  that of the eigenfrequencies of the tidally excited modes and sufficiently large coupling coefficients. These modes may be unstable on account of parametric resonance. Taking into account that the eigenspectrum of the inertial modes is dense, unstable  modes are expected to  exist for sufficiently weak dissipative processes. This question will be addressed in our future work.  Our formalism is general enough that it  can be extended to the case of convectively stable uniformly rotating stars and it can also be applied to other astrophysical situations where excitation of stellar pulsations with frequencies comparable to the angular velocity of rotation is important. In this particular work we assumed that the orbital and planetary angular momentum vectors were aligned. This is a very reasonable assumption when the system has many periastron passages and so readily attains a state of pseudo synchronisation during the circularisation process. However, misaligned angular momenta should be considered if the main issue is tidal capture, as in globular clusters, because in that case there is no reason for the initial angular momentum vectors to be correlated. These issues and extensions of the formalism will  be considered in future work.                               \\vfill \\eject      \\vspace{-0.7cm}"
    },
    "0512/astro-ph0512199_arXiv.txt": {
        "abstract": "{We study the H$\\alpha$ emission from jets using two-dimensional axisymmetrical simulations. We compare the emission obtained from hydrodynamic (HD) simulations with that obtained from magnetohydrodynamics (MHD) simulations. The magnetic field is supposed to be present in the jet only, and with a toroidal configuration. The simulations have time-dependent ejection velocities and different intensities for the initial magnetic field. The results show an increase in the H$\\alpha$ emission along the jet for the magnetized cases with respect to the HD case. The increase in the emission is due to a better collimation of the jet in the MHD case, and to a small increase in the shock velocity. These results could have important implications for the interpretation of the observations of jets from young stellar objects. ",
        "introduction": "Collimated outflows are observed in a variety of astrophysical objects, with typical spatial scales ranging from $\\sim 1$~pc for jets from young stellar objects (YSOs) up to several megaparsecs for extragalactic jets. All of these jets seem to be associated with accretion disks, which suggests the existence of a scale-independent physical mechanism responsible for the ejection and collimation of these outflows. The presently more accepted models are the magnetocentrifugal models (Blandford \\& Payne \\cite{BP1982}; Uchida \\& Shibata \\cite{US1985}), in which the ejection is driven by the presence of a dynamically important magnetic field in the accretion disk-central object system, and the collimation of the jet is due to the toroidal component of the magnetic field, which is able to collimate the outflows by pinching forces. The toroidal magnetic field is generated by the twisting of the magnetic field due to the rotation of the system. The region where this process acts is too close to the central object to be resolved observationally, and one possible way to obtain some insight into this region is by studying the properties of the outflows.  In particular, a lot of progress has recently been made regarding observations of the outflows from YSOs (see the review by Reipurth \\& Bally \\cite{RB2001}). The caracteristic spectral emission of these objects is believed to come from the region behind the shock, from recombination of the ionized gas (for the hydrogen lines) and electron excitation (and de-excitation) within ions (Schwartz \\cite{S1975}), and the typical knot structure visible along the jet could be interpreted as due to a time-dependent ejection from the young stellar object (YSO) (Reipurth \\cite{R1989}; Raga et al.\\cite{Ral1990}).  In the past few years, several authors have studied the effect of the magnetic field on the dynamical evolution of HH objects (e.~g. O'Sullivan \\& Ray \\cite{OR2000}; Stone \\& Hardee \\cite{SH2000}; Cerqueira et al. \\cite{Cal1997}). Frank et al. (\\cite{Fal1999}) showed that ambipolar diffusion could be important to smear out the magnetic field ejected with the jets, but only on timescales comparable or larger than the dynamical timescales of the jets. The determination of a magnetic field in the outflows would represent a test for the magnetocentrifugal mechanism. In particular, the presence of a dynamically important toroidal component of the magnetic field would represent an indirect proof of this mechanism. However, direct observations of magnetic fields in jets are very difficult, and there is yet no clear observational determination of the magnetic field intensity in jets.  Several hydrodynamic (HD) simulations with calculations of the spectral emission of HH objects have been presented in the past (e.~g. Blondin et al. \\cite{Bal1990}; Raga \\cite{R1994}), but such calculations have never been published for the magnetized case. The simulations with magnetic fields have concentrated on the dynamical aspects and the evolution of the jet rather than on obtaining predictions of the emitted spectrum. These simulations usually include a radiative cooling rate (given, e.g., by the coronal cooling function of Dalgarno \\& McCray \\cite{DM1972}), and different magnetic field configurations. The main features found in magnetized jets with respect to HD jets are the presence of a ``nose cone'', better collimation, an increase in the density along the direction of propagation, and some effects on Kelvin-Helmoltz and Rayleigh-Taylor instabilities (resulting in changes in the leading bow shock, e.~g. Todo el al \\cite{Tal1993}; Cerqueira \\& de Gouveia Dal Pino \\cite{CD1999}).  With respect to the emitted spectrum, Hartigan, Morse \\& Raymond (\\cite{Hal1994}) compared observed and predicted emission lines ratios (using plane-parallel shock models) to find an upper limit of 30 $\\mu$G for the magnetic field of the jet. Cerqueira \\& de Gouveia Dal Pino (\\cite{CD2001a}), using a semiempirical formula (valid for a shock velocity between $20$ and $80$ km s$^{-1}$) found the ratio between the H$\\alpha$ emission of a magnetized and a non-magnetized jet. They obtained that the H$\\alpha$ emission increases due to the presence of a magnetic field, and that the dominant cause of this increase is the toroidal component of the magnetic field.  Therefore, no direct calculation of the effect of the magnetic field on the emission from a jet has yet been made. Trying to fill this gap, we have carried out 2D, axisymmetric, MHD simulations of jets, looking for the differences in the predicted H$\\alpha$ emission from magnetized jets with respect to the hydrodynamic case. The paper is organized as follows. In section 2, we explain in some detail the numerical algorithm that has been used, the initial conditions of the simulations and the approximations used to calculate the emission. In section 3 we summarize and discuss the results obtained, and in section 4 we draw our conclusions.   \\begin{figure*} \\centering \\includegraphics[width=18cm]{densnew.eps} \\caption{Numerical density for the models HD (hydrodynamic), weak MHD ($\\beta=1$), and strong MHD ($\\beta=0.4$), after a 500 yrs integration time. The right bar gives the logarithm of the numerical density (in cm$^{-3}$) The axes are labeles in pixels, and the displayed domain has a physical size of ($L_z$,$L_r$) = ($3\\times 10^{17}$, $3\\times 10^{16}$) cm. A pixel correspond to a physical size of $1.66 \\times 10^{14}$ cm.} \\label{fdens} \\end{figure*} ",
        "conclusions": "We have presented 2D numerical simulations of HD and magnetized, variable jets propagating in a homogeneous medium. Models with increasing magnetic fields show an increase in the H$\\alpha$ luminosities of the successive knots (which correspond to internal working surfaces which result from the injection velocity variability). This result confirms the work of Cerqueira \\& de Gouveia Dal Pino (\\cite{CD2001a}, \\cite{CD2001b}), who estimated the H$\\alpha$ luminosity of the clumps along MHD jet simulations using a fit to predictions of plane-parallel shock models.  Somewhat surprisingly, our work presents the first predictions of emission line maps from MHD HH jet models. Therefore, our calculations for the first time show the emission line morphologies that would be expected for such models.  We find that the H$\\alpha$ emission of the leading head of the jet differs quite strongly between the HD and MHD cases. This is a result of the fact that our simulations develop extended ``nose cones'' (of somewhat dubious reality, as these structures might disappear in 3D jets without perfect axisymmetry, see Cerqueira \\& de Gouveia Dal Pino \\cite{CD2001b}).  For the knots along the jet, we find that for increasing magnetic field strengths we obtain emission structures with stronger peaks towards the symmetry axis. This can be clearly seen in Fig. 3, in which the $\\beta=0.4$ model has knots which are dominated by an elongated emission component along the jet axis.  This different type of knot morphology is interesting in terms of observations of HH jets. It has been a long-standing fact that while some HH jets (notably HH~111, see e.~g. Reipurth et al. \\cite{Ral1997}) show compact knots with ``bow shock-like'' morphologies which resemble the predictions from HD variable jet models (see Masciadri et al. \\cite{Mal2002}), other HH jets (e.~g., HH~30, see Lopez et al. \\cite{Lal1995}) have emission knots with axially elongated structures. This second kind of morphology could not be modeled successfully in terms of variable HD jet models, and suggested the presence of a different mechanism for knot formation. We now find that variable jets with a strong enough toroidal magnetic field do lead to the formation of axially elongated knots, which in principle might be used to model objects such as HH~30.  The present paper is limited to a study of the effect of a toroidal magnetic field on the H$\\alpha$ emission of variable jet flows. In a future paper, we will present a study of a more extended set of emission lines. From such a study, we will attempt to produce a set of line diagnostics which could be used to estimate the magnetic field strength along observed HH jets."
    },
    "0512/astro-ph0512220_arXiv.txt": {
        "abstract": "The project of a micro-TPC matrix of chambers of \\hetrois for direct detection of non-baryonic dark matter is presented. The privileged properties of \\hetrois  are highlighted. The double detection (ionization - projection of tracks) is explained and its rejection evaluated. The potentialities of MIMAC-\\hetrois for supersymmetric dark matter search are discussed ",
        "introduction": "In the last years our work on \\hetrois as a target for detecting WIMPs  allowed to confirm its privileged properties for direct detection \\cite{idm2002}. These properties can be enumerated  as follow : \\begin{itemize} \\item  its fermionic character opens the axial interaction with fermionic WIMPs as the neutralinos, \\item the extremely low Compton cross section reduces by several order of magnitude the natural radioactive background with respect to other targets, \\item the high neutron capture cross section gives a clear signature for neutron rejection, \\item its light mass allows a higher sensitivity to light WIMP masses than other targets, \\item the elastic energy transfer is bounded to a very narrow range of energy (a few keV) offering a high signal to noise ratio. \\end{itemize}  The extremely low Compton cross section and the possibility to detect events in the keV range ($\\leq $5.6 keV) have been demonstrated by the $^{57}$Co electron conversion detection recently reported \\cite{electrons}. The detection of  $\\rm 7 \\ keV$ electrons in the MACHe3 prototype with the source emitting 121 keV $\\gamma$-rays embedded in the \\hetrois is a clear demonstration of the virtual transparency of this medium to the electromagnetic radiation.      \\noindent ",
        "conclusions": ""
    },
    "0512/astro-ph0512016_arXiv.txt": {
        "abstract": "The deepest X-ray images of M31, obtained with  XMM-Newton, are examined to derive spectral and statistical properties of the population of the softest X-ray sources. Classifying supersoft X-ray sources (SSS) with  criteria based on  the same hardness ratios defined for recent Chandra observations, a quarter of the selected SSS  turn out to be supernova remnants (SNR). Another quarter of SSS are spatially coincident with recent classical novae (but they are less than 10\\% of the nova population observed in the last 25 years). Only 3 among 15 non-SNR SSS show clear variability with X-ray flux variation of more than one order of magnitude within  few months. Two of these sources display additional, smaller amplitude variability on time scales of several minutes.   Their broad band spectra and those of the novae are approximately fit with a blackbody or white dwarf atmospheric model at near-Eddington luminosity for the distance of M31.  Two SSS appear to reach very large, perhaps super-Eddington luminosities for part of the time, and probably eject material in  a wind until the luminosity decreases again after a few  months. One of the two objects has some characteristics in common with Ultra Luminous X-ray Sources observed outside the Local Group.  Most Quasi-Soft Sources (QSS),  among which also a few SNR are selected using the hardness ratio criteria, are repeatedly detected. Several QSS are better fit by a power law spectrum, but some faint, apparently blackbody-like QSS with temperatures T$_{\\rm bb}\\simeq$100-200 eV and luminosity 10$^{36}$ erg s$^{-1}$ at M31 distance do exist. I discuss the possibilities that most QSS may be SNR  in M31, or foreground neutron stars.  Two X-ray sources with both a soft and hard component are in the positions of a recurrent nova and another object that was tentatively classified as a symbiotic nova. These two sources may be black hole transients. ",
        "introduction": "One of the main discoveries of  Einstein and especially of ROSAT were the {\\it supersoft X-ray sources} (hereafter, SSS).  SSS emit detectable X-rays only at energy below 1 keV, with very large luminosity in the range 10$^{36}$-10$^{38}$ erg s$^{-1}$, and their spectrum is approximately fit with a blackbody at temperature in the range 150,000-1,000,000 K. Since these sources are very luminous, but very ``soft'', they are much more easily detected in Local Group galaxies,  towards which the column of neutral hydrogen N(H) is low, than in the Galaxy (see Greiner 2000 and references therein). The most luminous SSS are even detected outside the Local Group (Swartz et al. 2002, Di Stefano \\& Kong 2003,  2004, Kong \\& Di Stefano 2003 and 2005). Even if the whole sample of SSS, selected on purely phenomenological criteria, is not a homogeneous class, there is evidence that a large fraction of these sources in the Local Group   are extremely hot, accreting white dwarfs (WD) in close binaries, burning hydrogen in a shell (see Kahabka \\& van den Heuvel 1997).  Close binary SSS (CBSS) include post-outburst recurrent and classical novae, the hottest symbiotic stars, and other low mass X-ray binaries containing a WD, with typical orbital periods between 4 hours and 1 day.  The recently rekindled debate on the nature of the progenitors of type Ia supernovae (SNe Ia) focuses often on this class of sources, as prototypes of single degenerate binary SNe Ia progenitors  (e.g. van den Heuvel et al 1991., King et al. 2003, Yoon \\& Langer 2003, Starrfield et al. 2004). SSS appear often to be transient or recurrent in X-rays, or at least variable in X-ray flux (e.g. Greiner et al. 2000, 2004a).  Most CBSS are indeed expected to be variable. There are basically three mechanisms for variability of the binary SSS described above. 1) A periodic radius expansion is accompanied by variation in the mass transfer rate $\\dot m$, and feedback in the nuclear burning rate (e.g. Cal 83, Greiner \\& Di Stefano 2000),  recurrently causing increase the optical and UV luminosity to increase and the X-ray luminosity to decrease. 2) Repeated thermonuclear hydrogen shell flashes (unlike novae, without mass ejection) are expected in a regime of mass accretion rate, $\\dot m$, that varies depending on the model, but is around 10$^{-8}-10^{-7}$ M$_\\odot$ year$^{-1}$. Since the upper limits on the X-ray flux before the new generation of X-ray satellites were not very high and the observations of SSS were seldom repeated, we know very little about thermonuclear flashes. 3) At higher   $\\dot m$, all the energy produced in thermonuclear burning is immediately radiated (Fujimoto 1982, Kovetz \\& Prialnik 1994), and generally we expect to find persistent X-ray sources, which are the most likely candidates for progenitors of neutron stars born by accretion induced collapse (AIC) and/or of type Ia supernovae (see Yungelson et al. 1996, Starrfield et al. 2004). However, at high $\\dot m$ there is another mechanism for variability. A ROSAT source, RX J0513.9-6951, is optically bright and ``off'' in X-rays for the 140 days of a recurrent cycle, then it undergoes a rapid transition in less than 4 days to the supersoft X-ray source stage, with decreased optical luminosity. This stage lasts for about a month, until in less than 2 days, the source returns to the previous state (Reinsch et al. 2000). The optical variation is only 0.8 mag in V, while the ROSAT PSPC count rate was measured to vary by a factor of more than 20. According to Hachisu \\& Kato (2003) RX J0513.9-6951 burns hydrogen in shell at the constant rate necessary to reach the SN Ia explosion. When more mass transfer rate is triggered than the rate at which it can all be burned, the outer layers of the WD expand and mass accreted from the secondary is suddenly lost in a wind, causing the X-ray off state and optical brightening.  Independently of the final fate of a particular system, observing and studying shell-hydrogen burning   WD, we obtain a glimpse on the evolution of  potential  SN Ia progenitors in general, a central problem of modern astrophysics. Even what we know about the acceleration of cosmic expansion depends on type Ia SN.  M31 is the most luminous and massive galaxy in the Local Group, and as such provides a large selection of nearby extragalactic X-ray sources, including SSS. At visual wavelength it is almost five times as luminous as the other spiral M33, and almost 16 and 80 times more luminous as our two nearby satellites, respectively the LMC and the SMC (Sparke \\& Gallagher 2000). The content of neutral hydrogen in M31 is estimated to be $\\simeq 5.7 \\times 10^9$ M$_\\odot$, about 30\\% more than the Galaxy. The central bulge is affected by a relatively low column density of neutral hydrogen, N(H)$\\approx 8 \\times 10^{20}$ cm$^{-2}$, and the surrounding region has only N(H)$\\simeq 10^{21}$ cm$^{-2}$, as it is shown  Supper et al.'s (1997) simplified model derived from Unwin's (1981) maps. Most of the neutral hydrogen is concentrated in a thick star-forming ``ring of fire'' around it, with N(H)=7.7 $\\times 10^{21}$ cm$^{-2}$, but the gas extends to a large radius, and it has  a rather patchy structure, with both ``holes'' of low column density  and dark clouds causing higher column density in small regions (see Davies et al., 1976, and Hodge 1981).  Due to its large stellar population and many regions with low N(H),  M31    is ideal for investigating SSS. The M31 population was studied with  ROSAT by Supper et al (1997), Kahabka (1999) and Greiner et al. (2004), with  Chandra by De Stefano et al. (2004) and Greiner et al. (2004). XMM-Newton is the best suited satellite for this type of research outside the central region of the bulge (where the excellent spatial resolution of Chandra is very important), due to the high effective area, high sensitivity and quite good calibration of the EPIC detectors in the soft range. A comprehensive study of the XMM-Newton observations of SSS in M31 has not been published. De Stefano et al. (2004) and Greiner et al. (2004) examined most of the XMM-Newton exposures of M31, and recovered detections of the Chandra SSS as well as some originally found with ROSAT. Using criteria that are as equivalent as possible to  the ones adopted by these authors, I performed an unbiased search for  SSS in M31 in the XMM-Newton observations, using also two recent public images that were not yet available to the above authors. Moreover, I searched for common detections of a second large subset of ROSAT SSS identified by Kahabka (1999). In addition to a variability study, one of my aims is to make use of the large effective area and higher count rates of XMM-Newton to obtain as many broad band spectra as possible, to derive conclusions on the physical nature of these sources.  Last but not least, I also examined another class of sources, called Quasi-Soft-Sources (QSS) by Di Stefano et al. (2004), found in M31 and in other galaxies (Di Stefano \\& Kong 2004). These sources were discovered relaxing the search criteria of SSS based on the hardness ratios and extending them to sources with slightly harder spectra. Di Stefano et al. (2004)  define QSS spectra as blackbodies with temperatures in the range 100-350 eV. It is important to bear in mind that this class of sources are not necessarily related to SSS, nor is there any evidence yet that some of them may  be accreting WD.  Using the XMM-Newton data, I also investigated the nature of QSS: are they really a new kind of X-ray sources with uniform physical characteristics ? ",
        "conclusions": "I have performed a systematic search of SSS and QSS in the XMM-Newton observations of M31. When possible, I have studied the spectra and the time variability characteristics  of these sources and searched counterparts at other wavelengths. These are a few conclusions that can be drawn from this study:  1) This ``close look'' M31 SSS population clearly confirms that, as noted in other studies, using only hardness ratio criteria we  cannot select a uniform type of sources undergoing the same physical mechanisms.  2) Using hardness ratios based on broad spectral bands like in Di Stefano et al. (2004), at least 11 SNR are selected as  SSS. However, SNR tend to have on average harder spectra than H-burning WD in CBSS, so the number of SNR among SSS is reduced if hardness ratios with narrower band-passes are used, like for ROSAT (Supper et al. 1997) or in the XMM-Newton catalog (Pietsch et al. 2005a). The presence of SNR in  the SSS data bases in external galaxies may be assessed not only with multiwavelength observations, but especially complementing Chandra observations, necessary for the good spatial resolution, with XMM-Newton observations, in order to detect more signal in the soft band.  3) Post-outburst novae make up a  quarter of all  SSS, but only less than 1 in 10 nova of the last 25 years is detected as  SSS. Up to now, in M31 like in the Galaxy and the in the Magellanic Clouds, no nova was detected in X-rays after more than 10 years.  4) 1 out of 10 novae in outburst in M31 is associated with a SSS, indicating ongoing shell hydrogen burning. Yungelson et al. (1996) evaluated the SNe  Ia rate, the nova rate and the interacting binary formation rate (e.g. Yungelson et al. 1996). Following their reasoning, we  find that, if only 10\\% classical and recurrent novae keeps on burning hydrogen in a shell for several years, these systems are not the main class of progenitors of type Ia supernovae.  5) Large flux variations by more than one order of magnitude, consistent with the limit cycle of thermonuclear flash models or with ``wind-regulated'' sources, are detected in 20\\% of those  SSS that are not SNR (3 out of 15 sources), allowing for the possibility that they are CBSS, and that a relatively large number of CBSS may be burning hydrogen in a shell at the high rate required by type Ia supernova models,  6) At least one variable  SSS has a very large luminosity, of a few 10$^{39}$ erg s$^{-1}$ at M31 distance; the spectrum is surely not a simple blackbody and must be produced by two or more components. This source, detected with Chandra as well, may be the missing link with ULX of external galaxies. This is a very interesting finding. I dare suggest that monitoring this close-by source in Andromeda we may be able to find a substantial clue to the nature of ULX.  7) Looking at the present statistics, it seems that QSS may include a large number of SNR candidates, although it is puzzling that counterparts at FUV/NUV wavelengths are not detected.  QSS may also include foreground neutron stars, or even softer-than-usual AXP in M31.  8) Cross correlating the novae with with the positions of all X-ray sources observed with XMM-Newton (not only X-ray sources), two non- supersoft X-ray sources have been found.  Both are variable, one can be classified as ``transient'' and seems to be variable not only in flux but also in spectral characteristics. These two objects are candidate black-hole transients. It will be interesting to monitor the X-ray behaviour of  X-ray sources that are spatially coincident with classical novae, not only in order to study H-burning WD, but also with the aim of selecting black-hole transients in the Local Group galaxies.  I would like to conclude reminding that M31 gives us a wonderful opportunity to obtain statistics of whole classes of X-ray sources, and that both XMM-Newton (because of the large effective area)  and Chandra (because of the low background and especially of the high spatial resolution, of unique value in the very crowded central region) are exceptional instruments to study X-ray sources populations and monitor the variable sources."
    },
    "0512/astro-ph0512002_arXiv.txt": {
        "abstract": "We address the issue of electromagnetic pulsar spindown by combining our experience from the two limiting idealized cases which have been studied in great extent in the past: that of an aligned rotator where ideal MHD conditions apply, and that of a misaligned rotator in vacuum.  We construct a spindown formula that takes into account the misalignment of the magnetic and rotation axes, and the magnetospheric particle acceleration gaps.  We show that near the death line aligned rotators spin down much slower than orthogonal ones.  In order to test this approach, we use a simple Monte Carlo method to simulate the evolution of pulsars and find a good fit to the observed pulsar distribution in the $P-\\dot{P}$ diagram without invoking magnetic field decay. Our model may also account for individual pulsars spinning down with braking index $n<3$, by allowing the corotating part of the magnetosphere to end inside the light cylinder. We discuss the role of magnetic reconnection in determining the pulsar braking index.  We show, however, that $n\\sim 3$ remains a good approximation for the pulsar population as a whole. Moreover, we predict that pulsars near the death line have braking index values $n> 3$, and that the older pulsar population has preferentially smaller magnetic inclination angles. We discuss possible signatures of such alignment in the existing pulsar data. ",
        "introduction": "The current canonical pulsar paradigm is that of a magnetized rotating neutron star (see Mestel~1999 for a review). However, we feel that certain fundamental aspects of the paradigm still remain unclear.   One aspect of the paradigm which we hope to elucidate in the present work has to do with the way the neutron star spins down. A spinning down neutron star with mass $M_*$, radius $r_*$, and angular velocity $\\Omega$ loses rotational kinetic energy at a rate \\begin{equation} L= \\frac{2}{5}M_* r_*^2 \\Omega \\dot{\\Omega}\\ . \\label{kinetic} \\end{equation} Here, $M_*\\sim 1.4 M_{\\odot}$, $r_*\\sim 10$~km, and $\\dot{(...)}\\equiv {\\rm d}(...)/{\\rm d}t$.  Energy is lost through electromagnetic torques in the magnetosphere, although other physical processes have at times also been discussed (gravitational radiation, wind outflow, star-disk interaction, etc.).  To a first approximation, the stellar magnetic field may be considered as that of a rotating magnetic dipole.  Even under such a simplification the general description of the stellar magnetosphere is a formidable three dimensional problem, since, in general, the magnetic and rotation axes do not coincide. Awaiting the development of the general theory, one can still derive important conclusions based on two idealized limiting cases which have been studied in great extent: the case of an aligned magnetic dipole rotating in an atmosphere with freely available electric charges (i.e. with ideal MHD conditions), and that of a misaligned magnetic dipole rotating in vacuum.  The neutron star is not surrounded by vacuum, and one needs to take into consideration the electric fields that develop and the electric currents that flow in the rotating charged magnetosphere (Goldreich \\& Julian~1969).  The most recent numerical calculation of the simplest possible case, that of the magnetosphere of an aligned rotator in force-free approximation (Contopoulos~2005, hereafter C05), yielded the following rather general result for the electromagnetic energy loss \\begin{equation} L_{\\rm aligned}= \\frac{4\\Omega\\Omega_{F}\\psi_{\\rm open}^2}{6c}\\ \\label{C05} \\end{equation} (the quantities $\\Omega_F$ and $\\psi_{\\rm open}$ are defined below).  In that picture, the magnetosphere consists of a corotating region of closed fieldlines which extends up to a distance $r_c$ from the rotation axis, and an open fieldline region with enclosed magnetic flux \\begin{equation} \\psi_{\\rm open}\\equiv \\frac{1}{2\\pi}\\int_{\\rm open} {\\bf B}\\cdot {\\rm d}{\\bf S} =1.23\\frac{B_* r_*^3}{2r_c}, \\label{psi1} \\end{equation} where $B_*$ is the polar value of the magnetic field (Contopoulos, Kazanas \\& Fendt~1999, hereafter CKF; Gruzinov 2005; C05; Timokhin 2005; see also Appendix~A).  The above expression is valid when $r_c\\gg r_*$.  In the limit $r_c=r_*$, straightforward calculation yields $\\psi_{\\rm open}=B_* r_*^2/2$.  The neutron star spins down because of the establishment of a large scale poloidal electric current circuit flowing along open field lines, and returning along the edge of the open field line region (see CKF for a detailed description).  The electric current that flows between the magnetic axis (characterized by $\\psi=0$) and the edge of the open field line region (characterized by $\\psi=\\psi_{\\rm open}$) generates the spindown torque which leads to eq.~\\ref{C05}.  The quantity $\\Omega_{F}$ in eq.~\\ref{C05} is the angular frequency of rotation of the open field lines.  It is set by the electric potential drop that develops accross open field lines, between the magnetic axis and the edge of the open field line region. This potential in the magnetosphere is in general {\\rm smaller} than the corresponding electric potential drop on the surface of the star. The difference between the two potential drops is just the particle acceleration gap potential which develops {\\em along} open magnetic field lines in the vicinity of the polar cap. Consequently, the angular velocity of open field lines will in general be different (smaller) than $\\Omega$.  Models of particle acceleration and pair creation of rotation-powered pulsars yield values of the gap potential $V_{\\rm gap}$ of the order of $10^{12}$~Volts (e.g. Hibschmann \\& Arons 2001).  One can directly show (see Appendix~B) that the above together with the simplifying assumption that $\\Omega_F$ is uniform accross open field lines, yield \\begin{equation} \\Omega_F=\\Omega-\\Omega_{\\rm death}\\ , \\label{OmegaF} \\end{equation} where, \\begin{equation} \\Omega_{\\rm death}\\equiv \\frac{V_{\\rm gap}}{\\psi_{\\rm open}}c\\ . \\label{om_death} \\end{equation} This describes the so-called pulsar `death', i.e. the stopping of pulsar emission. As the neutron star slows down and $\\Omega$ drops below $\\Omega_{\\rm death}$, the gap potential cannot attain the value required for particle acceleration and consequent pulse generation, and the pulsar stops generating radio emission.  A misaligned dipole rotating in vacuum loses energy at a rate \\begin{equation} L_{vacuum}=\\frac{B_*^2\\Omega^4r_*^6}{6c^3}\\sin^2\\theta\\ , \\label{dipole1} \\end{equation} where $\\theta$ is the misalignment angle between the magnetic and rotation axes.  We know that in real life the neutron star is not surrounded by vacuum, and we may argue that, in analogy to the aligned case, the magnetosphere consists of a corotating and an open line region. We may thus rewrite eq.~\\ref{dipole1} in a more general form that expresses the energy loss rate of an orthogonal ($\\theta=90^o$) magnetic rotator as \\begin{equation} L_{\\rm orthogonal} \\sim\\frac{4\\Omega^2\\psi_{\\rm open}^2}{6c}\\ . \\label{dipole} \\end{equation}   We would like to emphasize at this point that eqs.~\\ref{C05} \\& \\ref{dipole1} being so similar, led most researchers to ignore the dependence on $\\theta$ and $\\Omega_F$ in estimates of stellar magnetic fields, and in most studies of the $P-\\dot{P}$ diagram.  The aim of the present work is to show that the dependence on $\\theta$ and $\\Omega_F$ is important and should not be ignored, especially in old pulsars approaching their death.  As we will see, pulsar death manifests itself in a most interesting way through its dependence on $\\Omega_F$.  Moreover, the $\\theta$ dependence `softens' the distribution of pulsars around the death line in the $P-\\dot{P}$ diagram. ",
        "conclusions": "Awaiting the development of a detailed three dimensional MHD theory for the rotating neutron star magnetosphere, we can describe a few general characteristics of its expected structure.  We argue that $\\Omega_F$ characterizes the reduced magnetospheric electric potential drop between the magnetic axis and the edge of the open field line region.  Poloidal electric currents will be generated as long as electric charges can be produced in the magnetosphere.  These charges are produced in the magnetospheric polar gaps.  In the axisymmetric case, $\\Omega_F$ can also be thought as the angular velocity of rotation of open magnetic field lines, which is in general smaller than $\\Omega$.  The electric current that flows between the magnetic axis and the edge of the open field line region generates the spindown torque. As shown in C05, this electric current and spindown torque are both proportional to $\\Omega_F$.  We argue that this picture describes also the situation when $\\theta\\neq 0$. As the neutron star rotates, the magnetic axis moves around the axis of rotation. In the rotating frame of the star, however, open field lines rotate at their own rate $\\Omega-\\Omega_F$ opposite to the direction of stellar rotation around the magnetic axis.  Observational evidence for this effect can be found in the well known sub-pulse drift phenomenon (e.g. Beskin 1997; Rankin \\& Wright 2003).  In the limit $\\Omega_F=\\Omega$, it is as if open magnetic field lines are anchored on the stellar surface.  In the opposite limit $\\Omega_F=0$, open field lines rotate around both the magnetic and rotation axes returning to the same position every rotation of the star when viewed by an inertial observer. For finite $\\Omega_F $ a phase shift about the magnetic axis will be accumulated after every turn of the star.  A simple analogy to the above picture might be that of a lawn watering system consisting of a hose rotating at angular velocity $\\Omega$, and a sprinkler at its end rotating with angular velocity $\\Omega_F-\\Omega$ in the rotating frame of the watering hose.  Note that the orthogonal component does not require the establishment of a poloidal electric current circuit in order for it to slow down the stellar rotation. The orthogonal component emits spiral electromagnetic waves (in general Alfven waves) which travel out to infinity through the rotating or non-rotating open field lines.  The main conclusions of the present work are: \\begin{enumerate} \\item The electromagnetic energy loss of a pulsar depends not only on the surface magnetic field $B_*$ and on the rotational angular velocity $\\Omega$, but also on the misalignment angle $\\theta$ and on the angular velocity of the open magnetic field lines $\\Omega_F$ (which is in general smaller than $\\Omega$). \\item The approach to pulsar death modifies the rate of energy loss. The death occurs when the pulsar slows down sufficiently so that the angular velocity of the open fieldlines $\\Omega_F=0$. \\item The energy loss close to the pulsar death is smaller than what is given by the standard dipolar spindown formula. This effect gives a good fit between the theoretical and observed distributions of pulsars near the death line, without invoking a magnetic field decay. \\item Our model may also account for individual pulsars spinning down with braking index $n<3$. However, $n\\sim 3$ remains a good approximation for the pulsar population as a whole. \\item Pulsars near the death line have braking index values $n> 3$. Such high braking index values may be observable. \\item Pulsars near the death line may have preferentially smaller inclination angles. \\end{enumerate}  A preliminary look at the ATNF pulsar catalog data suggests that the last point may have some observational support. One possible measure of the inclination angle of a pulsar is its fractional pulse width, or the ratio of the width of the pulse to the period of the pulsar. If the radio beam size is independent of inclination, then we would expect pulsars with smaller inclinations to be seen for a larger fraction of the period than pulsars with large inclinations. In order to test this hypothesis we took the available data for the pulse width at 50$\\%$ of the pulse peak from the ATNF catalog (1375 pulsars with $W_{50\\%} \\neq 0$).  In order to be able to compare the fractional pulse width for pulsars of different periods, we have to correct for the intrinsic size of the pulsar beam.  If the beam roughly follows the angular size of the open fieldlines, then the beam size falls as $P^{-1/2}$ with increasing period (we are assuming that the last closed field line extends out to the light cylinder). Therefore, the quantity that should relate to the degree of alignment is $F_{\\rm align} \\equiv (W\\times P^{1/2})/P=W P^{-1/2}$, where $W$ is the measured pulse width. In figure~\\ref{fig51} we plot the observed pulsars as circles with the radius of the circle linearly proportional to $F_{\\rm align}$.  A visual inspection of the plot shows that there is an excess of larger pulse fractions for older pulsars, and in particular for pulsars near the right edge of the $P-\\dot{P}$ triangle. This region is near the pulsar death line (dashed line), and therefore, this feature is particularly interesting. Obviously, a much better analysis needs to be performed, but these results are quite encouraging.  An overabandunce of pulsars with large pulse fractions near the death line in a smaller pulsar sample had also been interpreted by Lyne and Manchester (1988) as an indication of alignment.  In an effort to simplify the fits we have assumed in this paper that the misalignment angle stays the same during the evolution, and there is no field decay. The reality may of course be more complicated, and there could be some amount of field decay, and potential alignment (and consequent precession) during the lifetime of a pulsar, due to electromagnetic torques on the star. These effects would introduce extra degrees of freedom to fitting the pulsar distribution. We hope, however, that the physically motivated spindown law of the type introduced in this work would find its way into detailed population synthesis models."
    },
    "0512/astro-ph0512528_arXiv.txt": {
        "abstract": "{Heat conduction plays an important role in the balance between heating and cooling in many astrophysical objects, e.g. cooling flows in clusters of galaxies. Here we investigate the effect of heat conduction on the interaction between a cool disk and a hot corona around black holes.  Using the one-radial-zone approximation, we study the vertical structure of the disk corona and derive evaporation and coronal mass flow rates for various reduced thermal conductivities. We find lower evaporation rates and a shift in the evaporation maxima to smaller radii. This implies that the spectral state transition occurs at a lower mass flow rate and a disk truncation closer to the black hole. Reductions of thermal conductivity are thought to be magnetically caused and might vary from object to object by a different configuration of the magnetic fields. ",
        "introduction": "Heating by thermal conduction is an important process in many astrophysical objects. The balance between cooling and heating processes determines the structure of hot matter in contact with cooler regions. For the case of a dwarf nova accretion disk, we investigated this interaction (Meyer \\& Meyer-Hofmeister 1994) and found that a ``siphon flow'' leads to evaporation of the disk. The same happens in disks around neutron stars and in both galactic and supermassive black holes (Meyer-Hofmeister \\& Meyer 1999, Meyer et al. 2000). The heat conduction is an essential element; otherwise the coronal gas would not lose energy, but would instead become hotter and reach virial temperatures as it is present in advection-dominated flows.  Recently the same physical process was discussed by several authors in connection with the cooling flow problem for the hot intra-cluster medium (ICM). The work of Medveden \\& Narayan (2001) focuses on the question as to what degree chaotic magnetic fields suppress conduction relative to the Spitzer level. They find that thermal conduction in a weakly collisional plasma with turbulent magnetic fields approaches the Spitzer limit. Zakamska \\& Narayan (2003) find from the investigation of five galaxy clusters that the Spitzer formal with a conduction coefficient reduced to about 30\\% gives a good description of the observed radial profiles of electron density and temperature. Ghizzardi et al. (2004) discuss which fraction of the Spitzer value would be appropriate xplaining the Virgo/M87 observations.  Voigt \\& Fabian (2004), on the other hand, found support for an unhindered heat conduction from their analysis of 16 galaxy clusters using Chandra data. One third of the Spitzer value was used in hydrodynamic cosmological simulations of galaxy clusters by Dolag et al. (2004). In this context Soker et al. 2004 suggest a heat conduction along magnetic field lines.  Okabe \\& Hattori (2004) suggest a suppression of heat conduction by magnetic fields generated most strongly in the direction perpendicular to the temperature gradient. All these results point to a reduced heat conduction in many sources.  In the original context of disk/corona interaction, the effect of reduced thermal conduction has already been considered by a scaling procedure (Meyer et al.2000). In the present work we evaluate the effect of heat conduction on the evaporation of accretion disks in detail. Connected with the evaporation efficiency, this means a possible change in the resulting truncation radius. Motivation also comes from the fact that new computations including the irradiation of the coronal gas from the inner region (Meyer-Hofmeister et al. 2005 (hereafter MLM05), Liu et al. 2005) have led to radii that seem larger than indicated by observations (Yuan \\& Narayan 2004). In these works the thermal conduction was taken according to the standard value derived by Spitzer (1962). The question then arises whether a reduced heat conduction could be present in the disks around black holes. Here we study the effect of heat conduction on evaporation and the truncation of the accretion disk.  In Sect. 2 we give a short description of the accretion geometry and the interaction of corona and disk. In Sect. 3 we present the results for reduced heat conduction, and a discussion of the consequences and conclusions follow in Sect. 4. ",
        "conclusions": "We have evaluated how a reduced heat conduction affects the evaporation of gas from a cool disk to a hot coronal flow/ADAF. The physical situation of evaporation (or condensation) between hot (``coronal'') and cool (``disk'') gas is the same around stellar black holes and in galaxies and clusters of galaxies. For clusters of galaxies, reduced heat conduction seems to be supported by observations in several cases.  In our theoretical modeling, evaporation rates become lower with reduced heat conduction, and the location where the evaporation efficiency reaches its maximum moves inward by a factor of 7 for the reduction to 20\\%. This moves the truncation radii closer to an agreement with observations, but a significant difference exists still. On the other hand, the change of truncation radii with accretion rate $\\Delta \\log R_{\\rm{tr}}/ \\Delta \\log \\dot M$ in our results is in reasonable agreement with the numbers derived from observations for galactic and supermassive black holes.  Interesting is the strong dependence of spectral transition on heat conduction. A further reduction might arise from a different magnetic field situation. This can affect the transition from a very bright state to a very dim state in the AGN of elliptical galaxies as suggested by Churasov et al. (2005)."
    },
    "0512/astro-ph0512234_arXiv.txt": {
        "abstract": "We employ high-resolution dissipationless simulations of the concordance $\\Lambda$CDM cosmology ($\\Omega_0=1-\\Omega_{\\Lambda}=0.3$, $h=0.7$, $\\sigma_8=0.9$) to model the observed luminosity dependence and evolution of galaxy clustering through most of the age of the universe, from $z\\sim 5$ to $z\\sim0$.  We use a simple, non-parametric model which monotonically relates galaxy luminosities to the maximum circular velocity of dark matter halos ({\\Vmax}) by preserving the observed galaxy luminosity function in order to match the halos in simulations with observed galaxies.  The novel feature of the model is the use of the maximum circular velocity at the time of accretion, {\\Vin}, for subhalos, the halos located within virial regions of larger halos.  We argue that for subhalos in dissipationless simulations, {\\Vin} reflects the luminosity and stellar mass of the associated galaxies better than the circular velocity at the epoch of observation, {\\Vnow}.  The simulations and our model $L-${\\Vmax} relation predict the shape, amplitude, and luminosity dependence of the two-point correlation function in excellent agreement with the observed galaxy clustering in the SDSS data at $z\\sim 0$ and in the DEEP2 samples at $z\\sim 1$ over the entire probed range of projected separations, $0.1<r_p/(h^{-1}$ Mpc$)<10.0$.  In particular, the small-scale upturn of the correlation function from the power-law form in the SDSS and DEEP2 luminosity-selected samples is reproduced very well. At $z\\sim 3-5$, our predictions also match the observed shape and amplitude of the angular two-point correlation function of Lyman-break galaxies (LBGs) on both large and small scales, including the small-scale upturn.  This suggests that, like galaxies in lower redshift samples, the LBGs are fair tracers of the overall halo population and that their luminosity is tightly correlated with the circular velocity (and hence mass) of their dark matter halos. ",
        "introduction": "\\label{s:intro}  A generic prediction of high-resolution simulations of hierarchical Cold Dark Matter (CDM) models is that virialized regions of halos are not smooth, but contain subhalos --- the bound, self-gravitating dark matter clumps orbiting in the potential of their host halo. The subhalos are the descendants of halos accreted by a given system throughout its evolution, which retain their identity in the face of disruption processes such as tidal heating and dynamical friction. Their presence is in itself a vivid manifestation of the hierarchical build-up of halo mass.  In the CDM scenario, luminous galaxies form via cooling and condensation of baryons in the centers of the potential wells of dark matter halos \\citep{White78,Fall80,Blumenthal84}. In the context of galaxy formation, there is little conceptual difference between halos and subhalos, because the latter have also been genuine halos and sites of galaxy formation in the past, before their accretion onto a larger halo. We thus expect that each subhalo of sufficiently large mass should host a luminous galaxy and this is indeed supported by self-consistent cosmological simulations \\citep[e.g.,][]{Nagai05}. The observational counterparts of subhalos are then galaxies in clusters and groups or the satellites around individual galaxies.  In this sense, we will use the term {\\it halos} to refer to both distinct halos (i.e., halos not located within the virial radius of a larger system) and subhalos.  Although this general picture is definitely reasonable, it is not clear just how direct the relation between halos and galaxies is.  One may argue, for example, that the subhalos can be disrupted much faster than the more tightly bound stellar system they host, leaving behind ``orphan'' galaxies (\\citealt{Gao04a}; \\citealt*{Diemand04}). At the same time, properties of surviving subhalos, such as maximum circular velocity and gravitationally bound mass, are subject to strong dynamical evolution as they orbit within the potential of their host halo \\citep[e.g.,][]{Moore96,Klypin99,Hayashi03,Kravtsov04b,Kazantzidis04b}. This makes the relation between subhalo properties and galaxy luminosity ambiguous \\citep{Nagai05}, because the latter may be less affected by dynamical processes but may evolve due to aging of the stellar populations after ram pressure strips the existing gas and the accretion of new gas is suppressed.  The key question that we address in this paper is whether there is a one-to-one correspondence between populations of halos in dissipationless cosmological simulations and galaxies in the observable universe. As a test, we use comparisons of the predicted clustering of halos with the available observational measurements of galaxy clustering from $z\\sim 5$ to the present.  During the last decade, large observational surveys of galaxies both at low and high redshifts have tremendously improved our knowledge of galaxy clustering, its evolution, and the relation between the galaxy and matter distributions.  A coherent picture has emerged in which bright galaxies are strongly biased with respect to the matter distribution at high redshifts \\citep{Steidel98,Giavalisco98,Adelberger03, Adelberger05, Ouchi04b, Ouchi05, Lee05, Hamana05}, and in which the bias decreases with time in such a way that the amplitude of galaxy clustering is only weakly evolving \\citep[e.g.,][]{Ouchi04b}, as expected in hierarchical structure formation \\citep{Colin99, Kauffmann99}.  The bias is also in general scale- and luminosity-dependent.  Bright (red) galaxies are more strongly clustered than faint (blue) galaxies both in the local universe \\citep[][and references therein]{Norberg02a, Zehavi04, Zehavi05} and in the distant past \\citep{Coil04a, Coil05a}.  A recent development is the detection of a departure from a pure power law in the two-point correlation function of galaxies at $z\\sim 0$ \\citep{Zehavi04, Zehavi05}. This departure is expected in general because the two-point correlation function is a sum of two separate contributions: the one-halo term, which arises from pairs of galaxies within a distinct dark matter halo, and the two-halo term, which arises from pairs of galaxies from two different distinct halos \\citep[e.g.,][]{Cooray02}. The one-halo contribution dominates on small scales, while at scales larger than the size of the largest virialized regions clustering is due to the two-halo term. The two terms are not generically expected to combine so as to give a power-law correlation function. The deviation of the correlation function from a power-law was {\\it predicted} to be even stronger at higher redshifts \\citep{Zheng04,Kravtsov04}, and this has now been convincingly confirmed \\citep{Adelberger05, Ouchi05, Lee05, Hamana05} using galaxy samples identified with the Lyman-break technique \\citep{Steidel96,Steidel99}.  Despite impressive advances in the amount and quality of data on the galaxy distribution over a wide range of redshifts, the exact relation between dark matter halos and luminosity- or Lyman-break-selected galaxies is still rather uncertain \\citep[e.g.,][]{Mo99,Kolatt99, Wechsler01}.  The most popular attempts to connect the dark halos and luminous galaxies employ semi-analytic modeling (e.g., \\citealt{White91,Kauffmann94}; \\citealt{AvilaReese98, Somerville99,Cole94,Cole00, Croton05,Bower05}), which uses phenomenological recipes for specifying when, where, and how galaxies form within dark matter halos, often in conjunction with high-resolution dissipationless simulations. Another popular approach is to use the halo model \\citep[see e.g.][and references therein]{Cooray02}, which, in its simplest form, specifies the probability distribution for a particular halo of mass $M$ to host a given number of galaxies $N$ with specified properties, such as luminosity, color, etc.  More complex halo models include the conditional luminosity function (CLF) approach \\citep{Yang03,Cooray05a, Cooray05b} which specifies the luminosity function for halos of mass $M$, and models which connect other features \\citep*[e.g.,][]{Vale04, Neyrinck05}.  While these approaches manage to capture the general observational trends, they usually employ a large number of free parameters, making it difficult to glean the relevant underlying mechanism(s) responsible for the agreement between the model and data.  Most implementations also do not provide a direct constraint on the relation between luminous components of galaxies and dark matter halos. In fact, {\\it assumptions} are made about this relation in most cases.  In this study we use galaxy clustering to address two straightforward issues.  First, we question whether current data is consistent with a one-to-one correspondence between luminous galaxies and galaxy-sized dark matter halos in cosmological simulations. Second, we ask whether the observed clustering is consistent with a simple relation between galaxy luminosity and some property of its host halo.  In the context of the second question, we investigate which halo properties are most closely related to galaxy luminosity. We show with the currently available data the answer to both of these questions is yes, as we can reproduce luminosity-dependent clustering measurements at different redshifts from $z\\sim 5$ to $z\\sim 0$ with a simple, non-parametric model relating galaxy luminosity to the halo circular velocity.  Although a number of studies during the last decade have shown that galaxy clustering can be approximately matched by the clustering of dark matter halos in dissipationless simulations (\\citealt{Carlberg91, Brainerd92, Brainerd94a, Brainerd94b}; \\citealt*{Colin97}; \\citealt{Wechsler98, Colin99, Kravtsov99b, Kravtsov04,Neyrinck04}) the size of observational samples did not allow thorough tests of the galaxy-halo relation. For example, \\citet{Kravtsov04} and \\citet{Neyrinck04} compare the two-point correlation function of halos and bright galaxies from the SDSS and PSCz surveys, respectively, and find good agreement on scales $0.1\\lesssim r/$($h^{-1}$~Mpc$) \\lesssim 10$.  However, these studies do not attempt to match the clustering of fainter galaxies. \\citet{Tasitsiomi04} assign luminosities to halos and compute the galaxy-mass correlation function, finding good agreement for two broad luminosity bins, after a reasonable amount of scatter was introduced into the relation.  The current work extends these analyses by comparing results of very high-resolution simulations to the most current measurements of the two-point correlation functions over a wide range of luminosities and redshifts. The large size and wide luminosity range of observational samples allows us to test the relation between galaxy luminosities and properties of their host halos with unprecedented power.  The novel feature of the simulation analysis we present is that for each halo and subhalo we track the evolution of its properties, such as mass and maximum circular velocities. As we show below, this is a key piece of information for reasons that are easy to understand. For distinct halos, the current circular velocity is a measure of their potential well built-up during evolution, and can therefore be expected to be tightly correlated with the stellar mass (or more generally the baryonic mass) of the galaxy the halo hosts. The circular velocity of subhalos in dissipationless simulations, on the other hand, is a product of both mass buildup during the period when the halo evolved in isolation {\\it and} tidal mass loss after the halo starts to orbit within the virialized region of a larger object and experience strong tidal forces \\citep{Hayashi03,Kravtsov04b,Kazantzidis04b}. The stellar component of the galaxies, which should be more tightly bound than halo dark matter, should not be significantly affected by tidal forces and can stabilize the mass distribution (and hence {\\Vmax}) in the inner regions. We can therefore expect that luminosity and stellar mass of galaxies hosted by halos in dissipationless simulations should be correlated with the subhalo mass or circular velocity, \\Vin, {\\it at the epoch of accretion}, rather than with its current value. This is borne out by cosmological simulations, which include gas dynamics, cooling, and star formation \\citep{Nagai05} who show that selection using {\\Vin} results in subhalo properties similar to the selection based on stellar mass of galaxies subhalos host.  One can therefore argue that a reasonable approach is to relate galaxy luminosity to the current halo circular velocity for distinct halos and to the circular velocity at accretion for subhalos. The main result of this study is that this simple model reproduces the luminosity-dependence of galaxy clustering at different epochs with remarkable, and perhaps surprising, accuracy.  We note that \\citet{Vale05} have recently presented and used a semi-analytic model for subhalos, which employs a similar approach to luminosity assignment, except that they use the total bound halo mass instead of circular velocity.  The paper is organized as follows.  In $\\S$\\ref{s:sims} we briefly describe the simulations, halo finding algorithm, and the method for tracking the evolution of halos.  $\\S$\\ref{s:appr} details our method for relating halos to galaxies, and elaborates upon our motivation for using \\Vin as the basis for the luminosity assignment for subhalos.  In $\\S$\\ref{s:res} we compare observational clustering results over the redshift interval $0<z<5$ to the clustering of halos in dissipationless $\\Lambda$CDM simulations.  The halo occupation distribution implied by this model is described in $\\S$\\ref{s:hod}. In $\\S$\\ref{s:disc} we discuss the implications of our results and in $\\S$\\ref{s:conc} we summarize our main conclusions.  Throughout this paper we assume a $\\Lambda$CDM cosmology with $(\\Omega_m,\\Omega_{\\Lambda},h,\\sigma_8) = (0.3,0.7,0.7,0.9)$. ",
        "conclusions": "\\end{deluxetable}    \\label{s:disc}  Following the reliable identification of subhalos in $N$-body simulations, there has been a concerted effort to identify such subhalos with observed galaxies.  A persistent problem with this approach is that subhalos selected using the present-day mass or circular velocity tend to have radial distributions within their larger distinct halos that are shallower than those exhibited by observed galaxies \\citep{Gao04a,Diemand04,Nagai05,Maccio05}.  However, \\citet{Nagai05} find that the observed radial profile of galaxies can be reproduced by selecting subhalos using stellar mass, rather than the present day total mass.  We argue that the best proxy for the stellar mass in dissipationless simulations is the current maximum circular velocity for distinct halos and the circular velocity at the time of accretion, \\Vin, for subhalos.  We have shown that this simple, non-parametric model which relates galaxy luminosities to the maximum circular velocity of dark matter halos, \\Vmax, by preserving the observed galaxy luminosity function, accurately reproduces the observed luminosity-dependent clustering of galaxies over projected separations $0.1<r_p/(h^{-1}$ Mpc$)<10.0$, from $z\\sim 5$ to $z\\sim 0$. The only assumption of the model is that there is a monotonic relation between galaxy luminosity and halo circular velocity. To assign luminosities we use the maximum circular velocity at the time of accretion, \\Vin, for subhalos, and the maximum circular velocity at the time of observation, \\Vnow, for distinct halos.  This ingredient is crucial for accurately reproducing the observed clustering of galaxies at $z\\sim 0$.  The success of the model has several important implications.  First, our results indicate that dissipationless simulations do not suffer from significant overmerging because they reproduce both the amplitude and shape of the correlation function and the underlying HOD quite accurately. If our dissipationless simulations were missing a significant fraction ($\\gtrsim 40\\%$) of objects in groups and clusters, as suggested by some recent studies \\citep{Gao04a,Diemand04}, both the amplitude and shape of the correlation function in these simulations would be grossly incorrect.  We also note that recent analysis of the observed galaxy-mass cross correlation function disfavors the existence of a large population of satellites without associated dark matter halos \\citep{Mandelbaum05}.  In order to assess the quantitative impact of a fraction of ``orphan'' galaxies --- satellite galaxies with no identifiable subhalo \\citep{Gao04a} --- in our model, we perform the following test.  We artificially increase the fraction of subhalos in such a way that the subhalo occupation function, $\\langle N(M)\\rangle_{sat}$ increases in amplitude while maintaining the same shape.  In order to match the observed galaxy luminosity function, we simulatenously increase the overall \\Vmax threshold for a given sample, such that the number density does not change.  We find that an orphan fraction $>10$\\% is inconsistent with the correlation functions of the SDSS galaxy samples at $z\\sim0$ for all luminosities (for this cosmological model and the assumption of no scatter between $L$ and \\Vin).  Our results also imply that the central assumption of the luminosity assignment model --- the tight, monotonic relation between galaxy luminosity and halo circular velocity --- likely exists for the observed galaxies.  Such a relation is indeed expected \\citep[e.g.,][]{Mo98} for isolated galaxies, but our results indicate that this is true globally for galaxies of different types and for a wide range of redshifts.  We argue that for subhalos the dissipation should result in a centrally condensed, tightly-bound stellar system which would stabilize the halo circular velocity against tidal heating. Stellar mass and luminosity of galaxies should therefore correlate with the circular velocity of subhalos at the time they are accreted, before significant tidal evolution takes place. Given that we match luminosity at a particular epoch to the circular velocity at different epochs (the epoch of observation for distinct halos, and the epoch of accretion for subhalos), a subtle implication of our scheme is that the relation between luminosity and \\Vmax does not evolve strongly with time --- a result which may have been anticipated by the lack of scatter in the Tully--Fisher relationship in different environments.  The corollary is then that the clustering of a particular galaxy sample is largely determined by the clustering of halos and subhalos that host the galaxies. Clustering of halos and subhalos is governed by gravitational dynamics \\citep[e.g.,][and discussion therein]{Kravtsov99b,Zentner05}, while the particular subset of halos that host galaxies in a given sample is determined by the relations between observable galaxy properties and properties of their host dark matter halos {\\it and} selection criteria used to define the sample. In the case of the galaxy luminosity, the relation with halo circular velocity appears to be particularly tight.  The model agreement with the clustering properties of the LBG population at $3<z<5$ is perhaps more surprising, given that the LBG selection criteria are significantly more complicated than those of the luminosity-selected samples at $z\\leq 1$.  The success of our $L-$\\Vmax model indicates that the rest-frame LBG luminosity is likely to be tightly related to the halo circular velocity (and hence total mass). This would, in turn, suggest that LBGs are fair tracers of the overall halo population rather than a special subset of halos, such as halos undergoing mergers or collisions.  Although we have not explicitly investigated the latter scenarios (i.e., mergers or collisions), it seems unlikely that the observed LBG clustering would be well matched by a model in which the effective duty cycle for the LBGs is small and most LBGs are associated with minor mergers or small-mass collisions.  We note, however, that differences in the clustering properties and inferred host masses for the so-called ``massive halo'' model and the ``collisional starburst'' model may have been exaggerated --- in the former case because the massive subhalos were ignored, and in the latter because the efficiency of star formation in mergers was probably overestimated \\citep[e.g.,][]{Cox05}.  A key success of the simulations and our model is the correct description of the small-scale upturn in the correlation function.  At $z\\sim 0$, the upturn is detected for bright galaxies \\citep{Zehavi04}, while the correlation function of faint objects is very close to a power law \\citep{Zehavi05}. At $z\\gtrsim 1$, the upturn is more pronounced and is now unambiguously detected both in the DEEP2 data and in LBG samples at $z\\sim 3-5$, as predicted from the halo model arguments \\citep{Zheng04} and cosmological simulations \\citep{Kravtsov04}. The trend towards more pronounced deviation from the power-law form of CF at higher $z$ appears to be due to a couple different factors.  The accretion rate of subhalos is substantially higher at high $z$, so that the abundance of subhalos with masses close to the threshold mass of the sample, $M_{\\rm min}$, is also larger on average (see \\S~\\ref{s:hod}).  This makes the ``shoulder'' of the HOD of high-$z$ objects shorter and not as flat, increasing the number of small-separation pairs between central galaxies and satellites and between satellites themselves in {\\it small-mass halos}. This also makes the width of the HOD  at masses close to $M_{\\rm min}$ wider and closer to the Poisson  distribution. Although the overall fraction of subhalos in a given sample actually {\\it decreases} with increasing redshift (see Fig.~\\ref{f:cumn_z}), the average mass of the distinct halos hosting most of subhalos is decreasing as well. The contribution of the pairs of objects within the same halo to the 1-halo term of the correlation function comes then predominantly from the smallest halos in the sample and the 1-halo term is considerably ``peakier'' than at lower redshifts, where halos of a wide range of masses contribute.  We argue therefore that the more pronounced upturn in the correlation function at higher redshifts and for rarer objects at a given epoch is primarily due to the increased fraction of subhalos in host halos with masses closer to the threshold mass of the sample.  In the present work, we have not included any scatter in matching luminosity and velocity relations, and it is interesting that making this approximation works so well.  Some scatter in the $L-${\\Vmax} relation should certainly be expected given the scatter in the observed Tully-Fisher and Faber-Jackson relations.  This scatter likely depends on band, galaxy properties, and the extent to which internal dust extinction has been corrected for.  For our purposes, the main effect of adding scatter in this relationship to our model would be at the highest luminosities.  The insensitivity at low luminosities arises because 1) the luminosity function of galaxies is rather flat at $L<L_{\\ast}$ so that the scatter results in almost equal numbers of objects scattering in and out of the sample, and 2) the bias of halos corresponding to these luminosity cuts ($L<L_{\\ast}$), is only a weak function of mass and luminosity both in our simulations and in observations \\citep{Tegmark04}, so changes in the mass distribution of galaxies do not affect clustering significantly.  In contrast, at higher luminosities, both the luminosity function and the bias dependence on mass are steep, and adding scatter serves to decrease the average clustering and makes the luminosity a steeper function of mass. \\citet{Tasitsiomi04} find that some scatter is required to match the galaxy-mass correlations, but the conclusion is reached using \\Vnow for subhalos, which we expect to exhibit more scatter with luminosity than \\Vin. There is, however, an interesting tension between constraints provided by the observed galaxy-galaxy and galaxy-mass correlation functions for the most massive halos (cf. Fig.~\\ref{f:xigm}): the former constrains the scatter for bright galaxies to be small, while the latter appears to require significant amount of scatter.  We postpone a full discussion of these issues including combined constraints to future work.  \\vspace{1em}"
    },
    "0512/astro-ph0512144_arXiv.txt": {
        "abstract": "The gap caused by a strong electric field between the quark surface and nuclear crust of a strange star is studied in an improved model including gravity and pressure as well as electrostatic forces. The transition from gap to crust is followed in detail. The properties of the gap are investigated for a wide range of parameters assuming both color-flavor locked and non color-flavor locked strange star cores. The maximally allowed crust density is generally lower than that of neutron drip. Finite temperature is shown to increase the gap width, but the effect is significant only at extreme temperatures. Analytical approximations are derived and shown to provide useful fits to the numerical results. ",
        "introduction": "Strange stars are stars made of absolutely stable quark matter \\cite{Baym:1976yu,Witten:1984rs,Haensel:1986qb,Alcock:1986hz}.  Their existence (i.e., the stability of strange quark matter) depends on poorly constrained strong interaction properties, and remains to be decided by observation or experiment (see \\cite{Madsen:1998uh,Weber:2004kj} for reviews). If strange stars are stable, they contain roughly equal numbers of up, down, and strange quarks, but due to the higher mass of the s-quark, they will normally contain a slight deficit of strange quarks with negative charge, and thus have an overall positive quark charge to be compensated by electrons. Even color-flavor locked quark matter, which is electrically neutral in bulk \\cite{Rajagopal:2000ff}, has an overall positive quark charge due to surface effects \\cite{Madsen:2000kb,Madsen:2001fu,Usov:2004iz}.  The quark surface of a strange star is very sharp (the density drops from above nuclear matter density to zero within a few fm), and since the electrostatic force is weaker than the strong force, some of the electrons necessary to create an overall charge neutral object will form a thin atmosphere with a huge electric field (up to $10^{18}$ V/cm) outside the quark phase. This field is capable of sustaining a nuclear matter crust decoupled from the quark phase by an electron-filled gap.  The maximum mass of such a crust is approximately $2\\times 10^{-5} M_\\odot$, corresponding to the situation where the inner boundary of the crust reaches the neutron drip density, $4\\times 10^{11}\\; {\\rm g\\; cm}^{-3}$, where neutrons drip out of crust nuclei and get dissolved in the quark phase. Smaller crust mass limits would occur if the gap is sufficiently narrow to allow direct contact with the crust or if the rate of quantum tunneling of nuclei into the quark core of the star is large. The crust mass and moment of inertia as well as the coupling between core and crust play important roles in the understanding of strange star properties, and therefore the properties of the gap are very important.  The width of this gap is the main topic of the present investigation.  The properties of the electron atmosphere and the impact on a nuclear matter crust have been studied by many authors, including \\cite{Alcock:1986hz,Kettner:1994zs,Huang:1997,Xu:1999up,Cheng:2003hv, Usov:2004iz,Usov:2004kj,Jaikumar:2005,Zdunik:2002}.  These studies have involved solutions to the Poisson equation for the electrostatic potential with the boundary condition of electric neutrality deep within the strange star, as well as a condition for the potential at large distances. For bare strange stars, i.e., stars without a nuclear crust, the potential goes to zero at infinity. For strange stars accreting nuclear matter on the surface, the condition has been taken to be a matching of the electric potential from the electrons to the value of the potential required in the bulk of the nuclear matter crust (typically several MeV, depending on the crust density). This leads to a typical gap size (distance from quark surface to the inner surface of the nuclear crust) of $10^{2}-10^{3}$ fm, comparable to the distance over which the potential drops by a factor of a few.  The approaches just described do not account however for the detailed balance between electrical and gravitational forces and pressure in the transition from gap to crust. The potential and its derivative are taken to be continuous across the crust boundary while gravitational forces and pressure in the crust are assumed insignificant (though see \\cite{Huang:1997} for an approximate inclusion of gravity). In this work we expand earlier treatments, investigate the effects of including the Newtonian gravitational field, and find the detailed structure of the gap and the transition to the crust. This leads to more narrow gaps than previously found and thus constrains the density at the base of the crust below the neutron drip density. As the temperature increases, the additional thermal reservoir of electrons initially widens the gap then narrows it, but even for temperatures as high as $10^7$ K in the gap region this effect is small and the width remains constant up to temperatures in the range $10^8-10^{10}$ K depending on the crust density. ",
        "conclusions": "We have expanded the treatment of gaps below strange star crusts to include effects of pressure and gravity on the crust nuclei and found the structure of the transition from gap to crust. Overall our results are as one would expect with a very sharp transition from gap to crust, gap widths ranging from a few fm at densities near neutron drip over a few thousand fm at densities close to the maximum to $10^{10}$ fm for very thin crusts. Crust densities are consequently limited below neutron drip unless one assumes an extreme Coulomb barrier height. If the gap width must be at least a crust lattice distance ($\\sim 200$ fm) and a reasonable value of $e\\phi_e(R_S)\\simeq 20 \\text{MeV}$ (or $e\\phi_e(z=-1000\\text{ fm})\\simeq 20\\text{ MeV}$) is assumed then the crust density is limited to a few times $10^{10} \\text{ g/cm}^3$.  The variation of the gap width with temperature is noteworthy in that it initially increases with temperature which is a qualitatively new feature. However this increase only takes place at temperatures below 1 MeV ($10^{10}$ K) for crust densities which give very small gap widths of the order of 1 fm such as in the case of the HZ2003 crust at maximum density in Figs.~\\ref{artikeltemp} and \\ref{artikeltrans}. A realistic crust should have a sufficient gap width to be stable against strong interaction with the core, and examples of such crusts could be the HZ1990 and BPS crusts in Figs.~\\ref{artikeltemp} and \\ref{artikeltrans} -- the HZ1990 crust is almost stable, and lowering the density a little to make it so would not change the dependence on temperature much. If we apply the criterion in Eq.~\\eqref{stabtau} for crust stability we can find the highest stable density at given temperature and Coulomb barrier height, $e\\phi_e(R_S)$. Numerical results for this are shown in Fig.~\\ref{stabrho} and we see that temperature effects are then not important below $10^{10}$ K. If we demand at least a crust lattice distance in non accreting systems the crust becomes even more stable against variations in the temperature.  \\begin{figure}[!htbp] \\centering \\includegraphics[width=8.5cm]{fig14.ps} \\caption{Highest stable density according to Eq.~\\eqref{stabtau} at given temperature and $e\\phi_e(R_S)= 1,5,10,20,30$~MeV from below. The crust composition is from HZ2003 and we stop at the temperature corresponding to neutron drip density.} \\label{stabrho} \\end{figure}  The thermal structure of accreting strange stars was investigated by Miralda-Escud\\'e et al.\\ \\cite{Miralda:1990} who found temperatures around $10^7$ K at the crust boundary for strange stars accreting $10^{-11}$ to $10^{-9.83} \\text{ M}_\\odot \\text{ yr}^{-1}$ assuming that 10 MeV is released per nucleon converted to strange quark matter (bag constant $B=103 \\text{ MeV/fm}^3$). In this case the gap width would not be affected at all, and one might as well use the zero temperature results. Superbursters however are known to have much higher accretion rates around 0.1-0.3 times the Eddington accretion rate, $M_{\\text{Edd}}\\sim 10^{-8}\\text{ M}_\\odot\\text{ yr}^{-1}$, and in a recent investigation Page and Cumming \\cite{Page:2005} found temperatures for such stars in the range $10^8-10^9$ K at the crust boundary. This work was motivated by the difficulties in achieving carbon ignition at observed column depths in superburst models on neutron stars -- see Cumming et al. 2005 \\cite{Cumming:2005} for a recent review -- and it was found that superbursts may ignite at the right column density on strange stars provided that neutrino emission in the core is slower than direct Urca which implies that strange quark matter should be a color superconductor. During the superburst itself temperatures in the range $2-7 \\times 10^9$ K are reached for a few hours \\cite{Cumming:2001, Schatz:2003}. This is again too low for significant temperature effects and for stable crusts one may as well use the zero temperature expressions.  The present treatment could be improved by including general relativity instead of Newtonian gravity, and by using a less simplistic model for the equation of state for the crust. Such improvements are necessary to calculate the properties of the whole star, but we expect the changes in the gap properties to be minor compared to the results presented here. The same can be said about a more detailed treatment of the quark phase, which has here merely been treated as a source of a given total mass and starting value for the electron chemical potential. However, as pointed out in \\cite{Jaikumar:2005}, the electric potential in the outer layers of the strange core could be qualitatively different from the behavior assumed here if a mixed phase of quark nuggets and electrons can be realized.  Another important oversimplification involves the assumption of a constant pairing gap all the way to the quark core surface in the treatment of quark matter in the CFL-phase. In fact the pairing gap $\\Delta_{\\text{CFL}}$ may be envisaged to vanish within a surface distance of order $1/\\Delta_{\\text{CFL}}\\approx {\\text{2--20~fm}}$ for gap energies between 100 and 10~MeV. Possibly a range of phases appear in the extreme uppermost layers, including the possibility that the very surface resembles the ungapped ordinary quark phase also discussed above. A detailed treatment of gapped phases near a surface remains to be performed.  In conclusion the results of our expanded treatment, while quantitatively somewhat different, are qualitatively consistent with previous work in the literature with regards to gap width, transmission coefficient and possible crust densities. The increase in gap width with temperature is new but significant only at very high temperatures not likely realized. Perhaps the most useful results of the investigation are the various analytical approximations which have been derived and shown to fit the full numerical solutions very well, since these provide better physical insight and can be easily used in models for strange star rotation, glitches, accretion and instabilities."
    },
    "0512/astro-ph0512372_arXiv.txt": {
        "abstract": "We report the detection of broad absorptions due to \\ion{Si}{4} $\\lambda\\lambda$4088--4116 in the Luminous Blue Variable (LBV) AG Carinae during its last hot phase (2001-2003). Our NLTE spectral analysis, with the radiative transfer code CMFGEN, revealed the photospheric nature of these lines predicting, however, much narrower and deeper absorption profiles than observed. Using a recently-developed code to compute synthetic spectra in 2D geometry allowing for the effects of rotation, we could match the broad absorptions with a high projected rotational velocity of 190 $\\pm$ 30\\,km\\,s$^{-1}$ on 2001 April. Analysis of spectra obtained on 2002 March and 2003 January, when the star was cooling, yielded to a projected rotational velocity of 110\\,$\\pm$\\,10\\,km\\,s$^{-1}$ and 85\\,$\\pm$\\,10\\,km\\,s$^{-1}$, respectively. The derived rotational velocities are proportional to $R_{\\star}^{-1}$, as expected from angular momentum conservation. We discuss the effects of such high rotation on the spectral analysis of AG Car, and on the wind terminal velocity. Our results show direct spectroscopic evidence, for the first time, that a LBV may rotate at a significant fraction of its break-up velocity. Thus, AG Car (and possibly other LBVs) is indeed close to the $\\Gamma\\Omega$ limit, as predicted by theoretical studies of LBVs. ",
        "introduction": "Introduction}  Massive stars have a tremendous impact on both the chemical and dynamical evolution of a galaxy. They input energy and momentum, and eject their nucleosynthesis products into the interstellar medium \\citep{fhy03}. A detailed knowledge of massive star evolution is thus crucial to understanding their host galaxies, especially in the early Universe when massive stars are dominant \\citep{mc05}.  Over the last decade, significant efforts have been made to include the effects of rotation consistently in the evolutionary models for massive stars \\citep{mm_rot_paper1}. It has been found that the effects of rotation are far from negligible, especially for the most massive stars with high initial rotational velocity. As a consequence of rotation, the evolutionary tracks on the HR diagram, and the duration of the evolutionary stages (main sequence, Luminous Blue Variable -- hereafter LBV, and Wolf-Rayet phases), are substantially modified. For a recent review we refer to \\citet{maeder_araa00}.  Rotation is also expected to play an important role in both LBV variability and LBV eruptions \\citep{langer99,mm_omega00,alm04}. Even before reaching the Eddington limit, these stars may lose large amounts of material if they have a significant rotational velocity compared to their critical rotational velocity. Stellar rotation, during the LBV phase, can have an impact on the surrounding circumstellar medium, shaping their ejected nebulae with the characteristic bipolar geometry found around many of these stars \\citep{hd94,md01,do02}.  In this Letter we present high-resolution, high S/N optical spectra of the LBV AG Carinae. AG Car has pronounced S Dor type variations, changing its effective temperature from 9,000 to 25,000 K \\citep{leitherer94,stahl01}. AG Car is considered the prototype of the LBV class and has been extensively observed, from UV to radio wavelengths, with photometric, spectroscopic and polarimetric techniques over several decades. It also has a bipolar nebula which is seen at various wavelengths \\citep{nota92,voors00,dw02}  that was ejected approximately 10$^4$ years ago \\citep{lamers01}.  Here we report the detection of broad absorptions associated with the high-excitation lines of \\ion{Si}{4} $\\lambda\\lambda$4088--4116 in AG Car during its last hot phase (2001-2003). These data are described in Sec. \\ref{obs}, while in Sec. \\ref{modres} we present the results of our NLTE spectral modeling. Our modeling reveals the photospheric nature of the \\ion{Si}{4} lines and allow us to derive, for the first time, the rotational velocity ($v_{rot}$) of AG Car and how it varies during the S Dor cycle. We show that $v_{rot}$ is high compared to its critical rotational velocity ($v_{crit}$). In Sec. \\ref{disc} we discuss some of the effects expected to be present in fast-rotating stars, such as the reduction in the wind terminal velocity and the presence of asymmetric winds. The high $v_{rot}$ derived has important implications not only for AG Car and its wind, but for the general understanding of the LBV phase and the subsequent evolutionary stages of a massive star. ",
        "conclusions": ""
    },
    "0512/astro-ph0512658_arXiv.txt": {
        "abstract": "Galaxy colours are determined for two samples of 6C and 3CR radio sources at $z \\sim 1$, differing by a factor of $\\sim 6$ in radio power.  Corrections are made for emission line contamination and the presence of any nuclear point source, and the data analysed as a function of both redshift and the radio source properties.  The galaxy colours are remarkably similar for the two populations, and the UV excess evolves with radio source size similarly in both samples, depsite the fact that the alignment effect is more extensive for the more powerful 3CR radio galaxies. These results seem to suggest that the alignment effect at these redshifts does not scale strongly with radio power, and is instead more closely dependent on galaxy mass (which is statistically comparable for the two samples).   However, it is likely that the presence of relatively young ($\\lta$ several $10^8$ years old) stellar populations has considerably contaminated the $K-$band flux of these systems, particularly in the case of the more powerful 3CR sources, which are $\\sim 0.5$mag more luminous than the predictions of passive evolution models at $z \\sim 1$.  The higher luminosity of the 3CR alignment effect is balanced by emission at longer wavelengths, thereby leading to comparable colours for the two samples. ",
        "introduction": "\\footnotetext{E-mail: k.inskip@shef.ac.uk} The properties of powerful radio sources vary considerably with both radio power and redshift.  Up to redshifts of $z \\sim 1$, radio sources are usually hosted by massive elliptical galaxies, with scale sizes of $\\sim 10-15$kpc (e.g. Inskip et al 2005; McLure et al 2004). Whilst the average galaxy size does not appear to vary greatly, the host galaxies of the higher redshift radio sources within this redshift range ($z \\gta 0.6$) are intrinsically more luminous than those at lower redshifts (Inskip et al 2002a, 2005), to an extent which implys that factors other than simple passive evolution are involved.  At larger redshifts ($z > 2$), the host galaxies of the most distant radio sources are clearly still in the process of formation (e.g. van Breugel et al 1998, Pentericci et al 2001).  Redshift evolution is also observed in the emission surrounding the host galaxies.   At $z > 0.3$, the host galaxies are often seen to be surrounded by considerable excess rest-frame UV emission; at higher redshifts ($z \\gta 0.6$) this emission is usually more extensive, is generally observed to be closely aligned with the radio source axis (Chambers, Miley \\& van Breugel 1987; McCarthy et al. 1987; Allen et al 2002), and is known as the {\\it Alignment Effect}.  Both the alignment effect and the properties of the extended emission line regions surrounding these sources are seen to be more extreme (in terms of luminosity, alignment with the radio axis, physical extent and gas kinematics) both for more powerful radio sources, and also for the smaller radio sources (Best et al 1996, 2000b; Inskip et al 2002c, 2003, 2005). Several different mechanisms are thought to be responsible for producing these regions of extensive aligned emission.  These include: extended line emission and nebular continuum radiation (Dickson et al 1995), scattering of the UV continuum from the AGN (e.g. Tadhunter et al 1992; Cimatti et al 1993) and young stars (e.g. Chambers \\& McCarthy 1990) potentially produced in a radio jet induced starburst (McCarthy et al 1987).  Merger-induced starbursts may also be responsible for the presence of a relatively young stellar population in/around the host galaxy (Tadhunter et al 2005 and references therein).  McCarthy, Spinrad \\& van Breugel (1995) found that extended line emission can be observed around the majority of 3CR galaxies at $z \\gta 0.3$. The relative contribution of emission lines to the total aligned emission varies from source to source, although both the alignment effect and emission line flux are more extreme for smaller radio sources.  Typically, line emission provides from 2\\% up to 30\\% of the total rest-frame UV aligned emission for 3CR galaxies at $z \\sim 1$ (although this does depend on the distribution of emission lines within the wavelength range of the observed filter), with similar proportions found for the less powerful 6C radio sources at the same redshifts (Best 1996; Inskip et al 2003). In addition to line emission, nebular continuum radiation is also produced due to other radiative processes associated with the ionized gas. Although a significant process for some sources, the total flux provided by line emission and nebular continuum emission alone cannot account for all of the excess emission forming the alignment effect (e.g. Inskip et al 2003, Tadhunter et al 2002).  Although the active nucleus of a radio galaxy may be obscured from view, the powerful UV continuum emitted by the AGN may be scattered towards an observer by dust or electrons in the extended structures surrounding the galaxy. Emission polarised perpendicularly to the direction of emission from the AGN due to scattering has been observed from these extended regions, and is consistent with the orientation--based unification scheme for radio galaxies and quasars (e.g. Cimatti it et al 1993, Tran et al 1998). However, while polarized emission is frequently observed, the emission from the aligned structures surrounding many radio galaxies lacks the high levels of polarisation expected if scattering of the UV emission from the obscured quasar nucleus were the only mechanism occurring. A recent study of $0.15 < z < 0.7$ radio galaxies by Tadhunter et al (2002) found that scattering contributed a significant proportion of the UV excess in many cases, but was very rarely the dominant factor.  Further to this, some 3CR sources do not exhibit any polarization of their extended structures (e.g. Wills et al 2002).  Finally, one of the first explanations proposed for the excess UV continuum was that it was due to emission from young stars, whose formation was triggered by the passage of the expanding radio source (e.g. McCarthy et al 1987, Chambers, Miley \\& van Breugel 1987). The emission from such a population of hot, young stars would dominate the UV emission of the galaxy, and produce a fairly flat spectral shape, whilst at near infrared wavelengths, the emission would still be predominantly due to the old stellar population of the host galaxy. Any young stellar population would quickly (within $\\lta 10^7$years) evolve, accounting for the rapid evolution with radio size seen in the UV aligned structures at $z \\sim 1$ (Best, Longair \\& R\\\"{o}ttgering 1996). In order to account for the excess emission, the mass of stars formed in the interactions with the radio source is typically required to be only a few $10^8 M_\\odot$ (e.g. Best, Longair \\& R\\\"{o}ttgering 1997a). It is questionable how easily jet-induced star formation can occur.  Numerical simulations in the literature often disagree on whether clouds will be compressed or shredded/dissipated (e.g. Rees 1989; Begelman \\& Cioffi 1989; Klein, McKee \\& Colella 1994; Icke 1999; Poludnenko, Frank \\& Blackman 2002), although recent work including the effects of cooling (e.g. Mellema, Kurk \\& R\\\"{o}ttgering 2002) suggests that such triggered star formation is indeed plausible. Observational evidence for star formation triggered by the radio source jets is seen in isolated objects: 3C 34 (Best, Longair \\& R\\\"{o}ttgering 1997a), Minkowski's object (van Breugel et al 1985), 3C285 (van Breugel \\& Dey 1993).  Evidence that shocks strongly influence the ionization and kinematics of the emission line gas has been observed in the spectra of many distant radio sources (e.g. Best, R\\\"{o}ttgering \\& Longair 2000b; Sol\\'{o}rzano-I\\~{n}arrea, Tadhunter \\& Axon 2001; Inskip et al 2002b), particularly in the case of smaller radio sources, i.e. those with a projected physical size of $< 120$kpc.  We also find that the sources in which shocks have the greatest impact on the emission line gas properties  are those with the most extensive, luminous alignment effects (Inskip et al 2005, 2002c). The shocks associated with an expanding radio source can greatly influence the alignment effect.  Ionizing photons associated with the shocks may boost certain emission lines, and also lead to an increase in nebular continuum emission.  Star formation induced by the passage of radio source shocks through the cool dense gas clouds is also an obvious mechanism by which the alignment effect may be enhanced.  In addition, the passage of a fast shock can potentially cause the break-up of optically thick clouds (Bremer, Fabian \\& Crawford 1997), increasing the covering factor for scattering of the UV flux from the AGN.  The more numerous, smaller clouds will also have a larger cross section for ionization by the AGN, leading to an increase in the total flux of line emission.  Despite the compelling evidence for each of the alignment effect mechanisms outlined above, their relative balance is still poorly understood. A wide range of galaxy colours provides a useful means of probing the contributions from different physical processes, and their dependence on the properties of the radio source population (power, size, epoch). This is of paricular interest, since the dependence of different mechanisms on each of these parameters varies significantly. For example, line emission is known to be closely linked to both radio source power and size. Star formation, on the other hand, may be independent of radio power, despite evolving quickly with age.  It is necessary that these processes are better understood before we can interpret the clear redshift evolution of the alignment effect.  One problem with studies of 3CR radio galaxies is that the radio power of sources in a flux-limited sample such as this increases with redshift, leading to a degeneracy between redshift and radio power. The less powerful 6C sample provides a population of radio galaxies ideally suited for breaking this degeneracy.  The factor of $\\sim 6$ difference in radio power between the samples is small compared to the wide range of powers (spanning several orders of magnitude) observed for the radio source population as a whole.  However, it is comparable to the difference in power between 3C sources at low ($z \\sim 0.1-0.5$) and high ($z \\sim 1$) redshifts, and it is the evolution within this range that we hope to explain.   To this end, we have carried out a program of multiwavlength imaging and spectroscopic observations of a subsample of 11 6C radio sources at $z \\sim 1$ (Inskip et al 2003; Best et al 1999; Inskip et al 2002b), which are well matched to the 3CR subsample previously studied by Best et al (1997b, 2000a). Having already analysed the spectroscopic (Inskip et al 2002c) and morphological properties of these systems (Inskip et al 2005, which included an analysis of the variations in host galaxy size), we now turn our attention to investigating the effect of radio power on the galaxy colours, and the nature of the excess UV emission (including the relative contributions of line emission and nebular continuum). The structure of the paper is as follows.  In section 2, we briefly outline the sample selection and observations. Colours are determined for the two matched samples, including emission line corrections.  The results are presented in section 3, and analysed in more depth in section 4, where we consider the influence of radio source size and power.  We consider the influence of the host galaxy stellar populations in section 5, and present our conclusions in section 6. Values for the cosmological parameters $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$ and $H_{0}=65\\,\\rm{km\\,s^{-1}\\,Mpc^{-1}}$ are assumed throughout this paper.   \\begin{table*} \\caption{Observed (roman) and calculated (italics) magnitudes for the 6C and 3CR sources at $z \\sim 1$, together with a summary of the source redshifts, radio sizes and radio power at 178MHz.  All magnitudes for the sources in both samples were determined within a 4\\arcsec\\ diameter aperture and have been corrected for galactic extinction. The 6C data was previously presented in Inskip et al (2003).  The 3CR HST, $J$ and $K$ band data were initially analysed in 5\\arcsec\\ and 9\\arcsec apertures (Best et al 1997), but magnitudes have been re-extracted in 4\\arcsec\\ diameter apertures for the purposes of this paper.  The $H$ band 3CR data were obtained via the UKIRT service program. For 6C1256+36, the flux due to the unresolved companion object has been modelled (Inskip et al 2005, Paper 2) and removed from the F702W and K-band data (magnitudes which remain contaminated by this object are marked by a '*').  No corrections have been made for flux contamination in the HST data by adjacent objects, as it is usually impossible to disentangle adjacent galaxies from any aligned line/continuum emission. } \\begin{center} \\begin{tabular} {ccccr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}l}\\\\ Source    & Redshift & L$_{178}$ (log$_{10}$W\\,Hz$^{-1}$)&D$_{rad}$ (kpc) &\\multicolumn{2}{c}{F702W} & \\multicolumn{2}{c}{F814W} & \\multicolumn{2}{c}{$J$}& \\multicolumn{2}{c}{$H$}& \\multicolumn{2}{c}{$K$}\\\\\\hline 6C0825+34 & 1.467 &28.31 &64  & {\\it 22.59}& {\\it 0.31}&22.10&0.17 &19.86 &0.17 &19.68 &0.29 & 19.12& 0.12 \\\\ 6C0943+39 & 1.035 &28.07 &92  & 22.08&0.14 & {\\it 21.55}& {\\it 0.20}&19.70&0.12&19.27&0.14&18.09&0.07\\\\ 6C1011+36 & 1.042 &28.02 &444 & 21.73&0.07 & {\\it 21.23}&{\\it 0.15} &19.68&0.20&18.63&0.10&17.83&0.06\\\\ 6C1017+37 & 1.053 &28.10 &65  & 21.77&0.06 & {\\it 21.10}&{\\it 0.14} &19.89&0.16&19.54&0.15&18.57&0.09\\\\ 6C1019+39 & 0.922 &28.03 &67  & {\\it 21.04}&{\\it 0.29} & 20.07&0.04 &18.47&0.07&17.71&0.06&16.80&0.04\\\\ 6C1100+35 & 1.440 &28.32 &119 & {\\it 22.31}& {\\it 0.26}& 21.80&0.06 &19.59&0.12&19.02&0.11&17.99&0.07\\\\ 6C1129+37 & 1.060 &28.03 &141 & 22.02&0.07 & {\\it 21.64}&{\\it 0.16} &19.35&0.11&18.31&0.09&17.81&0.07\\\\ 6C1204+35 & 1.376 &28.45 &158 & {\\it 21.84}&{\\it 0.28} & 21.30&0.10 &19.31&0.11&18.76&0.11&18.01&0.07\\\\ 6C1217+36 & 1.088 &28.12 &38  & {\\it 21.91}&{\\it 0.25} & 20.89&0.05 &19.50&0.12&18.51&0.07&17.55&0.06\\\\ 6C1256+36 & 1.128 &28.20 &155 & 22.86&0.09 & {\\it 22.60}&{\\it 0.16*} &19.74&0.20*&18.58&0.09*&18.14&0.06\\\\ 6C1257+36 & 1.004 &28.06 &336 & {\\it 21.55}&{\\it 0.30} & 20.98&0.05 &19.39&0.11&18.27&0.07&17.50&0.05\\\\\\\\  3C13      & 1.351 &29.20 & 259 & {\\it 21.16}&{\\it 0.25}&20.57 & 0.03&18.74 &0.13 &\\multicolumn{2}{c}{-}&17.47 &0.11\\\\ 3C22 \t  & 0.938 &28.77 & 209 & {\\it 20.00}&{\\it 0.15}&19.29 &0.02 &17.53 &0.08 & 16.95&0.09 &15.66&0.05\\\\ 3C34 \t  & 0.690 &28.51 & 359 & {\\it 20.70}&{\\it 0.25} &{\\it 20.03} &{\\it 0.25} &18.30 &0.11 &\\multicolumn{2}{c}{-} &16.46 & 0.07\\\\ 3C41 \t  & 0.795 &28.48 & 194 & {\\it 20.42}&{\\it 0.20}&{\\it 19.79} &{\\it 0.20} &18.78 &0.14 & \\multicolumn{2}{c}{-}&15.89 & 0.05\\\\ 3C49 \t  & 0.621 &28.25 & 7   & {\\it 19.97}&{\\it 0.15} &19.33 &0.05 &\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &16.25 & 0.06\\\\ 3C65 \t  & 1.176 &29.09 & 160 & {\\it 22.08}&{\\it 0.16} &21.07 &0.03 &18.93 &0.14 &\\multicolumn{2}{c}{-} &17.19 & 0.10\\\\ 3C68.2 \t  & 1.575 &29.33 & 218 & {\\it 22.37}&{\\it 0.30} &{\\it 21.87 }& {\\it 0.30}&19.78 &0.22 & \\multicolumn{2}{c}{-}&18.18 & 0.16\\\\ 3C217 \t  & 0.897 &28.69 & 110 & {\\it 20.81}&{\\it 0.17} &20.27 & 0.02&18.81 &0.13 &18.99 &0.23 &17.88&0.13 \\\\ 3C226 \t  & 0.820 &28.75 & 259 & {\\it 20.46}&{\\it 0.18} &{\\it 19.64} &{\\it 0.16} &18.46 &0.11 &17.82 &0.13 &16.83 &0.08\\\\ 3C239 \t  & 1.781 &29.60 & 111 & {\\it 21.60}&{\\it 0.20} &21.16 &0.03 &19.01 & 0.14&\\multicolumn{2}{c}{-} &17.90 &0.13 \\\\ 3C241 \t  & 1.617 &29.39 & 8   & {\\it 22.19}&{\\it 0.20} &21.62 &0.04 &19.19 & 0.16&\\multicolumn{2}{c}{-} &17.82 &0.13 \\\\ 3C247 \t  & 0.749 &28.44 & 110 & {\\it 20.35}&{\\it 0.24} &19.46 &0.02 &\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &16.04 &0.06 \\\\ 3C252 \t  & 1.105 &28.98 & 501 & {\\it 21.10}&{\\it 0.13} &20.58 &0.03 &\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &17.54 & 0.11\\\\ 3C265 \t  & 0.811 &28.88 & 636 & {\\it 19.51}&{\\it 0.24}&{\\it 19.06} &{\\it 0.30} &17.81 &0.08 &17.52 &0.11 &16.39 &0.07\\\\ 3C266 \t  & 1.272 &29.13 & 41  & 21.22&0.03 &20.55 &0.03 &\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &17.99 &0.14\\\\ 3C267 \t  & 1.144 &29.11 & 339 & 21.50& 0.04&{\\it 20.95} &{\\it 0.10} &19.20 &0.16 &\\multicolumn{2}{c}{-} &17.46 &0.11 \\\\ 3C277.2   & 0.766 &28.62 & 422 & {\\it 20.16}&{\\it 0.20} &19.82 &0.02 &18.59 &0.12 &\\multicolumn{2}{c}{-} &17.32 & 0.10\\\\ 3C280     & 0.996 &29.13 & 118 & {\\it 20.92}&{\\it 0.06} &20.29 &0.02 &18.48 &0.11 &18.10 &0.15 &17.05 &0.09\\\\ 3C289 \t  & 0.967 &28.81 & 90  & {\\it 21.49}&{\\it 0.07}&20.50 &0.03 &18.66 &0.12 &17.97 &0.14 &17.06 &0.09\\\\ 3C324 \t  & 1.206 &29.18 & 100 & 21.49&0.03&{\\it 20.84} &{\\it 0.20} &19.18 &0.17 &18.30 &0.16 &17.33 &0.10\\\\ 3C337 \t  & 0.635 &28.32 & 326 & {\\it 20.79}&{\\it 0.25} &19.94 &0.02 &18.45 &0.11 &\\multicolumn{2}{c}{-} &16.84 &0.08 \\\\ 3C340     & 0.775 &28.48 & 363 & {\\it 20.88}&{\\it 0.18} &{\\it 19.71} &{\\it 0.21} &18.52 &0.12 &17.89 &0.12 &17.08 &0.09\\\\ 3C352 \t  & 0.806 &28.61 & 100 & {\\it 20.63}&{\\it 0.08} &20.14 & 0.02&\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &17.09 & 0.09\\\\ 3C356 \t  & 1.079 &28.96 & 638 & {\\it 21.22}&{\\it 0.09} &20.60 &0.03 &\\multicolumn{2}{c}{-}& \\multicolumn{2}{c}{-}&17.54 &0.11 \\\\ 3C368 \t  & 1.132 &29.17 & 75  & 20.47&0.02 &{\\it 19.55 }&{\\it 0.08} &19.01 &0.11 & 17.83& 0.13&17.17 &0.10\\\\ 3C437 \t  & 1.480 &29.32 & 339 & {\\it 22.38}&{\\it 0.30} &{\\it 21.86} & {\\it 0.30}&\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} &18.23 & 0.16\\\\ 3C441     & 0.708 &28.51 & 257 & {\\it 20.07}&{\\it 0.15} &{\\it 19.28} & {\\it 0.15}&18.15 &0.10 &\\multicolumn{2}{c}{-} &16.49 &0.07 \\\\ 3C470 \t  & 1.653 &29.20 & 228 & {\\it 22.79}&{\\it 0.46} &{\\it 22.35} &{\\it 0.46} &\\multicolumn{2}{c}{-}&\\multicolumn{2}{c}{-} & 18.20&0.15 \\\\  \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "The major results of our analysis of the $z \\sim 1$ galaxy colours can be summarised as follows. \\begin{enumerate} \\item[$\\bullet$] The observed infrared colours of the 6C and 3CR sources are indistinguishable, and well explained by passively evolving galaxies up to redshifts of $\\sim 1.3$. \\item[$\\bullet$] The galaxy colours begin to deviate from passive old stellar populations at higher redshifts ($z \\gta 1.3$), as increasing amounts of redshifted aligned emission lie in the wavelength range of the infrared filters. \\item[$\\bullet$] The observed F702W$-K$ and F814W$-K$ colours become increasingly blue out to redshifts of $z \\approx 1.1$, and then redden again at higher redshifts. This is linked to the fact that the excess rest-frame UV emission produced by the aligned structures becomes increasingly important between redshifts of 0.6 and 1, and that as redshift increases beyond that, the underlying old stellar population of elliptical galaxies can be expected to have increasingly red colours in these observed wavebands. \\item[$\\bullet$] The observed F702W$-K$ and F814W$-K$ colours for both galaxy samples are also statistically indistinguishable, suggesting that either the predominant alignment effect mechanisms do not scale strongly with radio power, or that some additional effect works to counterbalance this. \\item[$\\bullet$] Just as the most extreme rest-frame UV morphologies are generally associated with the smaller radio sources in the sample, these sources were also observed to display bluer colours due to the increased excess rest-frame UV emission observed on the HST/WFPC2 images. We see some signs that galaxy environment affects the strength of the observed alignment effect. \\end{enumerate} \\vspace{-3pt}  The interpretation of these results is not totally straightforward. Whilst the overall redshift evolution of the galaxy colours can be readily understood, the (lack of) variation in the observed optical-IR colours between the two $z \\sim 1$ subsamples is less easily accounted for.   The 3CR radio sources are $\\sim 6$ times more powerful than those of the 6C subsample at $z \\sim 1$, but the difference in mean $K$-band flux between the two data sets is much lower, at a factor of $\\lta 2$. Given the fact that the emission produced by many alignment effect mechanisms should scale with radio power one might expect the more powerful sources to display bluer colours. Several mechanisms for producing the excess rest-frame UV aligned emission will not necessarily produce more emission in the presence of a more luminous radio source.  Star formation triggered by the expanding radio source through the surrounding IGM could potentially be more efficient for the 6C sources.  The passage of a less powerful radio source may lead to lower levels of cloud shredding, and thus allow greater amounts of star formation.  Additionally, the AGN of the more powerful 3CR galaxies are likely to heat and ionize the gas clouds to a greater extent, impeding the star formation rate within such clouds. However, observations of the mJy radio sources LBDS 53W091 \\& LBDS 53W069 at $z \\sim 1.5$ (Dunlop 1999; Spinrad et al 1997) would seem to contradict this hypothesis.  These very low power radio galaxies have very red stellar populations consistent with ages of over 3.5Gyr.  The observations of both galaxies are consistent with a de Vaucouleurs law luminosity profile, and neither display any aligned emission. But, this may be a reflection of the radio properties of these sources (largest angular sizes are 4.2 arcsec and $<5.1$ arcsec respectively).  The disturbance to the IGM caused by such small, low power radio galaxies may not be sufficient to trigger any extra star formation at all.  It should be noted that evidence for star formation is not restricted to sources displaying a strong alignment effect, and that a large variety of stellar populations are observed in different systems.  UV HST observations of very low redshift 3CR sources (Allen et al 2002) show that a high proportion of sources display some level of recent or ongoing star formation, without any need for a large-scale alignment effect.  Very young (several Myr old) stellar populations accounting for $\\lta 1\\%$ of the total stellar mass are observed in the host galaxies of the intermediate redshift compact radio sources PKS1345+12 (Rodriguez et al, {\\it priv. comm.}) and 9C1503+4528 (Inskip et al, {\\it in prep}), whilst a wider study of 2Jy and 3CR radio galaxies (Tadhunter et al 2005, Wills et al 2002, Holt {\\it priv. comm.}) revealed evidence for older (0.05 to 2 Gyr), more massive (up to 50\\% of the total stellar mass) stellar populations.  Such varied stellar populations are likely to account for much of the scatter within different samples.  However, this does not prevent us from drawing conclusions regarding the more general effects of properties such as radio power on the samples as a whole.  One important consideration is whether star formation induced by an expanding radio source may instead depend more strongly on the mass of available gas, rather than any other parameters.  Scattering processes are also likely to depend on the available mass of gas and dust as well as on the power of the rest-frame UV emission from the AGN.  If true, the fact that the 6C and 3CR host galaxies at $z \\sim 1$ are of comparable size would be a point in favour of similar masses and hence observed colours.  However, the total mass of cool clouds in the regions of the IGM surrounding the host galaxy may or may not scale with galaxy mass.  If, for example, the radio source is triggered by a galaxy merger, the scattering processes may depend most strongly on the amount of dust and gas brought in by that merger. A further consideration is whether the 3CR sources lie in particularly rich environments, which may boost both their observed radio luminosity and also the availability of gas in the surrounding IGM.  Two main mechanisms are known/expected to vary with radio source size (and plausibly age).  Emission line flux scales strongly with radio source size as well as radio power, due to the increased importance of shock excitation in smaller, younger radio sources.  Blue galaxy colours due to jet induced star formation should also weaken at larger radio sizes, due to the rapid aging of the recently formed young stellar population. Although radio power does not appear to strongly influence the observed galaxy colours, it is noticeable that the trends observed with radio source size in the case of the more powerful 3CR data are considerably weaker (or absent) at the lower powers of the $z \\sim 1$ 6C sample, particularly after the removal of emission line contamination. The lack of any clear cut trend in these data does not rule out jet induced star formation as an important mechanism for producing the alignment effect, but rather indicates that the high level of scatter, due to other processes and the influence of the local environment/IGM, has swamped any underlying evolution of the alignment effect with source age/size. Given the comparable colours between the samples and the distinct lack of discrete blue components at any great distance from the 6C host galaxies, it is certainly clear that the bulk of the excess rest-frame UV emission lies closer to the host galaxies (or indeed within them) in the case of the 6C sources.      Finally, it seems likely that the most important factor in explaining these data is not variations in the origin or nature of the rest-frame UV excess in each sample, but the longer wavelength emission of our $K$-band data.  The $K-$band emission from the more powerful 3CR sources is increased over the levels expected from passive evolution scenarios, suggesting that the presence of a young  stellar population (or at least one more youthful than the majority old stellar population) may be skewing the galaxy colours.    The presence of either a reddened young stellar population of similar age to the radio source itself (i.e $\\gta 10^7$years), or an older population of up to a Gyr or so in age, can account for the excess $K-$band emission from the more powerful sources without leading to substantially bluer emission from the host galaxy.  However, given that the rest-frame UV excess must be balanced by any long-wavelength emission from an additional stellar population (in order that the lack of any colour difference between the samples is thereby maintained), it seems certain that reddening effects within the host galaxy are also an important factor."
    },
    "0512/hep-ph0512317_arXiv.txt": {
        "abstract": "The dwarf spheroidal galaxy Draco has long been considered likely to be one of the brightest point sources of gamma-rays generated through dark matter annihilations. Recent studies of this object have found that it remains largely intact from tidal striping, and may be more massive than previously thought.  In this article, we revisit Draco as a source of dark matter annihilation radiation, with these new observational constraints in mind. We discuss the prospects for the experiments MAGIC and GLAST to detect dark matter in Draco, as well as constraints from the observations of EGRET. We also discuss the possibility that the CACTUS experiment has already detected gamma-rays from Draco. We find that it is difficult to generate the flux reported by CACTUS without resorting to non-thermally produced WIMPs and/or a density spike in Draco's dark matter distribution due to the presence of an intermediate mass black hole. We also find that for most annihilation modes, a positive detection of Draco by CACTUS would be inconsistent with the lack of events seen by EGRET. ",
        "introduction": "It has long been thought that dark matter particles could be observed indirectly by detecting the products of their annihilations. Such products, including gamma-rays, neutrinos and anti-matter, have been searched for using a wide range of experimental techniques~\\cite{review}. Gamma-rays from dark matter annihilations, in particular, have been sought after using both satellite and ground based experiments.  The potential astrophysical source of gamma-rays from dark matter annihilations which is most often studied is the central region of our galaxy. Recently, observations by the Atmospheric Cerenkov Telescopes (ACTs) HESS \\cite{hess}, Whipple \\cite{whipple} and Cangaroo \\cite{cangaroo} have revealed the presence of a very bright gamma-ray source from this direction. The spectrum of this source has been measured in steadily increasing detail by the HESS collaboration \\cite{hess}. Although the first HESS data from this source was not inconsistent with a spectrum from annihilating dark matter \\cite{hessdark}, it is now becoming difficult to reconcile the HESS data with such a spectrum. Instead, it appears more likely that an astrophysical accelerator is responsible for this bright gamma-ray emission. As a result, future dark matter searches in this region will face a background that will be very challenging to overcome \\cite{gabi}.  Given this newly discovered background, it is important to consider other possible regions in which an observable rate of dark matter annihilation radiation may be generated. Such gamma-rays may appear as point sources external to our own galaxy, such as Andromeda (M31), M87 or the Large Magellanic Cloud~\\cite{fornengo}, or as a diffuse spectrum generated by a large number of distant sources \\cite{diffuse,closer}. Observable quantities of gamma-rays may also be generated in dark substructure within our own galactic halo. Again, this may appear as a diffuse spectrum from a large number of dark matter clumps \\cite{mwclumps}, or may be dominated by a few of the most massive dwarf galaxies within the Milky Way, such as Draco, Sagittarius and Canis Major \\cite{sarkar,tyler}.  In this article, we will discuss the prospects for detecting gamma-rays from dark matter annihilations in the dwarf galaxy Draco. We focus on this particular object for several reasons. First, of the most nearby and massive dwarf galaxies, the halo profile of Draco is the most tightly constrained by observations. Although other dwarfs may actually be brighter sources of dark matter annihilation radiation (this is likely for both Sagittarius and Canis Major \\cite{sarkar}), the rates from these objects cannot be estimated with as much confidence. Second, since dwarf galaxies are dark matter dominated, containing very few baryons, gamma-ray searches for dark matter in these regions are very unlikely to be complicated by the presence of astrophysical sources. In light of the challenges faced for dark matter searches in the galactic center, this is clearly an important consideration.  A third reason that we chose to focus on Draco is the potentially exciting results of the CACTUS gamma-ray experiment. In recent conferences \\cite{cactus}, the CACTUS collaboration has stated that they have detected an excess of $\\sim$100 GeV gamma-rays from the direction of Draco. Although still preliminary, this result, if confirmed, would have dramatic implications for dark matter.  The remainder of this article is organized as follows. In the following section, we calculate the annihilation rate of dark matter in Draco, and the resulting gamma-ray flux. We then discuss the prospects for MAGIC and GLAST to detect this flux, and then lastly turn our attention to the possible detection of Draco by CACTUS, and the implications of such an observation for dark matter. ",
        "conclusions": "In this article, we have revisited the possibility of detecting gamma-rays produced in dark matter annihilations in the dwarf galaxy Draco. Draco is the most well constrained of the Milky Way's satellite galaxies, and therefore provides the best opportunity to make reliable predictions of dark matter annihilation rates and corresponding gamma-ray fluxes.  Using the constraints on the dark matter distribution of Draco put forth in Ref.~\\cite{couchman}, we have calculated maximal and minimal annihilation rates (and corresponding gamma-ray fluxes), considering both a cusped (NFW) profile and a profile with flat core. The variation that we find in the annihilation rate between even these two extreme scenarios is less than three orders of magnitude.  We then proceeded to compare these rates to the sensitivity of MAGIC and GLAST. MAGIC is a currently operating ground based gamma-ray telescope, while GLAST is a satellite based gamma-ray detector scheduled to be deployed in 2007. We find that while both MAGIC and GLAST have the ability to detect dark matter in Draco in some scenarios ({\\it ie.} a maximal NFW profile, low WIMP mass and favorable annihilation modes), dark matter in Draco can go undetected by these experiments in other cases. In some extreme cases (such as the non-thermal generation of neutralinos in anomaly mediated supersymmetry breaking models, for example), however, the lack of a detection by GLAST of gamma-rays from Draco could successfully rule out models, even if the most conservative halo model were assumed.  Finally, we have also discussed the implications of the recent possible detection of Draco by the ground based gamma-ray detector, CACTUS. We find that to produce the signal reported by the CACTUS collaboration, an annihilation rate of dark matter in Draco is needed which is three to four orders of magnitude larger than can be accommodated for an NFW profile and an annihilation cross section consistent with thermally generated dark matter. Non-thermally produced dark matter and/or extremely high densities of dark matter in Draco would therefore be required to generate this signal. We also find that the CACTUS signal appears to be in conflict with the null results from the region by the EGRET experiment for most choices of the dark matter's dominant annihilation modes. If the dark matter almost entirely annihilates to tau pairs, however, this conflict can be (marginally) avoided. If CACTUS is in fact detecting this very large flux of gamma-rays from Draco, both MAGIC and GLAST should easily be able to confirm this result.    \\bigskip     We would like to thank Gianfranco Bertone, Joakim Edsj\\\"o, Francis Halzen, Peter Marleau, Karl Mannheim and Mani Tripathi for helpful discussions. DH is supported by the US Department of Energy and by NASA grant NAG5-10842. LB is supported by the Swedish Science Research Council (VR)."
    },
    "0512/astro-ph0512308_arXiv.txt": {
        "abstract": "We present an analysis of Chandra observations of the galaxy clusters A2670 and A2107. Their cD galaxies have large peculiar velocities ($>200~\\rm\\: km\\: s^{-1}$) and thus the clusters appear to be undergoing mergers.  In A2670, we find a comet-like structure around one of the brightest galaxies. At the leading edge of the structure, there is a cold front. The mass of the X-ray gas in the comet-like structure suggests that the galaxy was in a small cluster or group, and its intracluster medium (ICM) is being stripped by ram-pressure. The regions of cool interstellar medium (ISM) of the cD galaxies in A2670 and A2107 are very compact. This is similar to the brightest galaxies in the Coma cluster, which is also a merging cluster. In each galaxy, the short cooling time of the ISM requires a heating source; the compact nature of the ISM makes it unlikely that the heating source is a central active galactic nucleus (AGN). ",
        "introduction": "It is generally believed that dark matter constitutes a large fraction of the mass in the Universe. Among various theories of dark matter, cold dark matter (CDM) theory provides a remarkably successful description of large-scale structure formation and is in good agreement with a large variety of observational data. This model predicts that small objects are the first to form and that these then amalgamate into progressively larger systems. In this model, clusters of galaxies are considered to be the objects that have recently formed via mergers of subclusters.  A cluster merger is one of the most spectacular events in the Universe. In a major merger, subclusters collide at velocities of $\\gtrsim 1000\\rm\\; km\\; s^{-1}$ and release gravitational energies of $\\sim 10^{64}$~erg.  A giant elliptical galaxy (cD galaxy) is often located at the center of a cluster. In general, the peculiar velocities of the cD galaxies are much smaller ($\\lesssim 200\\:\\rm km\\: s^{-1}$; \\authorcite{oeg01} \\yearcite{oeg01}) than the velocity dispersions of other galaxies in the clusters ($\\sim 1000\\:\\rm km\\: s^{-1}$). This means that the cD galaxies are nearly at rest in the cluster potential wells. However, the cD galaxies in some clusters have large peculiar velocities ($\\gtrsim 200\\:\\rm km\\: s^{-1}$).  Since a cluster merger may disturb a cD galaxy from its resting place at the bottom of the potential, the large peculiar velocity of the cD galaxy is a good indicator of mergers.   In this paper, we present the results of Chandra observations of two clusters containing rapidly moving cD galaxies in order to study the effects of cluster mergers on the cluster galaxies (including the cD galaxies) and the nature of substructures within clusters.  Recently, Oegerle and Hill (\\yearcite{oeg01}) studied the redshifts of cluster galaxies and their distributions intensively. They observed redshifts of $50$--$300$ galaxies per cluster and studied the peculiar velocities of cD galaxies in the clusters.  For several clusters, they compiled published data. Among 25 clusters, they found 4 clusters for which the redshifts of the cD galaxies are significantly different from the cluster means within $1.5\\; h_{75}^{-1}$~Mpc of the cDs, where the Hubble constant is $H_0=75\\; h_{75}\\;\\rm km\\; s^{-1}\\; Mpc^{-1}$. The clusters are A2052, A2107, A2199, and A2670. The peculiar velocities of the cD galaxies are $\\sim 250$--$400\\;\\rm km\\; s^{-1}$. Among them, A2052 and A2199 have already been observed by Chandra.  For these clusters, although interesting X-ray structures are observed in the cluster cores, no X-ray features related to the motion of the cD galaxies have been reported \\citep{bla01,joh02,kaw03}.  In this paper, we report on Chandra X-ray observations of the remaining two clusters that have fast moving cDs, A2107 and A2670.  The peculiar velocities of the cD galaxies are $v_{\\mathrm{P}}=433\\:\\rm km\\: s^{-1}$ for A2670, and $270\\:\\rm km\\: s^{-1}$ for A2107 \\citep{oeg01}.  For A2670, Hobbs and Willmore (\\yearcite{hob97}) analyzed the existing X-ray and optical data. While the X-ray morphology on large scales is regular, the cluster has complicated structures at the center. Hobbs and Willmore (\\yearcite{hob97}) detected four point sources near the center of the ROSAT HRI image as is shown in their figure~2.  Besides the cD galaxy, they argued that two of the sources (1: \\timeform{23h54m09s.5}, \\timeform{-10D25'48''}, and 3: \\timeform{23h54m07s.0}, \\timeform{-10D25'17''}) were coincident with cluster galaxies.  The X-ray contours near the cluster center are elongated towards east and west, which suggests that the cD galaxy and sources~1 and/or~3 are interacting. \\citet{bir94} indicated that the cluster consists of four subclumps from the spatial and redshift distributions of the galaxies. The centers of three subclumps (Clumps A, B, and C) are very close to the cluster center (within a few arcminutes from the cD galaxy).  This may show that these subclumps are near their point of closest approach to one another.  \\citet{bir94} showed that the masses of the subclumps are comparable; the subclumps may induce strong gas motion in the hot intracluster medium (ICM). On the other hand, the X-ray morphology of A2107 is not reported to be irregular \\citep{buo96}, although the detailed X-ray structures in the central region have not been investigated. Both A2670 and A2107 have relatively weak central X-ray peaks, which may show that the gas motions in the ICM disrupt the central gas structures of the clusters. In terms of classical cooling flows, their peak strengths are equivalent to the mass deposition rates of $\\dot{M}\\lesssim 50\\: M_\\odot \\rm yr^{-1}$ \\citep{whi97}, although the classical cooling flow model is no longer valid \\citep{mak01,pet01,kaa01,tam01}.  In the following analysis, we use the cosmological parameters of $\\Omega_0=0.27$, $\\lambda=0.73$, and the Hubble constant of $H_0=70\\:\\rm km\\: s^{-1}\\: Mpc^{-1}$ unless otherwise mentioned. This means that $1\\arcsec$ corresponds to 1.45~kpc for A2670 ($z=0.0765$) and 0.813~kpc for A2107 ($z=0.0411$). ",
        "conclusions": "We present an analysis of Chandra observations of the galaxy clusters A2670 and A2107. The peculiar velocities of the cD galaxies are unusually large ($>200~\\rm\\: km\\: s^{-1}$) and the clusters are undergoing mergers.  To the west of the cD galaxy of A2670, we find a galaxy having a comet-like X-ray structure (comet galaxy). The leading edge of the structure is a cold front. From the pressure profile across the cold front, we estimate the velocity of the galaxy in the cluster A2670. The mass of X-ray gas in the comet-like structure is too large to have been produced by stellar mass loss within the galaxy itself.  Thus, it is likely that the galaxy was at the center of a small cluster or group, and that its intracluster or intragroup medium is being stripped by ram-pressure.  The sizes of the cooler, X-ray interstellar medium (ISM) regions of the cD galaxies in A2670 and A2107 are very small. This is similar to the brightest galaxies in the Coma cluster, which is also a merging cluster.  Since the cooling time of the ISM is small, there must be some heating sources in the galaxies.  The compactness of the ISM indicates that the heating sources are not the AGNs in the galaxies.  We suggest thermal conduction or hydrodynamical heating by `tsunamis' as possible heat sources to balance the observed X-ray emission.  \\vspace{5mm}  We are grateful to Tracy Clarke for useful comments.  Y. F. is supported in part by a Grant-in-Aid from the Ministry of Education, Culture, Sports, Science, and Technology of Japan (14740175).  Support for this work was provided by the National Aeronautics and Space Administration primarily through Chandra award GO4-5137X, but also through GO4-5133X, GO4-5149X and GO4-5150X, issued by the Chandra X-ray Observatory, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-39073. G. R. S. acknowledges the receipt of an ARCS fellowship."
    },
    "0512/astro-ph0512552_arXiv.txt": {
        "abstract": "In this letter, we comment on the robustness of putative cool ($kT \\simeq 0.2$~keV) accretion disc components in the X-ray spectra of the most luminous ($L_{X} \\simeq 10^{40}$~erg/s) ultra-luminous X-ray sources (ULXs) in nearby normal galaxies.  When compared to stellar-mass black holes, the low disc temperatures observed in some ULXs may imply intermediate-mass black hole primaries.  Recent work has claimed that such soft excesses are unlikely to be actual disc components, based on the lack of variability in these components, and in the overall source flux.  Other work has proposed that alternative phenomenological models, and complex Comptonisation models, rule-out cool disc components in ULX spectra.  An inspection of the literature on Galactic stellar-mass black holes and black hole candidates demonstrates that the flux behaviours seen in specific ULXs are consistent with phenomena observed in well-known Galactic X-ray binaries.  Applying Comptonisation models to simulated disc blackbody plus power-law spectra shows that at the sensitivity achieved in even the best ULX spectra, Comptonisation fits are highly model-dependent, and do not yield meaningful constraints on the accretion flow.  In contrast, the need for a soft, thermal component does not appear to be model-dependent.  As we have previously noted, soft thermal components in ULX spectra may not represent accretion discs, but present alternatives to this interpretation are not robust. ",
        "introduction": "The nature of ultra-luminous X-ray sources (ULXs, $L_{X} \\geq 2\\times 10^{39}$~erg/s) in nearby normal galaxies has been a topic of great interest in the {\\it Chandra} and {\\it XMM-Newton} era.  The apparent luminosity of these sources can exceed the isotropic Eddington limit for a $10~M_{\\odot}$ black hole, leading to the possibility that some ULXs may harbour intermediate-mass black holes ($10^{2-5}~M_{\\odot}$). Initially, this debate was framed in absolute terms, e.g., What is the nature of the ULX phenomenon?  Such questions were likely posed in an erroneously simple manner.  Every new class of sources is soon divided into subclasses with further study, and is found to be a broad designation covering a heterogeneous group of sources.  The fact that Galactic stellar-mass black holes can appear to be mildly super-Eddington (see, e.g., McClintock \\& Remillard 2005) may argue that many ULXs are stellar-mass black holes.  However, a growing subset (fewer than 10 presently) is emerging that may represent a class of intermediate mass black holes.  The lessons learnt from decades of X-ray spectral and timing studies of Galactic stellar-mass black hole binaries and black hole candidates would seem to indicate that a small number of the most luminous ULXs ($L_{X} \\simeq 10^{40}$~erg/s, and above) may harbour black holes of a few hundred or few thousand solar masses (see, e.g., Colbert \\& Mushotzky 1999; Strohmayer \\& Mushotzky 2003; Miller et al.\\ 2003; Cropper et al.\\ 2004; Miller, Fabian, \\& Miller 2004a; Kaaret, Ward, \\& Zezas 2004; Miller \\& Colbert 2004).  Intermediate mass black holes are implied in these sources via Eddington luminosity scaling, and/or multi-wavelength properties that argue against beaming, and/or scaling characteristic frequencies found in the X-ray flux, and/or scaling apparent inner disc temperatures in the X-ray spectra.  In the case of inner disc temperatures, intermediate mass black holes may be implied because the temperatures implied ($kT \\simeq 0.2$~keV) are well below the temperatures typically measured in stellar-mass black holes accreting near their Eddington limits ($kT \\simeq 1-2$~keV), and $T \\propto M^{-1/4}$ for standard black hole accretion discs (see, e.g., Miller, Fabian, \\& Miller 2004a).  However, the robustness of these soft components has recently been questioned.  It has been suggested that the variability properties of ULXs and the soft components in particular, argue against associating them with discs (e.g. Dewangan, Griffiths, \\& Rao 2005; Goad et al.\\ 2005).  Other work has suggested that alternative spectral models may be more appropriate, and may rule-out the possibility of cool discs and intermediate mass black holes in the most luminous ULXs (e.g. Goad et al.\\ 2005, Roberts et al.\\ 2005).  In this letter, we examine these arguments in detail, in the context of better-understood Galactic stellar-mass black hole binaries and AGN. ",
        "conclusions": "We have shown that a variety of arguments against interpreting apparent cool thermal components in the spectra of very luminous ULXs are contradicted by published results from stellar mass black holes and black hole candidates, by simple physical considerations, and by critically examining the ability of complex models to yield meaningful results in low signal-to-noise regimes.  The absence of direct correlations between apparent disc temperatures and fluxes observed in some ULXs is a phenomenon commonly observed in stellar-mass black holes and black hole candidates.  Moreover, limited fast X-ray variability has been observed in stellar-mass black holes and black hole candidates, even in hard phases at low fractional Eddington luminosities.  Spectroscopy of stellar-mass black holes and physical considerations strongly argue for associating thermal distributions with the low energy portion of a spectrum, and scattered or non-thermal distributions with the high energy portion.  Finally, we have shown that even at the sensitivity achieved in the best present ULX spectra, complex Comptonisation models can yield false inferences and give an incorrect picture of the accretion flow geometry.  In the case of NGC 1313 X-1 and M81 X-9, present spectra have been sufficient to show that cool disc plus power-law models are at least 5$\\sigma$ better than single-component models, broken power-law models, and models with significantly sub-solar abundances in gas along the line of sight (Miller et al.\\ 2003; Miller, Fabian, \\& Miller 2004b).  The same work has shown that cool discs are robust against specific choices of disc models and specific choices of hard component models.  Even in cases where Comptonisation models are invoked to argue against interpreting soft excesses as discs, cool disc components are required to achieve acceptable fits (e.g. Goad et al.\\ 2005).  This further highlights the robustness of cool thermal components, and the perils of Comptonisation models in low signal-to-noise spectra.  The cool disc interpretation is merely the most plausible one, based on a comparison to other black holes, and principally stellar-mass black holes.  As we have previously noted, however, cool thermal components in the spectra of very luminous ULXs are not necessarily disc components, and not all soft components in the spectra of accreting black holes are well-understood.  The soft excess observed in some Seyfert-1 galaxies is too hot to easily be attributed to a disc, and its nature is uncertain.  Miller, Fabian, \\& Miller (2004a) noted this aspect of Seyfert-1 spectra, in anticipation of the possibility that some of the ULXs which have cool thermal components may actually be background AGN.  Indeed, this proved to be the case for the ULX Antennae X-37 (Clark et al.\\ 2005).  It is possible that the soft excess in genuine ULXs is similar to that in some Seyfert-1 spectra.  However, if both phenomena are due to relativistically-blurred disc emission lines (Crummy et al.\\ 2005), then the soft component in ULXs indicates that the emission is isotropic and intermediate-mass black holes may still be required via simple Eddington scaling arguments.  We note that a preliminary investigation we have undertaken reveals that {\\it XMM-Newton} spectra of ULXs with apparent cool thermal components can also be fit acceptably with blurred disc reflection models.  While this letter demonstrates that cool thermal components are likely more robust than than inferences derived from models for the hard component in ULX spectra, it also serves to demonstrate that present ULX spectra are intrinsically a low signal-to-noise regime.  All spectral fits, whether simple or complex, must be regarded cautiously in such circumstances, but the most robust conclusions probably derive from simple models that can be constrained by the data.  Very deep observations of ULXs with {\\it XMM-Newton} (350~ksec or longer) may be the only means of parsing the nature of soft and hard components in very luminous ULXs in more detail."
    },
    "0512/astro-ph0512287_arXiv.txt": {
        "abstract": "Verification of theoretical predictions of an oscillating behavior of the fine-structure constant, $\\alpha$, with cosmic time requires high precision measurements at individual redshifts, while in earlier studies the mean $\\Delta\\alpha/\\alpha$ values averaged over wide redshift intervals were usually reported. This requirement can be met via the Single Ion Differential $\\alpha$ Measurement (SIDAM) procedure. We apply SIDAM to the FeII lines associated with the damped Ly$\\alpha$ system observed at $z = 1.15$ in the spectrum of HE0515--4414. The weighted mean calculated on base of carefully selected 34 FeII pairs is $\\langle \\Delta\\alpha/\\alpha \\rangle = (-0.07\\pm0.84)\\times10^{-6}$. The precision of this estimate represents the absolute improvement with respect to what has been done in the measurements of $\\Delta\\alpha/\\alpha$. ",
        "introduction": "The Sommerfeld fine-structure constant, $\\alpha \\equiv e^2/\\hbar c$, which describes electromagnetic and optical properties of atoms, is the most suitable for time variation tests in both laboratory experiments with atomic clocks and astronomical observations (for a review see, e.g.,  \\cite[Barrow 2005]{B05}).  The question whether or not the fine-structure constant varied at different cosmological epochs can be answered only through observations of quasar absorption-line spectra. The main requirement of such studies~-- precise line position measurements at the level of 10$^{-7}$-10$^{-8}$~-- can be fulfilled only at giant optical telescopes equipped by high resolution spectrographs.  Theoretically the effects of inhomogeneous space and time evolution of $\\alpha$ were considered by \\cite{Ma84} and \\cite{MB04}. Most recently \\cite{F05} suggested a damped-oscillation-like behavior of $\\alpha$ as a function of cosmic time $t$. It is apparent that to study such irregular changes in $\\alpha$, we need to achieve high precision in the measurements of $\\Delta\\alpha/\\alpha$ at {\\it individual} redshifts, contrary to the averaging procedure over many redshifts which is usually used to decrease uncertainties of the mean values $\\langle \\Delta\\alpha/\\alpha \\rangle$ (\\cite[Murphy et al. 2004]{Mu04}, and references therein).  The uncertainties of individual values of $\\Delta\\alpha/\\alpha = (\\alpha_z - \\alpha_0)/\\alpha_0$,\\, (here $\\alpha_0$ and $\\alpha_z$ are the values of $\\alpha$ at epoch $z=0$ and at redshift $z$, respectively) are currently known at the level of a few ppm (parts per million) (\\cite[Quast et al. 2004]{Q04};  \\cite[Chand et al. 2004]{Ch04}). In both cases the standard many-multiplet (MM) method (\\cite[Dzuba et al. 2002]{Dz02}) has been used. Further modification of the MM method (\\cite[Levshakov et al. 2005a]{L1}) resulted in a new methodology for probing the cosmological variability of $\\alpha$ on base of pairs of FeII lines observed in {\\it individual exposures} from a high resolution spectrograph (henceforth referred to as SIDAM~-- Single Ion Differential $\\alpha$ Measurement).  The basic idea behind SIDAM was to avoid the influence of small spectral shifts due to ionization inhomogeneities within the absorbers and due to non-zero offsets between different exposures. The individual offsets can affect the shape of the line profiles during rebinning and coadding procedures when exposures are combined together to increase signal-to-noise, S/N, ratio (examples are given in \\cite[Levshakov et al. 2005a]{L1}).  \\firstsection % ",
        "conclusions": ""
    },
    "0512/astro-ph0512215_arXiv.txt": {
        "abstract": "We present low-resolution ($64 < R < 124$) mid-infrared (8--38 $\\micron$) spectra of two $z \\approx 1.3$ ultraluminous infrared galaxies with $\\lfir \\approx 10^{13} \\Lsun$: \\firstobjectlong \\ and \\secondobjectlong.  The spectra were taken with the \\textit{Infrared Spectrograph} (\\irs) \\ on board the \\spitzer.  Both objects were discovered in a \\textit{Spitzer}/\\mips \\ survey of the \\bootes \\ field of the NOAO Deep Wide-Field Survey (NDWFS).  \\firstobjectlong \\ is a bright 160 $\\micron$ source with a large infrared-to-optical flux density ratio.  Previous authors provide evidence for a foreground lens and estimate an amplification of $\\la$10, although this factor is presently poorly constrained.  The 6.2, 7.7, 11.3, and 12.8 $\\micron$ PAH emission bands in its \\irs \\ spectrum indicate a redshift of $z \\approx 1.3$.  The large equivalent width of the 6.2 $\\micron$ PAH feature indicates that at least 50\\% of the mid-infrared energy is generated in a starburst, an interpretation that is supported by a large [\\ion{Ne}{2}]/[\\ion{Ne}{3}] ratio and a low upper limit on the X-ray luminosity.  \\secondobjectlong \\ has the brightest 24 $\\micron$ flux (10.55 mJy) among optically faint ($R > 20$) galaxies in the NDWFS.  Its mid-infrared spectrum lacks emission features, but the broad 9.7 $\\micron$ silicate absorption band places this source at $z \\approx 1.3$.  Optical spectroscopy confirms a redshift of $z = 1.293 \\pm 0.001$.  Given this redshift, \\secondobjectlong \\ has among the largest rest-frame 5 $\\micron$ luminosities known.  The similarity of its SED to those of known AGN-dominated ULIRGs and its lack of either PAH features or large amounts of cool dust indicate that the powerful mid-infrared emission is dominated by an active nucleus rather than a starburst.  Our results illustrate the power of the \\irs \\ in identifying massive galaxies in the well-known ``redshift desert'' between $1 < z < 2$ and in discerning their power sources.  Because they are bright, \\firstobjectlong \\ (pending future observations to constrain its lensing amplification) and \\secondobjectlong \\ are useful $z > 1$ templates of a high luminosity starburst and AGN, respectively. ",
        "introduction": "\\label{sec:intro}  Observations at long wavelengths, particularly those made by the \\textit{Infrared Astronomical Satellite} (\\textit{IRAS}) and the \\textit{Infrared Space Observatory} (\\textit{ISO}), have revealed a low-redshift population of galaxies which emit most of their luminosity in the far infrared. Ultraluminous infrared galaxies \\citep[ULIRGs;][]{Soifer84,Sanders88,SandersMirabel96}, which have far-infrared luminosities greater than $\\lfir = 10^{12}~{\\rm L}_{\\odot}$, represent the luminous tail of this population.  Infrared number counts and the cosmic infrared background provide strong evidence that the ULIRG population becomes progressively more important at high redshift \\citep{Dole01, Elbaz02, LeFloch04}.  In addition, submillimeter galaxies (SMGs) have been revealed as likely ULIRG analogs at high redshift \\citep{Smail97b, Barger98, Hughes98, Ivison98, Eales00, Scott02}.  However, although the energy generation mechanisms of these extremely luminous objects have been a topic of great interest and study, the relative fraction of high redshift SMGs powered by buried AGN is still a matter of debate, as are the relative contributions of AGN and starbursts to the energy output of ULIRGs at low redshift.  Mid-infrared diagnostics of the energetics of galaxies have been developed using ground-based \\citep{Roche91} and space-based \\citep{Genzel98,Lutz00,Laurent00} observations. In general, the mid-infrared spectra of starbursts are characterized by low-excitation fine structure lines, polycyclic aromatic hydrocarbon (PAH) features, and a weak 3--6 $\\micron$ rest-frame continuum.  In contrast, the mid-infrared spectra of AGN exhibit high-excitation emission lines, very weak or absent PAHs, and a strong 3--6 $\\micron$ rest-frame continuum.  Until recently, the sensitivity of available spectrometers has limited the application of these diagnostics to the brightest galaxies at low redshift.  The highly sensitive instruments aboard the \\spitzer \\ \\citep{Werner04}, especially the \\textit{Infrared Spectrograph} \\citep[\\irs;][]{Houck04}, are powerful tools for performing detailed studies of ULIRGs at higher redshifts than were previously possible with \\textit{IRAS} or \\textit{ISO}.  In order to identify high-redshift ULIRGs for detailed study, the \\irs \\ and \\mips \\ instrument teams have completed a mid-infrared imaging survey of the 9 deg$^2$ \\bootes \\ region of the NOAO Deep Wide-Field Survey (NDWFS; Jannuzi \\etal, in preparation).  The NDWFS is a $B_WRIK$ imaging survey reaching 3$\\sigma$ point-source depths of approximately 27.7, 26.7, 26.0, and 19.6 Vega magnitudes, respectively.  In particular, we used the NDWFS Data Release 3 catalog.  The mid-infrared imaging was carried out with the \\textit{Multiband Imaging Photometer for Spitzer} \\citep[\\mips;][]{Rieke04} and reaches 5$\\sigma$ limits of 0.28, 35, and 100 mJy at 24, 70, and 160 $\\micron$, respectively. Most (8.5 deg$^2$) of this field has also been surveyed using the \\textit{Infrared Array Camera} \\citep[\\irac;][]{Fazio04}, also aboard \\textit{Spitzer}, to 5$\\sigma$ sensitivities of 6.4, 8.8, 51, and 50 $\\mu$Jy at wavelengths of 3.6, 4.5, 5.8, and 8.0 $\\micron$, respectively \\citep{Eisenhardt04}.  The ACIS-I instrument \\citep{Garmire03} aboard the \\textit{Chandra X-ray Observatory} was also used to map the entire \\bootes \\ field down to a limiting sensitivity of $\\sim 4 \\times 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ in the energy range 0.5--7 keV \\citep{Murray05}.  The \\bootes \\ field also overlaps with the 20 cm FIRST Survey, which has a limiting sensitivity of approximately 1 mJy \\citep{White97}.  The NDWFS \\bootes \\ field was chosen for the MIPS imaging survey because of its low infrared background and high-quality, deep $B_WRI$ imaging.  These properties allow the selection of mid-infrared sources which are bright enough for \\irs \\ spectroscopy but which have very faint optical counterparts, and are therefore likely at high redshift. \\citet{Houck05} describe the results of \\irs \\ observations of 30 extreme ($f_{\\nu}(24\\micron) > 0.75$ mJy; $R > 24$ mag) sources identified in \\bootes.  The vast majority of the 17 objects for which the mid-infrared spectra yielded redshifts are probably obscured AGN at $z \\approx 2$.  We present here a detailed analysis of two objects which were also selected from the \\bootes \\ field, but which have unique infrared and optical properties.  In particular, we present \\irs \\ and optical spectroscopy, and constrain the spectral energy distributions of these objects using data from the previously described multiwavelength \\bootes \\ surveys, as well as new submillimeter data.  \\firstobjectlong \\ (hereafter \\firstobject) was selected from a catalog of objects that were detected in all three MIPS bands (thus the prefix ``MIPS'').  Following the naming scheme of \\citet{Borys05}, we use the NDWFS $I$-band coordinates\\footnote{For completeness, the 24 $\\micron$ coordinates are $14^{\\rm h}28^{\\rm m}24\\fs07$, +35$\\degr$26$\\arcmin$19\\farcs14 (J2000).} to form the name of this object. \\firstobject \\ was chosen for follow-up because it is bright at 160 $\\micron$ ($f_{\\nu}(160\\micron) = 430 \\pm 90$ mJy), has red optical--near-infrared colors ($R-K > 5$), and is unresolved in the optical.  The large 160 $\\micron$ flux indicates that \\firstobject \\ contains large amounts of cold dust.  The red optical--near-infrared colors and compact morphology suggest that it resides at $z > 1$.  The combination of the 160 $\\micron$ flux and the high redshift implies that \\firstobject \\ probably has a large far-infrared luminosity ($\\lfir \\simeq 10^{13} {\\rm L}_{\\odot}$).  \\citet{Borys05} presented the spectral energy distribution (SED) of this object from the optical through the radio.  They argued that while the emission shortward of 1 $\\micron$ (observed frame) is likely dominated by a $z \\approx 1$ lens, the near-infrared emission is dominated by the 24 $\\micron$ source.  Based on a cool ($\\sim$43 K) dust temperature, adherence to the far-infrared--radio correlation, and the presence of a prominent 1.6 $\\micron$ stellar bump, they argued that \\firstobject \\ is a dusty starburst.  As discussed in \\S \\ref{sec:firstPAH}, the \\irs \\ spectrum we present provides a rough redshift of $z \\approx 1.3$. \\citet{Borys05} used this estimate to identify H$\\alpha$ in the near-infrared spectrum and obtained $z = 1.325 \\pm 0.002$.  In combination with their SED, this redshift confirms a high luminosity of $\\lfir = (3.2 \\pm 0.7) \\times 10^{13} \\Lsun$.  They also noted that the instrinsic luminosity may be up to a factor of 10 lower if lensing is important.  Here we present additional diagnostics to determine the energetics of \\firstobject.  \\secondobjectlong \\ (hereafter \\secondobject), was selected from a catalog of objects detected at 24 $\\micron$ (thus the prefix ``SST24'').  Following the convention of \\citet{Houck05}, we use the 24 $\\micron$ position to derive the name of this object. \\secondobject \\ has the brightest 24 $\\micron$ flux ($f_{\\nu}(24 \\micron) = 10.55$ mJy) among mid-infrared sources with faint optical counterparts ($R > 20$).  The faint $R$-band magnitude ($R = 22.78$) implies that \\secondobject \\ is at high redshift, and the 24 $\\micron$ flux indicates a large luminosity.  The plan of this paper is as follows: \\S \\ref{sec:Observations} describes the mid-infrared \\irs \\ observations of both sources, as well as optical spectroscopy and submillimeter imaging of \\secondobject.  The results for each object are presented in \\S \\ref{sec:firstobjectlong} and \\S \\ref{sec:secondobjectlong}. Finally, our conclusions are discussed in \\S \\ref{sec:Conclusions}. Throughout, we use the following cosmological parameters: $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_{0} = 70 $ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:Conclusions}  We have presented the \\irs \\ spectra of two $z \\approx 1.3$ ULIRGs selected from the \\bootes \\ region of the NDWFS.  \\firstobject \\ has a high 6.2 $\\micron$ PAH equivalent width, high [\\ion{Ne}{2}]/[\\ion{Ne}{3}] ratio, and low ${\\rm L}_{2-10{\\rm keV}}/\\lfir$. These properties, in combination with the prominent 1.6 $\\micron$ stellar bump, cold ($\\sim$43 K) dust temperature, and concordance with the radio-FIR correlation noted by \\citet{Borys05}, indicate that \\firstobject \\ is dominated by a starburst.  With $\\lfir = 3.2 \\times 10^{13} \\Lsun$ (modulo its unknown lensing amplification), \\firstobject \\ is extremely luminous compared to low-redshift starbursts.  However, PAH-dominated spectra of comparably luminous objects are beginning to be discovered at $z \\ga 2$ \\citep{Lutz05, Yan05}.  \\firstobject \\ may be analogous to these more distant luminous starbursts, but is at a redshift where it is more easily studied in detail.  Its use as a template for higher redshift populations rests on our ability to understand its lensing properties \\citep{Borys05}, which will require high-resolution imaging to separate out the contribution from the foreground object.  At a similar redshift, \\secondobject \\ is among the most luminous mid-infrared sources known (as measured at 5 $\\micron$, rest-frame). It appears to be powered mainly by an AGN: it has a very small 6.2 $\\micron$ PAH equivalent width; exhibits a mid-infrared SED similar to those of AGN-dominated ULIRGs; and lacks large amounts of cool dust, as evidenced by its moderate silicate optical depth and non-detections at 160 and 350 $\\micron$.  The first sizable high-redshift samples of dusty sources observed with the \\irs \\ have focused on objects at $z \\ga 2$ \\citep{Houck05, Yan05}.  The submillimeter galaxies also lie at this redshift. Relatively less is known about the population of dusty sources at $1 < z < 2$. This redshift range has been traditionally difficult to access through optical spectroscopy, and has come to be known as the ``redshift desert''.  Indeed, the optical spectra of both objects and the near-infrared spectrum of \\firstobject \\ show very few emission lines.  A redshift identification for either source would have been dubious without its mid-infrared \\irs \\ spectrum.  Nevertheless, a large fraction of the stars in local L$_{\\ast}$ galaxies were likely formed in this redshift range \\citep[e.g.][]{Dickinson03,Drory05}. Objects such as \\firstobject \\ and \\secondobject \\ provide examples of how the \\irs \\ will allow the identification and detailed study of objects in this critical redshift range."
    },
    "0512/astro-ph0512509_arXiv.txt": {
        "abstract": "Sterile neutrinos with masses in the keV range are viable candidates for the warm dark matter. We analyze existing data for the extragalactic diffuse X-ray background for signatures of sterile neutrino decay. The absence of detectable signal within current uncertainties of background measurements puts model-independent constraints on allowed values of sterile neutrino mass and mixing angle, which we present in this work. ",
        "introduction": "\\label{sec:intro}  At present time there exists an extensive body of evidence that most of the matter in the Universe is composed of new, yet undiscovered particles -- dark matter (DM).  Observations of (i) galactic rotation curves, (ii) cosmic microwave background radiation, (iii) gravitational lensing, and (iv) X-ray emission of hot gas in galaxy clusters provide independent measurements of DM content of the Universe.  Another major experimental discovery of the recent decade is that of neutrino oscillations.  There are separate measurements of neutrino oscillations in solar neutrinos (\\cite{SNO}), atmospheric neutrinos (\\cite{superK}), and reactor neutrinos (\\cite{Kamland}).  Neutrino oscillations can be explained if neutrino is a massive particle, contrary to the Standard Model assumption. This means that along with the usual left-handed (or \\emph{active}) neutrinos there may exist also right-handed or \\emph{sterile neutrinos}.  Conventional sea-saw mechanism (\\cite{Minkowski,Yanagida,gell,ramond,mohapatra,glashow}) of generation of small active neutrino masses implies that the sterile neutrinos are heavy (usually of the order of GUT energy scale $\\sim 10^{10}- 10^{15}$~GeV) and that their mixing with usual matter is of the order $\\sin\\theta \\sim 10^{-10} - 10^{-15}$. In addition of the smallness of neutrino masses, models of this type can explain baryon asymmetry of the Universe via thermal leptogenesis (\\cite{fukugita}) and anomalous electroweak fermion number non-conservation (\\cite{kuzmin}). However they do not offer a DM candidate.  Recently it was proposed that neutrino oscillations, the origin of the dark matter, and baryon asymmetry of the Universe can be consistently explained in the model called \\emph{neutrino Minimal Standard Model} ($\\nu$MSM)~(\\cite{Misha05a,Misha05b}). This model is a natural extension of minimal Standard Model (MSM), where three right-handed neutrinos are introduced into the MSM Lagrangian. In this extension neutrinos obtain Dirac masses via Yukawa coupling analogous to the other quarks and leptons of MSM and in addition Majorana mass terms are allowed for right-handed neutrinos. Unlike conventional see-saw scenarios, all of these Majorana masses (which are roughly equal to the masses of corresponding sterile neutrinos) are chosen such that the mass of the lightest sterile neutrino is in the keV range and the other ones are $\\lesssim 100$~GeV --- below electroweak symmetry breaking scale. In this model a role of the dark matter particle is played by the lightest sterile neutrino.  The existence of a relatively light sterile neutrino has nontrivial observable consequences for cosmology and astrophysics. It was proposed in~\\cite{Dodelson:93} that a sterile neutrino with the mass in the keV range may be a viable ``warm'' DM candidate. The small mixing angle ($\\sin \\theta \\sim 10^{-6} - 10^{-4}$) between sterile and active neutrino ensures that sterile neutrinos were never in thermal equilibrium in the early Universe and thus allows their abundance to be smaller than the equilibrium one.  Moreover, a sterile neutrino with these parameters is important for the physics of supernova (\\cite{fryer}) and was proposed as an explanation of the pulsar kick velocities~(\\cite{kusenko,2003PhRvD..68j3002F,2004PhRvD..70d3005B}).  In addition to the dominant decay mode into three active neutrinos, the light (with mass $m_s\\lesssim 1$~MeV) sterile neutrino can decay into an active one and a photon with the energy $E_\\gamma=m_s/2$.  Thus, there exists a possibility of direct detection of neutrino decay emission line from the sources with big concentration of DM, e.g. from the galaxy clusters~(\\cite{Fuller:01b}).  Similarly, the signal from radiative sterile neutrino decays accumulated over the history of the Universe could be seen as a feature in the diffuse extragalactic background light spectrum. This opens up a possibility to study the physics beyond the Standard Model using astrophysical observations.  Recently there has been a number of works devoted to the analysis of the possibility to discover sterile neutrino radiative decays from X-ray observations~(\\cite{Fuller:01b,Mapelli:05}). For example, it was argued by \\cite{Fuller:01b} that if sterile neutrinos composed 100\\% of all the DM, one should be able to detect the DM decay line against the background of the X-ray emission from the Virgo cluster. According to \\cite{Fuller:01b} the non-detection of the line puts an upper limit $m_s<5$~keV on the neutrino mass (this limit was, however recently revised in \\cite{abazajian05}, who finds the restriction $m_s<8$~keV).  It was also noted by~\\cite{Fuller:01b,Mapelli:05} that one can obtain even stronger constraints $m_s \\lesssim 2$~keV -- from diffuse extragalactic X-ray background (XRB) under the assumption, that the dark matter in the Universe is uniformly distributed up to the distances, corresponding to red shifts $z\\ll 1$. Together with the claim of~\\cite{Hansen:01,Viel:05}, putting lower bound $m_s>2$~keV on the neutrino mass from Lyman $\\alpha$-forest observations, this would lead to a very narrow window of allowed sterile neutrino masses, if not exclude it completely.  In this paper we re-analyze the limit imposed on the parameters of sterile neutrino by the observations of the diffuse X-ray background (XRB). For that we are processing actual astrophysical data of HEAO-1 and XMM-Newton missions. There are several motivations for this. Namely \\begin{enumerate} \\item All the above restrictions on sterile neutrino mass (\\cite{Fuller:01b,Dolgov:02,Mapelli:05,abazajian05}) are model dependent and based on the assumption that sterile neutrinos were absent in the early Universe at temperatures larger than few GeV.  Depending on the model, the relation between the mass of sterile neutrino, the mixing angle and the present-day sterile neutrino density $\\Omega_s$ does change. In fact, to compute the sterile neutrino abundance one needs to know whether there is any substantial lepton asymmetry of the universe at the time of sterile neutrino production, what is the coupling of sterile neutrino to other particles such as inflaton or super-symmetric particles, etc.\\footnote{For example, if the coupling of sterile neutrino to inflaton is large enough, the main production mechanism will be the creation of sterile neutrinos in inflaton oscillations rather than active-sterile neutrino transition.} Moreover, even if these uncertainties were removed, the reliable computation of the relic abundance of sterile neutrinos happens to be very difficult as the peak of their production falls on the QCD epoch of the universe evolution, corresponding to the temperature $\\sim 150$ MeV (\\cite{Dodelson:93}), where neither quark-gluon nor hadronic description of the plasma is possible. Therefore, before the particle physics model is fully specified and the physics of hadronic plasma is fully understood, one can not put a robust restriction on one single parameter of the model such as~$m_s$.  Therefore we aim in this paper at clear separation between the model independent predictions, based solely on astrophysical observations and any statements that depend on a given model and underlying assumptions. To this end we treat $m_s$ and $\\sin\\theta$ as two \\emph{independent} parameters and present the limits in the form of an ``exclusion plot'' in the $(m_s; \\Omega_s \\sin^22\\theta)$ parameter space.  It should be stressed that our data analysis is not based on any specific model of sterile neutrinos and as such can be applied to any ``warm'' DM candidate particle which has a radiative decay channel. In case of sterile neutrino the full decay width of this process is related to parameter $\\sin\\theta$ via Eq.~(\\ref{eq:4}) (see below).  \\item Contrary to the previous works~(\\cite{Fuller:01b,Mapelli:05}), we argue that non-isotropy or ``clumpyness'' of the matter distribution in the nearby Universe \\emph{does not} relax the limit on the neutrino mass.  Indeed, the fact that significant part of the dark matter at red shifts $z\\lesssim10$ is concentrated in galaxies and clusters of galaxies just means that the strongest signal from the dark matter decay should come from the sum of the signals from the compact sources at $z\\lesssim10$. Taking into account that DM decay signal from $z\\lesssim10$ is some 2 orders of magnitude stronger than that of from $z\\gtrsim10$, while the subtraction of resolved sources reduces the residual X-ray background maximum by a factor of 10, we argue that it would be wrong to subtract the contribution from the resolved sources from the XRB observations when looking for the DM decay signal. The form of XRB background spectrum with sources subtracted is, in fact, unknown and the assumption that it has a shape of initial spectrum, scaled down according to resolved fraction (as in \\cite{Fuller:01b}) requires additional justification.  \\item We find that more elaborate analysis of the data enables to put tighter limits on the allowed region of the parameter space $(m_s,\\Omega_s\\sin^2 2\\theta)$ from the XRB observations. The idea is that the cosmological DM decay spectrum is characterized not only by the total flux but also by a characteristic shape. Being present in the XRB spectrum, it would produce a local feature with some clear maximum and a width greater than spectral resolution of the instrument. Features of such a scale are clearly absent in the data (tinier features could be present in the spectrum due to e.g. element lines, but they can not produce a bump wider than spectral resolution). Therefore, one can find that adding to the standard broad continuum model of XRB the DM decay component in a wrong place results in decrease of the overall quality of the fit of the data by such a two-component model (increase of the $\\chi^2$ of the fit).  The condition that the two-component model provides an acceptable fit to the data imposes an upper limit on the flux in the DM decay component which is much more restrictive than the limit following from the condition that the flux of the DM component should not exceed the flux in the continuum component. \\end{enumerate}  The paper is organized as follows. In Section~\\ref{sec:dm-flux} we compute the contribution of the radiative decay of sterile neutrino to the diffuse X-ray background (XRB) and compare its shape with that of measured XRB. We discuss the effects of non-uniformness of the DM distribution in Section~\\ref{sec:uniform}. In Section~\\ref{sec:diffuse-bg} we obtain a model-independent exclusion region from HEAO-1 and XMM-Newton observations. ",
        "conclusions": ""
    },
    "0512/astro-ph0512023_arXiv.txt": {
        "abstract": "The dark energy crossing of the cosmological constant boundary (the transition between the quintessence and phantom regimes) is described in terms of the implicitly defined dark energy equation of state. The generalizations of the models explicitly constructed to exhibit the crossing provide the insight into the cancellation mechanism which makes the transition possible. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512353_arXiv.txt": {
        "abstract": "We show that contrary to what is expected from 1D stationary model atmospheres, 3D hydrodynamical modeling predicts a considerable influence of convection on the spectral properties of late-type giants. This is due to the fact that convection overshoots into the formally stable outer atmospheric layers producing a notable granulation pattern in the 3D hydrodynamical models, which has a direct influence on the observable spectra and colors. Within the framework of standard 1D model atmospheres the average thermal stratification of the 3D hydro model can not be reproduced with any reasonable choice of the mixing length parameter and formulation of the turbulent pressure.  The differences in individual photometric colors -- in terms of 3D versus 1D -- reach up to $\\sim0.2$\\,mag, or $\\Delta T_{\\rm eff}\\sim70$\\,K. We discuss the impact of full 3D hydrodynamical models on the interpretation of observable properties of late-type giants, briefly mentioning problems and challenges which need to be solved for bringing these models to a routine use within the astronomical community in 5-10 years from now. ",
        "introduction": "Convection plays an important role in governing the interior structure and evolution of late-type giants (i.e., stars on the red and asymptotic giant branches, RGB/AGB). Besides of aiding the energy transport from stellar interiors to the outer layers, convection is important in delivering heavy elements from the nuclear burning layers to the outer atmosphere. Since convective mixing changes the chemical composition both in the outer atmosphere and stellar interiors, it eventually alters also the stellar structure because of changes in atomic and molecular opacities. This affects observable properties of a star, lifetimes in different evolutionary stages, and so forth. Obviously, proper understanding of convection is of fundamental importance for building realistic evolutionary models, which form the basis of our understanding of individual stars and stellar populations.  Convection in current theoretical models is treated in a rather simplistic way, typically within a framework of mixing length theory (MLT). Inevitably, this has a number of drawbacks. For instance, the efficiency of convection within the framework of MLT is scaled by a-priori unknown mixing-length parameter, $\\alpha_{\\rm MLT}$ (defined as the ratio of the mixing length to the pressure scale height), usually calibrated with Solar models. Almost certainly, $\\alpha_{\\rm MLT}$ needs not to be the same in main-sequence stars, subgiants, giants and supergiants, as is assumed in current evolutionary models (see, e.g., \\cite[Freytag \\& Salaris 1999]{FS99} for more details on this issue). Not surprisingly, there have been many attempts during the last few decades to improve the treatment of convection in stellar models (for example, implementing the concept of convective overshooting, which allows for convection to penetrate beyond the classical boundaries of a convective layer). While such efforts are incremental steps towards a more realistic modeling of convection in stellar interiors, a fundamental breakthrough in this area is likely to go beyond the classical stationary 1D modeling. Obviously, this may be possible with full 3D hydrodynamical models, as they account for time-dependent and three-dimensional character of convection from first principles, providing a degree of realism in the treatment of non-stationary phenomena (and convection in particular) that is beyond reach with classical approach.   \\begin{figure} \\centering \\includegraphics[width=6cm]{kucinskas1_fig01.eps} \\caption{Snapshot of the emergent white light intensity during the temporal evolution of a hydrodynamical red giant model. Note the spatial scales of the granulation pattern. The relative rms intensity contrast of the granulation pattern is 22.5\\,\\% at this particular instant in time.}\\label{fig:3Dsurface} \\end{figure} ",
        "conclusions": "Clearly, our 3D model of a prototypical late-type giant predicts considerably different thermal structures from those inferred with the classical 1D model atmospheres. Spectral properties of the 3D model are rather sensitive to convection too, which results in significant differences between photometric colors calculated with the 3D and 1D model atmospheres. All this should be properly taken into account both with evolutionary models and model atmospheres of late-type giants. At the same time, there is a number of issues that will have to be tackled to improve the 3D models, just to name a few:  \\begin{itemize}  \\item Implement a possibility of direct spectral synthesis with 3D stellar atmosphere models, preferably in non-LTE;  \\item Improve spectral line databases, atomic and (especially!) molecular opacities;  \\item Investigate the properties of stars in the regions of HR diagram that were poorly covered with 3D models up to now, at different metallicities (especially RGB/AGB stars).  \\end{itemize}  While 3D hydrodynamical models are still computationally expensive today, grids of synthetic spectra calculated in 3D approach may be available in 5-10 years from now (see \\cite[Ludwig \\& Ku\\v{c}inskas 2005]{LK05} for a discussion). No doubt, this will open new possibilities for improving models of stellar evolution (especially those on RGB/AGB), and will alow to study a variety of new phenomena that are beyond reach with classical 1D models.   \\begin{figure} \\centering \\includegraphics[width=7cm]{kucinskas1_fig03.eps} \\caption{Influence of surface granulation on the broad-band photometric colors of red giant, as reflected by magnitude differences in various band-passes (indicated on the top of the panel) between the predictions of 3D hydrodynamical and classical 1D model atmospheres. The different lines depict various approximations employed in the treatment of the 3D radiative transfer (see \\cite[Ku\\v{c}inskas \\etal\\ 2005]{K05} for details).}\\label{fig:3Dcolors} \\end{figure}"
    },
    "0512/astro-ph0512486_arXiv.txt": {
        "abstract": "X-ray absorption spectroscopy provides a potentially powerful tool in determining the metal abundances in various phases of the interstellar medium (ISM). We present a case study of the sight line toward \\source\\ (Galactic coordinates $l, b=2^\\circ.79, -7^\\circ.91$ and distance = 7.6 kpc), based on {\\sl Chandra} Grating observations. The detection of \\oi, \\oii, \\oiii, \\ovii, \\oviii, and \\neix\\ K$\\alpha$ absorption lines allows us to measure the atomic column densities of the neutral, warm ionized, and hot phases of the ISM through much of the Galactic disk. The hot phase of the ISM accounts for about 6\\% of the total oxygen column density $\\sim8\\times 10^{17} {\\rm~cm^{-2}}$ along the sight line, with the remainder about evenly divided between the neutral and warm ionized phases. By comparing these measurements with the 21 cm hydrogen emission and with the pulsar dispersion measure along the same sight line, we estimate the mean oxygen abundances in the neutral and total ionized phases as 0.3(0.2, 0.6) and 2.2(1.1, 3.5) in units of \\citet{and89} solar value (90\\% confidence intervals). This significant oxygen abundance difference is apparently a result of molecule/dust grain destruction and recent metal enrichment in the warm ionized and hot phases. We also measure the column density of neon from its absorption edge and obtain the Ne/O ratio of the neutral plus warm ionized gas as 2.1(1.3, 3.5) solar. Accounting for the expected oxygen contained in molecules and dust grains would reduce the Ne/O ratio by a factor of $\\sim 1.5$. From a joint-analysis of the \\ovii, \\oviii, and \\neix\\ lines, we obtain the Ne/O abundance ratio of the hot phase as 1.4(0.9, 2.1) solar, which is not sensitive to the exact temperature distribution assumed in the absorption line modeling. These comparable ISM Ne/O ratios for the hot and cooler gas are thus considerably less than the value ($2.85 \\pm 0.07$; 1$\\sigma$) recently inferred from corona emission of solar-like stars (Drake \\& Testa 2005). ",
        "introduction": "}  The measurement of the ISM metal abundances plays a key role in our understanding of the universe. Although metals were all produced in stars, the ISM is the depository of stellar feedback (e.g., via supernovae). The metal abundances of the ISM thus give a unique measure of the integrated stellar feedback. The ISM is also a reservoir from which stars formed; the similarity or dissimilarity of the abundances in the ISM and stars  provides important constraints on the physical processes involved in star formation and evolution (e.g., metal diffusion and settling). Of course, the abundances are also important in the study of various physical and chemical processes of the ISM itself (e.g., heating and cooling as well as dust grain formation). Furthermore, the interpretation of existing observational data often relies heavily on the assumption of the abundances in various phases of the ISM. Examples are the modeling of hot plasma spectra and the correction for ISM absorption to infer the intrinsic spectra of X-ray sources. Therefore, the measurement of the abundances represents a major task of astrophysics.  But a complete accounting of the metal abundances in the ISM has been a challenge. The ISM is a heterogeneous ensemble of various forms (atomic, molecular, and solid dust grain) and various states (phases): cold ($\\lsim 100$ K), warm ($\\sim8000$ K; including warm neutral and warm photo-ionized), and hot ($\\sim10^{6}$ K; collisionally ionized) (e.g., Ferri\\'ere 2001 and references therein). For easy of reference in the present work, we will not distinguish the cold and warm neutral phases and will call their combination as the ``cold'' phase. The ``warm'' phase will, in stead, refer only to the above warm ionized phase; the total ``ionized'' phase therefore refers to a combination of warm and hot phases. Also these phases are used for the {\\sl atomic} gas form only, whereas molecule and solid dust grains are denoted as the ``compound'' form without any phase separation. The abundances of individual elements could be significantly different in these forms and phases, because of the varying level of depletion and molecule/grain destruction as well as recent chemical enrichment from stellar feedback. The traditional methods for abundance measurement in the optical and ultraviolet (UV) are sensitive mainly to the cold and warm phases. The reliability of such measurements strongly depends on how well the physical and ionization conditions are modeled (e.g., Savage \\& Sembach 1996 and  references therein).  Stellar abundance measurements also suffer large uncertainties, especially for solar-like  stars. Recently, solar abundances of light elements such as C, N, O and Ne have been revised downward by 25-45 per cent \\citep{asp05} from the values of Anders \\& Grevesse (1989; AG89 hereafter). The revised values reasonably agree with UV and X-ray measurements of the O abundance in the cold and warm ISM phases (e.g., Sofia \\& Meyer 2001; Takei \\etal 2002; Andr\\'e \\etal 2003; this work). However, this revision has broken the abundance accordance with helioseismological measurements and theoretical solar model predictions (e.g., Bahcall et al. 2005a). It is argued that the Ne abundance in the Sun is poorly determined and that if the Ne/O ratio is in fact substantially larger (by a factor of $\\gtrsim 2.5$) the models can then be brought back into agreement with helioseismological measurements  (e.g., Bahcall et al. 2005b). Based on the spectroscopy of the \\oviii, \\neix, and \\nex\\ emission lines in the \\chandra HETG-ACIS spectra of nearby solar-like stars, \\citet{drake05} estimate that the averaged Ne/O number ratio is $0.41\\pm0.01$ (1 $\\sigma$), or $2.85\\pm0.07$ of the AG89 value. However, two latest independent studies suggest no such high Ne/O ratio at the solar surface \\citep{sch05, young05}.  X-ray absorption spectroscopy can, in principle, be the ideal tool to measure the metal abundances of the ISM, in essentially all forms and phases. The X-ray bandpass contains almost all the K- and L-transitions of different charge states of the abundant elements from carbon to iron. X-ray spectroscopy, less affected by extinction, also probes larger column densities than is possible in the optical and UV, and is thus especially useful for measuring the general ISM through much of the Galactic disk. With the still limited spectral resolution and sensitivity of existing X-ray telescopes, however, only the most abundant elements (like oxygen and neon) can produce significant absorption features in the spectra of bright background sources such as AGNs and X-ray binaries. With the grating instruments aboard \\chandra and {\\sl XMM-Newton}, several groups have indeed studied the abundances of X-ray-absorbing gas, based chiefly on measurements of absorption edges, which are contributed by the cold and warm phases in both atomic and compound forms \\citep{pae01, sch02b, tak02, jue03, jue05, cunn04, ued05}. Most of these studies could only get relative metal abundances (e.g., Ne/O). But \\citet{tak02} and \\citet{cunn04} were also able to estimate the absolute abundance (e.g., O/H) for the sight lines to Cyg X--2 and X Persei, by using emission line data (e.g, 21 cm) and UV absorption lines. Some success has been achieved even in distinguishing the contributions from the atomic and compound forms through calibrating the still uncertain wavelengths and cross-sections of the absorption edges \\citep{tak02}.  More recently, Yao \\& Wang (2005; hereafter Paper I) and Wang \\etal (2005; hereafter Paper II) have demonstrated the possibility of measuring the abundances of the hot ISM phase. By jointly fitting various highly-ionized species, significantly detected or not, one can simultaneously determine multiple parameters of the modeled absorbers such as the gas temperature, velocity dispersion, and element abundance ratio. Even with the limited counting statistics of the LMC X-3 observation, \\citet{wang05} are able to obtain a meaningful lower limit to the Ne/O ratio, which is about the solar value of AG89.  We also note that the detection of \\ion{O}{1}, \\ion{O}{2}, and \\ion{O}{3} absorption lines allow for direct measurements of the cold and warm oxygen column densities. These lines have large oscillator strengths relative to the expected velocity dispersions of the ISM phases; there is little flat part in the curve-of-growth (CoG) of such a line.  Therefore, one can reliably measure the column densities in the typical square-root ranges of the CoG.  Here we capitalize on the high-quality detections of multiple absorption lines in the \\chandra\\ spectra of \\source\\ (Futamoto \\etal 2004; Paper I) to estimate the abundances in the multi-phase ISM of the Galaxy. Throughout this paper, we use the AG89 solar abundances as a convenient reference; the number ratios of O/H and Ne/O are 8.5$\\times10^{-4}$ and 0.144; in comparison, the recently revised solar values are $4.6\\times10^{-4}$ and 0.151, respectively (Asplund, Grevesse, \\& Sauval 2005; further discussion in \\S~5). We also assume a collisional ionization equilibrium (CIE) in the hot intervening X-ray-absorbing gas; this should be quite a good approximation for the oxygen and neon in typical hot ISM conditions \\citep{sut93}. The quoted parameter errors are all at 90\\% confidence levels unless otherwise specified. ",
        "conclusions": "\\end{deluxetable}  For ionized gas, we can use the free electron column density $N_e$ from the pulsar DM (\\S 2). The oxygen column densities in the warm and hot phases are $N$(\\ion{O}{2}+\\ion{O}{3}) and $N$(OVII+OVIII), neglecting \\oiv\\ - \\ovi\\ and \\oix, which are only important in thermally very unstable ``intermediate'' temperatures ($T\\lsim10^{5.5}$ K) and in probably rare low-density regions with temperatures $\\gtrsim 10^{6.5}$ K. We can then define the oxygen abundances in the two phases as $({\\rm O/H})_w = N$(\\ion{O}{2}+\\ion{O}{3})/$\\eta (1-\\xi) N_e$ and $({\\rm O/H})_h = N($\\ion{O}{7}+\\ion{O}{8}$)/\\xi \\eta N_e$, where $\\xi$ is  the hot phase fraction  of the electrons and $\\eta = 0.84$ accounts for the contribution from helium. Letting $({\\rm O/H})_h = \\alpha$ (O/H)$_w$, we have $\\xi = 1/(1+\\alpha r_N$), where $r_N = N($\\ion{O}{2}+\\ion{O}{3}$)/N($\\ion{O}{7}+\\ion{O}{8}$) \\sim 8$. We expect that dust grain destruction and chemical enrichment occurs primarily in the hot phase; i.e., the metal abundance in the hot phase should be comparable to, or higher than, that in the warm phase. Therefore, $\\alpha \\gtrsim 1$; specifically, $({\\rm O/H})_w = 1.8(0.8, 3.1)\\times10^{-3}$ for $\\alpha =1$ and $1.6(0.7,2.8) \\times10^{-3}$ for $\\alpha = \\infty$. We thus adopt 1.7(0.7, 3.1)$\\times10^{-3}$ [or 2.0(0.8, 3.6) solar] to account for this small $\\alpha$-dependent uncertainty. Systematically,  however, this oxygen abundance could be an overestimate, because of the potential contamination to the \\ion{O}{2} line by the uncertain compound oxygen line as mentioned in \\S~3. For instance, if the compound oxygen column density is $\\sim0.5N{\\rm _{OI}}$, as suggested by \\citet{tak02}, then $({\\rm O/H})_w \\sim 9.2\\times10^{-4}$ and $6.7\\times10^{-4}$ for $\\alpha =1$ and $\\infty$, respectively. On the other hand, $({\\rm O/H})_h \\propto \\alpha$ cannot be determined without knowing $\\alpha$. Nevertheless, we may estimate the mean oxygen abundance in total ionized gas: [$N({\\rm O{\\small II}+O{\\small III}})$ + $N({\\rm O{\\small VII}+O{\\small VIII}})$]/$\\eta N_e = 2.2(1.1, 3.5)$ solar, or $\\sim1.1$ solar if the potential contamination from the compound oxygen is accounted for.   }  We have used the \\chandra HETG and LETG observations of the LMXB \\source\\ as a test case to explore the potential for metal abundance measurement  in different atomic phases of the ISM. We have concentrated on the measurements of atomic oxygen and neon, as summarized in Table~\\ref{tab:summary}, and on comparison with previous X-ray studies. Our main results and conclusions are as follows:  1. We have separately measured the column densities of \\oi, \\oii, and \\oiii\\ from their heavily saturated K$\\alpha$ lines. These measurements are only weakly dependent on the exact velocity dispersion assumed. A comparison of the \\oi\\ column density and the 21 cm hydrogen emission in the field gives a cold (neutral) oxygen abundance of 0.3(0.2, 0.6) solar (AG89). The ratio of our measured \\oii\\ plus \\oiii~column density to the pulsar DM along the same sight line further gives an estimate of the warm oxygen abundance as 2.0(0.8, 3.6)  solar.  2. We have constrained the neon column density from its absorption edge, giving an  abundance of 1.2(1.0, 1.5) solar and a Ne/O ratio of 2.1(1.3, 3.5) solar for the cold plus warm gas. The Ne/O ratio with the inclusion of the compound ISM is likely to be a factor of $\\sim 1.5$ lower \\citep{tak02}.  3. We have also measured the oxygen column density and Ne/O ratio in the hot ISM, based on a joint-analysis of the detected \\ovii~K$\\alpha$, \\oviii~K$\\alpha$, and \\neix~K$\\alpha$ absorption lines, together with the non-detection of the \\ovii~K$\\beta$ line. Assuming that \\ovii~K$\\alpha$ and \\oviii~K$\\alpha$ trace all hot gas (i.e., $T \\sim 10^{5.5-6.5}$ K), we estimate that the hot phase accounts for about 6\\% of the total oxygen column density along the sight line and obtain a Ne/O ratio of 1.4(0.9, 2.1)  solar, which is insensitive to the exact temperature distribution assumed.  4. Our abundance estimates for the atomic phases, together with complementary X-ray spectroscopic studies of the total abundances, may have strong implications for understanding various chemical enrichment and depletion processes both in the ISM and during star formation. There is evidence for an enhanced enrichment in the hot gas, especially in comparison with the cold phase. The Ne/O ratios obtained in the ISM are significantly smaller than the value indicated in the recent emission line measurement of solar-like stars \\citep{drake05}.  The existing X-ray measurements of the ISM abundances (including the present work) are still preliminary, subject to various systematic uncertainties, both theoretical and observational, which are yet difficult to fully quantify. Nevertheless, the outcome of this study demonstrates the unique potential for a comprehensive characterization of the metal abundances in various ISM forms and phases."
    },
    "0512/astro-ph0512603_arXiv.txt": {
        "abstract": "The neutron capture cross section of the unstable nucleus $^{186}$Re is studied by investigating the inverse photodisintegration reaction $^{187}$Re($\\gamma$,n). The special interest of the {\\it s}-process branching point $^{186}$Re is related to the question of possible {\\it s}-process contributions to the abundance of the {\\it r}-process chronometer nucleus $^{187}$Re. We use the photoactivation technique to measure photodisintegration rates. Our experimental results are in good agreement with two different statistical model calculations. Although the cross sections predicted by both models for the inverse reaction $^{186}$Re(n,$\\gamma$) is too low to remove the overproduction of $^{186}$Os; the two predicted neutron-capture cross sections differ by a factor of $2.4$; this calls for future theoretical study. ",
        "introduction": "Introduction}  Almost all elements above mass $A \\approx 60$ can be produced in neutron capture reactions\\,\\cite{burb57}. Two different neutron induced processes are necessary to explain the abundance distribution of heavy elements. The first one is the slow neutron capture process ({\\it s}-process). The neutron densities are of the order of $n_{\\rm n} \\approx 10^8\\,{\\rm cm}^{-3}$ and the time scale $\\tau_{\\rm n}$ between two subsequent neutron capture reactions is typically of the order of years. The {\\it s}-process path propagates along the valley of stability. Whenever an unstable nucleus with a mean lifetime $\\tau \\ll \\tau_{\\rm n}$ is reached, this nucleus $\\beta$-decays. If $\\tau \\approx \\tau_{\\rm n}$, a branching occurs and the {\\it s}-process path splits. Thus, nuclei with $\\tau \\approx \\tau_{\\rm n}$ are called branching points of the {\\it s}-process. The second process is the rapid neutron capture process ({\\it r}-process). High neutron densities ($n_{\\rm n} \\gg 10^{20} \\,{\\rm cm}^{-3}$) lead to the production of very neutron rich nuclei up to 20 mass units away from stable nuclei. During freeze out, these  nuclei $\\beta$-decay back to the valley of stability.  There are at least two scenarios known where the {\\it s}-process takes place. It occurs during helium burning in red giant stars and during helium shell flashes in low mass asymptotic giant branch (AGB) stars\\,\\cite{Kaep90, Arla99}. The former scenario is mainly responsible for the production of elements between iron and yttrium. The latter, for the production of elements between zirconium and bismuth. For a detailed discussion see e.g.\\,\\cite{Wall97}. In the following we will focus on the mass region $A \\approx 185$ and, hence, restrict our discussion to the so-called main component of the {\\it s}-process.  \\begin{figure}[ht] \\centering\\includegraphics[angle=0,keepaspectratio,width=8.5cm]{mueller_fig1.eps} \\caption{\\label{fig:fig0} The {\\it s}-process path in the W-Re-Os mass region. Unstable nuclei are marked by dashed boxes (except $^{187}$Re). The indicated values are laboratory half-lives. However, the half-life of $^{187}$Re decreases by 10 orders of magnitude at typical {\\it s}-process temperatures of $T=3 \\times 10^8$\\,K and $^{187}$Os becomes unstable\\,\\cite{Bosc96, Yoko83, Taka83}.} \\end{figure}  Due to its very long half-life ($t_{1/2}=5 \\cdot 10^{10}\\,{\\rm a}$) the nucleus $^{187}$Re can be used as a {\\it r}-process chronometer\\,\\cite{Clay64, Brow76}. The ratio $N(^{187}{\\rm Re}) / N_{\\rm c}(^{187}{\\rm Os})$ is related to the starting point of the {\\it r}-process in our galaxy and, hence, to its age. $N$ denotes the total and $N_{\\rm c}$ the cosmoradiogenic part of the abundance stemming from the decay of $^{187}$Re. To extract the cosmoradiogenic part of the $^{187}$Os abundance one has to subtract the {\\it s}-process abundance $N_{\\rm s}$ from the total abundance $N$. In Fig.\\,\\ref{fig:fig0} the {\\it s}-process flow through the W-Re-Os isotopes is shown. The {\\it s}-process abundance of $^{187}$Os can be derived from the abundance of the neighboring {\\it s}-only nucleus $^{186}$Os via the local approximation\\,\\cite{Clay61}: \\begin{equation} N_{\\rm s}(^{187}{\\rm Os})/N_{\\rm s}(^{186}{\\rm Os}) \\approx F \\,\\bar{\\sigma}_{\\rm n}(^{186}{\\rm Os})/\\bar{\\sigma}_{\\rm n}(^{187}{\\rm Os}), \\label{eq01a} \\end{equation} where $\\bar{\\sigma}_{\\rm n}$ are the Maxwellian-averaged radiative neutron capture cross sections (MACS) from the ground state, and $F$ accounts for the correction of the cross section due to neutron capture on thermally excited states in $^{187}\\rm{Os}$, in particular on the first excited state at $9.75$\\,keV. This correction factor was first calculated in\\,\\cite{Woos79} (see\\,\\cite{Arno84} for discussion). The neutron capture cross sections of $^{186}{\\rm Os}$ and $^{187}{\\rm Os}$ were measured by Browne \\& Berman\\,\\cite{Brow76}, Browne, Lamaze \\& Schroder\\,\\cite{Brow76b}, by Browne \\& Berman\\,\\cite{Brow81} and Winters \\& Macklin\\,\\cite{Wint82}, resulting in an uncertainty of about 20\\% for the ratio $R=\\bar{\\sigma}_{\\rm n}(^{186}{\\rm Os})/\\bar{\\sigma}_{\\rm n}(^{187}{\\rm Os})$. Recently, these cross sections were measured by the n\\_TOF collaboration\\,\\cite{Mosc04, Kaep05}.  The use of the Re/Os clock is not free of problems. First of all, the half-life of $^{187}{\\rm Re}$ strongly depends on temperature and $^{187}{\\rm Os}$ becomes unstable under stellar conditions\\,\\cite{Bosc96, Yoko83, Taka83}. Thus, it is necessary to use chemical evolution models of the galaxy\\,\\cite{Yoko83} to include irradiation effects on the abundance ratio $R$. The two branchings at $^{185}$W and $^{186}$Re also affect the {\\it s}-process abundances in this region. Finally, the $N_s \\bar{\\sigma}_{\\rm n}$ correlation Eq.\\,(\\ref{eq01a}) for the two {\\it s}-only isotopes $^{186,187}{\\rm Os}$ is not fulfilled. This can be caused by two facts. Either the branchings are not correctly modeled or the capture cross sections are strongly affected by stellar conditions. The branching at $^{185}$W has already been studied and an overproduction of $^{186}{\\rm Os}$ was reported due to the new experimental value\\,\\cite{Sonn03a}. Thus the radiative neutron capture cross section of $^{186}$Re is the only relevant cross section in this mass region which is not known experimentally yet.  In this paper we study the radiative neutron capture cross section of the branching point nucleus $^{186}$Re using an indirect method. The unstable nucleus $^{186}$Re decays via $\\beta^-$-decay to $^{186}$Os or via electron capture to $^{186}$W with a half-life of $t_{1/2}=3.7\\,{\\rm d}$. Due to the fact that neutron capture experiments with such short-lived targets are nearly impossible, we choose the inverse reaction $^{187}$Re($\\gamma$,n)$^{186}$Re for our investigation. After neutron emission the $^{186}$Re nucleus is in the ground state or some excited state and the measured cross section is a sum over several channels. With maximum excitation energies just above the neutron threshold, only the lowest states can be reached, e.g., the first excited state in $^{186}$Re at $59\\,\\rm{\\,keV}$. On the other hand, these low lying states in the $^{186}$Re nucleus are also thermally populated - however not in the same proportions - under {\\it s}-process conditions and contribute to the neutron capture cross section. Thus, the $^{187}$Re($\\gamma$,n)$^{186}$Re cross section and the cross section of the inverse reaction $^{186}$Re(n,$\\gamma$)$^{187}$Re are related via the principle of detailed balance.  In section \\ref{para:par2} we describe our experimental setup. Section \\ref{para:par3} explains the analysis of our data and the results for the $^{187}$Re($\\gamma$,n) cross section are presented. The results are compared to calculations using the NON-SMOKER\\,\\cite{Raus00a, Raus04} and MOST\\,\\cite{Gori02, Goriely} codes. Both computer codes are based on the statistical Hauser-Feshbach model but use different input parameters. In section \\ref{para:par3b} the implications on the Re/Os clock are briefly discussed. ",
        "conclusions": ""
    },
    "0512/gr-qc0512158_arXiv.txt": {
        "abstract": "In this work we study to which extent the knowledge of spatial topology may place constraints on the parameters of the generalized Chaplygin gas (GCG) model for unification of dark energy and dark matter. By using both the Poincar\\'e dodecahedral  and binary octahedral spaces as the observable spatial topologies, we examine the current type Ia supernovae (SNe Ia) constraints on the GCG model parameters. We show that the knowledge of spatial topology does provide  additional constraints on the $A_s$ parameter of the GCG model but does not lift the degeneracy of the $\\alpha$ parameter. ",
        "introduction": "Recently, the generalized Chaplygin gas (GCG) model \\cite{Kamenshchik:2001cp,Bilic:2001cg,Bento:2002ps} has attracted considerable attention given its potential to account for the observed accelerated expansion of the Universe \\cite{observations}, and to describe in a simple scheme, both the negative pressure dark energy component as well as the pressureless dark matter component. In terms of the critical density, the contribution of each component is about two thirds for dark energy and one third for dark matter \\cite{Bahcall}.  In the GCG proposal, the dark components are described through a perfect fluid of density $ \\rho_{ch}$ and pressure $p_{ch}$ with an exotic equation of state \\begin{align} p_{ch} = - {A \\over \\rho_{ch}^\\alpha}~, \\label{rhoGCG} \\end{align} where $A$ and $\\alpha$ are positive constants.  For $\\alpha=1$, the equation of state is reduced to the Chaplygin gas scenario \\cite{Kamenshchik:2001cp}. The striking feature of this model is that it allows for an unification of dark energy and dark matter \\cite{Bilic:2001cg,Bento:2002ps}.  The parameters of the GCG or indeed any dark energy model are known to be affected by the spatial geometry the Universe. Physicists describe the Universe as a manifold, which is characterized by its geometry and its topology. Two fundamental questions regarding the nature of the Universe concern the geometry and topology of the $3$--dimensional space. Geometry is a local feature related with the intrinsic curvature of the $3$--dimensional space and can be tested by studies of the cosmic microwave background radiation (CMBR) such as the Wilkinson Microwave Anisotropy Probe (WMAP). Topology is a global property that characterizes its shape and size. Geometry constrains but does not fix the topology of the spatial sections. In a locally spatially homogeneous and isotropic universe the topology of its spatial section dictates its geometry. Within the framework of the standard Friedmann--Lema\\^{\\i}tre--Robertson% --Walker (FLRW) cosmology, the universe is modeled by a space-time manifold $\\mathcal{M}_4$ which is decomposed into $\\mathcal{M}_4 = \\mathbb{R} \\times M_3$ and endowed with a locally (spatially) homogeneous and isotropic  metric \\begin{equation} \\label{RWmetric} ds^2 = -dt^2 + a^2 (t) \\left [ d \\chi^2 + f^2(\\chi) (d\\theta^2 + \\sin^2 \\theta  d\\phi^2) \\right ] \\;, \\end{equation} where $f(\\chi)=(\\chi\\,$, $\\sin\\chi$, or $\\sinh\\chi)$ depending on the sign of the constant spatial curvature ($k=0,1,-1$).   The $3$--dimensional space where we live in is usually taken to be one of the following simply-connected spaces: Euclidean $\\mathbb{R}^3$, spherical $\\mathbb{S}^3$, or  hyperbolic space $\\mathbb{H}^3$. However, given that the connectedness of the spatial sections $M_3$ has not been determined by cosmological observations, and since geometry does not fix the topology, our $3$--dimensional space may equally well be one of the possible multiply connected quotient manifolds $M_3 = \\widetilde{M}/\\Gamma$, where $\\Gamma$ is a fixed point-free group of isometries of the covering space $\\widetilde{M}=(\\mathbb{R}^{3},\\mathbb{S}^{3}, \\mathbb{H}^{3})$.   Thus, for instance, for the Euclidean geometry ($k=0$) besides $\\mathbb{R}^{3}$  there are 10 classes of topologically distinct compact $3$--spaces consistent with this geometry, while for both the spherical ($k=1$) and hyperbolic ($k=-1$) geometries there are an infinite number of topologically inequivalent compact manifolds with non-trivial topology that admit these geometries.  Recently, different strategies and methods to probe a non-trivial topology of the spatial sections of the Universe have been devised (see, e.g., the review articles Refs.~\\cite{CosmTopReviews} and also Refs.~\\cite{CCmethods} for details on cosmic crystallographic methods). An immediate observational consequence of a detectable non-trivial topology% \\footnote{The extent to which a non-trivial topology may have been detected was discussed in Refs.~\\cite{TopDetec}.} of the $3$--dimensional space $M_3$ is that the sky will exhibit multiple (topological) images of either cosmic objects or specific spots on the CMBR. The so-called ``circles-in-the-sky\" method, for example, relies on multiple  images of correlated circles in the CMBR maps~\\cite{CSS1998}. In a space with a detectable non-trivial topology, the sphere of last scattering intersects some of its topological images along pairs of circles of equal radii, centered at different points on the last scattering sphere (LSS), with the same distribution % of temperature fluctuations, $\\delta T$. Since the mapping from the last scattering surface to the night sky sphere preserves circles~\\cite{CGMR05}, these pairs of matching circles will be imprinted on the CMBR anisotropy sky maps regardless of the background geometry or detectable topology. As a consequence, to observationally probe a non-trivial topology one should scrutinize the full-sky CMBR maps in order to extract the correlated circles, whose angular radii and relative position of their centers can be used to determine the topology of the Universe. In this way, a non-trivial topology of the space section of the Universe is observable, and can be probed for all locally homogeneous and isotropic geometries. %  In this regard, in a recent work~\\cite{RAMM} in the context of the $\\Lambda$CDM  model, the Poincar\\'e  dodecahedral space was used as the observable spatial topology of the Universe to reanalyze the current type Ia supernovae (SNe Ia) constraints on the density parameters associated with dark matter ($\\Omega_m$) and dark energy ($\\Omega_{\\Lambda}$).  As a result, it has been shown that the knowledge of the Poincar\\'e dodecahedral space topology through the ``circles-in-the-sky\" method gives rise to stringent constraints on the energy density parameters  allowed by the conventional SNe Ia observations, reducing considerably the inherent degeneracies of the current measurements. Given this encouraging result it is natural to assess to what extent this method can be useful for determining the parameters of more complex dark energy models. In this paper, we address these questions by focusing on the constraints that cosmic topology \\footnote{In line with current literature, by topology of the Universe we mean the topology of the spatial section $M_3$.} together with current SNe Ia data pose on the parameters of the GCG model. To this end, we use the Poincar\\'e dodecahedral and the binary octahedral spaces as the topologies of the spatial sections of the Universe% \\footnote{These spatial topologies account for the low value of the CMBR quadrupole and octopole moments measured by the WMAP team, and fit the temperature two-point correlation function, for values of the total density within the reported range ~\\cite{Poincare,Aurich1,Aurich2,WMAP-Spergel}.} to reanalyze current constraints on the parameters of the GCG model, as provided by the so-called \\emph{gold} sample of 157 SNe Ia~\\cite{Riess:2004nr}. ",
        "conclusions": "The so-called ``circles-in-the-sky\" method makes apparent that a non-trivial detectable topology of the spatial section of the Universe can be probed for any locally homogeneous and isotropic universe, with no assumption on the cosmological density parameters.  In this paper we have shown that the knowledge of $\\mathcal{D}$ and $\\mathcal{O}$ spatial topologies does provide some additional constraints on the $A_s$ parameter of the GCG model, even though it does not help in lifting the degeneracy on the $\\alpha$ parameter.  In any case, our results indicate that the introduction of topological considerations into the analysis of the large scale structure of the Universe is an interesting complementary strategy to constrain and eventually characterize the nature of dark energy and dark matter. In the particular case of the GCG, the complexity of the model does not allow for obtaining striking constraints on its parameters as is the case for the $\\Lambda$CDM model.  Finally, the question arises whether topology may play a significant role for other dark energy and modified gravity models, an issue we plan to analyze in a future publication."
    },
    "0512/astro-ph0512435_arXiv.txt": {
        "abstract": "Corrections to the magnitudes of high redshift objects due to intergalactic attenuation are computed using current estimates of the properties of the intergalactic medium. The results of numerical simulations are used to estimate the contributions to resonant scattering from the higher order Lyman transitions. Differences of $0.5-1$ magnitude from the previous estimate of Madau (1995) are found. Intergalactic $k_{\\rm IGM}$-corrections and colours are provided for high redshift starburst galaxies and Type I and Type II QSOs for several filter systems used in current and planned deep optical and infra-red surveys. ",
        "introduction": "\\label{sec:introduction}  Over the past decade, deep optical and infra-red surveys have enabled giant strides to be taken in elucidating the nature and properties of objects that populate the high redshift universe. The band-dropout method has unveiled a population of Lyman break galaxies at $z\\approx3$ (Guhathakurta \\etal 1990; Bithell 1991; Steidel \\& Hamilton 1992, 1993; Steidel, Pettini \\& Hamilton 1995; Steidel \\etal 2003) and higher (Sawicki \\& Thompson 2005). The selection method was successfully applied to the Hubble Deep Field (HDF) (Giavalisco, Steidel \\& Macchetto 1996; Steidel \\etal 1996), broadening the redshift range and volume coverage over previous surveys. Most recently the Ultra Deep Field (UDF) was exploited to discover objects as distant as $z\\gsim6$ (Stanway, Bunker \\& McMahon 2003; Bouwens \\etal 2004; Yan \\& Windhorst 2004; Giavalisco \\etal 2004).  An alternative selection method for identifying high redshift objects relies on combinations of broadband colours to estimate photometric redshifts, for which the most likely redshift is assigned based on predicted spectral energy distributions (Sawicki, Lin \\& Yee 1997; Csabai \\etal 2000).  The modelling of the high redshift objects through population synthesis, applied to a combination of spectroscopic data and broad-band colours, suggests that most of the high redshift objects are star-forming galaxies (Madau \\etal 1996; Metcalfe \\etal 2001; Papovich, Dickinson \\& Ferguson 2001; Pettini\\etal 2001; Shapley \\etal 2003). A few of the Lyman break galaxies contain Active Galactic Nuclei (Steidel \\etal 2002), used to determine the faint end of the QSO luminosity function at $z\\approx3$ (Hunt \\etal 2004).  Parallel to these surveys have been several searches for high redshift Quasi-Stellar Objects (QSOs). The discovery of a few dozen $z>3.6$ QSOs by the Sloan Digital Sky Survey (SDSS) has made possible a new evaluation of the bright end of the QSO luminosity function and its evolution at these high redshifts (Fan \\etal 2001). The results are currently being revised (Richards \\etal 2005) based on the much larger numbers now detected, including over 500 at $z>4$ (Schneider \\etal 2005).  Similar surveys are expected to continue well into the future, including some now in progress, such as the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS\\footnote {{\\rm www.cfht.hawaii.edu/Science/CFHLS/ }}) and the UKIRT Infrared Deep Sky Survey (UKIDSS\\footnote{{\\rm www.ukidss.org }}; Hewett \\etal, in preparation), or planned for telescopes currently under development, such as the Visible \\& Infrared Survey Telescope for Astronomy (VISTA\\footnote{ {\\rm www.vista.ac.uk }}) and the {\\it James Webb Space Telescope} (JWST\\footnote{{\\rm www.ngst.nasa.gov }}).  Crucial to all these analyses is an accurate estimate of the amount of intergalactic attenuation due to intervening absorption systems. Most have relied on the standard work of Madau (1995), whose assessment was based on the then most current understanding of the properties and distributions of intervening systems. Madau (1995) estimated the blanketing due to the resonant (Lyman series) scattering of photons assuming idealised forms for the \\HI column density distribution of the absorbers. The contributions of the different orders in the Lyman series depend on the full line-shape of the absorber, precluding a direct scaling of line-centre optical depths based on pure atomic physics considerations. Instead the broadening of the absorbers must be included, the distribution of which has since been shown to be sensitive to column density (Kirkman \\& Tytler 1997; Kim, Cristiani \\& D'Odorico 2002b). Madau (1995) adopted a constant Doppler parameter for all absorption systems, varying the value to probe the sensitivity of the total amount of attenuation to this variable. A blanketing formalism based on Poisson placement of the absorbers was used to predict the effective optical depths, although in principle small scale clustering of the absorption systems will affect the total amount of blanketing, and such correlations have been detected (Kirkman \\& Tytler 1997; Kim \\etal 2002a).  Since Madau's seminal work, numerical simulations of the intergalactic medium have yielded results matching the measured distributions of the \\Lya flux distributions to an accuracy of a few percent, as well as the line parameters (allowing for extra heating) (Meiksin, Bryan \\& Machacek 2001). Numerical simulations have also reproduced the \\HI column density dependence of the Doppler parameter envelope (Misawa \\etal 2004), as well as correlations in the \\HI flux distribution (Croft \\etal 2002; Meiksin \\& White 2004). The understanding of the mean \\Lya intergalactic optical depths has improved substantially over the past decade (see Meiksin \\& White 2004 for a summary and Kirkman \\etal 2005 for subsequent results). The simulations contain the information necessary to extract the contributions from all higher order transitions to the blanketing. Although simulations do not recover the full numbers of Lyman Limit Systems observed (Gardner \\etal 1997; Meiksin \\& White 2004), an assessment of their numbers over the redshift range $0.3\\lsim z\\lsim4$ has been made by Stengler-Larrea \\etal (1995). These improvements now permit a much more secure determination of the amount of intergalactic attenuation to be made. ",
        "conclusions": "The amount of the attenuation of light from high redshift objects on passing through the IGM has been computed based on recent measurements of the mean transmitted \\Lya flux through the IGM and recent assessments of the numbers of Lyman Limit Systems. The properties of the IGM required to compute the attenuation due to resonant Lyman photon scattering were based on numerical simulations matching the measured properties of the IGM. Differences from the predictions of the model of Madau (1995) are found for $k_{\\rm IGM}$ and colours of $0.5-1$ magnitude for bands containing restframe \\Lya and shorter wavelengths.  Intergalactic $k_{\\rm IGM}$-corrections are provided for filter systems used in current or planned deep optical and infra-red surveys, {\\it viz.} the Sloan $u^\\prime g^\\prime r^\\prime i^\\prime z^\\prime$ system, Steidel $U_n G {\\cal R}$ and the $I$-band, UKIDSS $ZYJHK$, the {\\it HST} $B_{435}$, $V_{606}$, $i_{775}$, $z_{850}$, $J_{110}$ and $H_{160}$ bands, and the {\\it JWST} $F070W$ and $F150W$ bands.  Colours based on the above bands are provided for starbursts of ages 3~Myr and 600~Myr, typical of $z\\approx3$ Lyman Break Galaxies, and Type I and Type II QSOs, over the redshift range $2.5<z<8.5$. The results show which colours are most effective for distinguishing between different objects and the typical values to expect, which may be used for the planning and analysis of current and upcoming deep surveys."
    },
    "0512/astro-ph0512573_arXiv.txt": {
        "abstract": "{Some rapidly variable extra-galactic radio sources show very high brightness temperatures $T_{\\rm B}>10^{12\\,}$K and high degrees of circular polarisation ($\\sim1\\%$). Standard synchrotron models that assume a power-law electron distribution cannot produce such high temperatures and have much lower degrees of intrinsic circular polarisation.} {We examine the synchrotron and inverse Compton radiation from a monoenergetic electron distribution and discuss the constraints placed upon it by radio, optical and hard X-ray/gamma-ray observations.} {The standard expressions of synchrotron theory are used. Observational constraints on the source parameters are found by formulating the results as functions of the source size, Doppler boosting factor, optical depth to synchrotron self-absorption and maximum frequency of synchrotron emission, together with a parameter governing the strength of the inverse Compton radiation.} {The model gives brightness temperatures $T_{\\rm B}\\sim10^{13}$ to $10^{14}\\,$K for moderate ($\\lesssim10$) Doppler boosting factors together with intrinsic degrees of circular polarisation at the percent level. It predicts a spectrum $I_\\nu\\propto \\nu^{1/3}$ between the radio and the infra-red as well as emission in the MeV to GeV range. If the energy density in relativistic particles is comparable to or greater than the magnetic energy density, we show that electrons do not cool within the source, enabling the GHz emission to emerge without absorption and the potentially catastrophic energy losses by inverse Compton scattering to be avoided. Magnetically dominated sources can also fulfil these requirements at the cost of a slightly lower limit on the brightness temperature.} {We suggest that sources such as \\PKSfifteennineteen, \\PKSofourofive\\ and \\Jeighteennineteen\\ can be understood within this scenario without invoking high Doppler boosting factors, coherent emission mechanisms, or the dominance of proton synchrotron radiation.} ",
        "introduction": "The rapidly varying radio flux density observed in several extra-galactic sources implies a very high brightness temperature $\\Tbright$ at the source \\citep{wagnerwitzel95}. In cases such as \\PKSfifteennineteen\\ and \\PKSofourofive, where the variability is identified as interstellar scintillation \\citep{macquartetal00,rickettetal02} realistic models of the scattering screen require $T_{\\rm B}>10^{13}\\,$K. The recently discovered diffractive scintillation in \\Jeighteennineteen\\ requires $T_{\\rm B}>2\\times10^{14}\\,$K \\citep{macquartdebruyn05}. In other sources, the variability can be interpreted as intrinsic, in which case the implied temperature can be even higher.  Such high temperatures are difficult to understand within standard models of synchrotron emission, which assume a power-law distribution of electrons. The energy radiated by inverse Compton scattering in the Thomson regime rises dramatically when the intrinsic brightness temperature exceeds a certain threshold, roughly equal to $3\\times10^{11}\\,$K in the case of a power-law electron distribution, although a flaring source might exceed this limit for a short time \\citep{kellermannpaulinytoth69,slysh92,melrose02}. However, as first pointed out by \\citet{crusius-waetzel91}, the threshold temperature is higher for a monoenergetic electron distribution than for a power-law distribution, since, in the former case, photons of low frequency (compared to the characteristic synchrotron frequency) can escape without absorption by low energy electrons.  High brightness temperature sources frequently display circular polarisation at the~$1\\%$ level or above \\citep{macquart03}. Because this is much higher than the value $m_{\\rm e}c^2/k_{\\rm B}T_{\\rm B}$ conventionally estimated for the intrinsic emission of a power-law electron distribution, propagation effects are the favoured explanation \\citep{jonesodell77a,wardleetal98,macquartmelrose00,beckertfalcke02,ruszkowskibegelman02,broderickblandford03,wardlehoman03}. A monoenergetic electron model, on the other hand, predicts {\\em intrinsic} circular polarisation at the $\\%$ level, obviating the need for a conversion process. Additional predictions are a hard $I_\\nu\\propto\\nu^{1/3}$ synchrotron spectrum extending from the radio at least up to the infra red, and an inverse Compton component in the MeV to GeV range. ",
        "conclusions": "The electron Lorentz factor implied by the above analysis is: \\eqb \\gamma&=& 2.8\\times10^3 \\left({\\xi\\doppler_{10}\\nu_{{\\rm max}14}^{2/3}\\over \\tausync(1+z)}\\right)^{1/5} \\nuGHz^{-1/3} \\label{gammavalue} \\eqe A key ingredient of this model is the absence of electrons of lower Lorentz factor, since these would absorb the GHz emission, leading to a reduction of the brightness temperature. Specifically, we require a quasi-monoenergetic distribution such that $\\diff \\ln n/\\diff\\ln \\gamma>-1/3$ at Lorentz factors lower than that given by Eq.~(\\ref{gammavalue}). Such a distribution is not a natural consequence of, for example, the first-order Fermi process at relativistic shocks \\citep[e.g.,][]{kirk05}. On the other hand, a relativistic thermal distribution, which rises at low energy as $\\gamma^2$ is well-approximated by a monoenergetic distribution of energy roughly equal to the temperature. The addition of a power-law tail to higher energy would not change this conclusion.  The Lorentz factor implied by Eq.~(\\ref{gammavalue}) is higher than the cut-off conventionally assumed when modelling radio sources \\citep[e.g.,][]{gopal-krishnabiermannwiita04}. Nevertheless, scenarios exist which suggest such values. One example is an electron-proton jet with a bulk Lorentz factor $\\Gamma\\sim 10$ which is thermalised at a shock front. If the downstream electron and ion temperatures are equal, the distribution can be approximated as monoenergetic with an electron Lorentz factor of $\\Gamma m_{\\rm p}/m_{\\rm e}$, where $m_{\\rm p}$ and $m_{\\rm e}$ are the proton and electron masses. Another possibility is that the electrons are accelerated at a current sheet in an electron-proton plasma in which the magnetic energy density is comparable to the rest-mass energy density \\citep{kirk04}. Each of these possibilities relies on the composition of the source plasma being electron-proton. Interestingly, so does the relatively high degree of intrinsic circular polarisation given by Eq.~(\\ref{circular}).  Although it is conceivable that continuous re-acceleration prevents the accumulation of low energy electrons, both the current sheet and the shock scenario envisage a finite escape rate of particles from the acceleration or thermalisation region. Escaping particles subsequently cool by synchrotron and inverse Compton emission. Therefore, the model electron distribution is self-consistent only if these particles can be evacuated from the source in a time shorter than the cooling timescale. The ratio of the electron cooling time $\\tcool$ to the light-crossing time of the source can be written as: \\eqb {c \\tcool/ R}&=&{\\eta/\\xi} \\label{crossing} \\eqe where $\\eta$ is the ratio of the energy density in relativistic electrons to that in the magnetic field. Writing $R=0.01\\rmtwo\\,$parsec, we find: \\eqb \\eta&=& {\\gamma n m c^2\\over \\left(B^2/8\\pi\\right)} \\,=\\, 2.9\\left({\\doppler_{10}^{13} \\xi^8\\over(1+z)^{13}\\tausync^3\\nu_{{\\rm max}14}^{8}}\\right)^{1/5} \\nuGHz^{-1}\\rmtwo^{-1}\\sin^2\\theta \\label{etadef} \\eqe Clearly, very small sources tend to be particle dominated and, since $\\xi<1$, they satisfy the self-consistency requirement $c\\tcool/R>1$. However, Eq.~(\\ref{etadef}) shows that $\\eta$ is also quite sensitive to $\\doppler$, $\\xi$ and $\\numax$, so that magnetically dominated sources are by no means ruled out, provided they have $\\xi\\lesssim\\eta$. Since the brightness temperature is proportional to $\\xi^{1/5}$ (Eq.~\\ref{bright1}) it is slightly lower for magnetically dominated sources.  Although lacking a compelling physical justification, the assumption of equipartition, $\\eta=1$, can be used to define an \\lq\\lq equipartition Doppler factor\\rq\\rq. This leads, in the standard model, to a relatively low limit on the brightness temperature $T_{\\rm B, eq}\\lesssim 3\\times10^{10}\\,$K \\citep{singalgopalkrishna85,readhead94}, which has some observational support \\citep{cohenetal03}. In the monoenergetic model, however, this assumption re-introduces a dependence on the source size, but does not substantially constrain the brightness temperature, as can be seen from equations (\\ref{bright1}) and (\\ref{etadef}).  The parameter \\eqb \\zeta&=&{\\gamma h\\nusynch\\over m_{\\rm e}c^2} \\,=\\, 2.3\\times10^{-4}\\left({(1+z)^4\\xi\\over\\tausync\\doppler_{10}^{4}}\\right)^{1/5} \\nu_{{\\rm max}14}^{17/15} \\nuGHz^{-1/3} \\eqe which gives the ratio of the energy of a photon of the characteristic synchrotron frequency to the electron rest-mass, as seen in the rest frame of a relativistic electron is also independent of source size. For $\\zeta\\ll1$, the first inverse Compton scattering takes place in the Thomson regime. In this case, it is consistent to require $\\xi\\lesssim1$ in order to avoid an excessively large energy demand on the source i.e, in order to avoid the Compton Catastrophe. The first generation of inverse Compton photons has a frequency of approximately \\eqb \\nu_1&\\approx& 4.3\\left({\\doppler_{10}^{2}\\xi^2\\over(1+z)^2\\tausync^2}\\right)^{1/5} \\nu_{{\\rm max}14}^{19/15} \\nuGHz^{-2/3}\\,\\textrm{MeV} \\eqe and its flux can be estimated to be \\eqb \\fluxic &\\approx&4.5\\times10^{-6} \\left({(1+z)^3\\xi^3\\tausync^2\\over\\doppler_{10}^{2}}\\right)^{1/5} \\nu_{{\\rm max}14}^{1/15} \\nuGHz^{1/3}\\fluxghz \\eqe where $\\fluxghz$ is the flux observed in the radio at frequency $\\nuGHz\\,$GHz. This estimate of the inverse Compton flux is generally above the detection threshold of instruments on the INTEGRAL satellite, as noted by \\citet{protheroe03}. Subsequent generations of inverse Compton scattered photons are likely to fall into the Klein-Nishina regime, and the maximum photon energy achieved by multiple inverse Compton scattering is ultimately limited by the electron energy, as seen in the observer's frame, which takes the value: \\eqb {\\doppler\\gamma mc^2\\over 1+z}&=& 14 \\left({\\doppler_{10}^6\\xi\\over(1+z)^6\\tausync}\\right)^{1/5} \\nu_{{\\rm max}14}^{2/15} \\nuGHz^{-1/3}\\,\\textrm{GeV} \\eqe  To summarise, synchrotron radiation from a monoenergetic electron distribution reproduces the extremely high brightness temperatures observed in variable extra-galactic radio sources, and explains the observed levels of circular polarisation. Therefore, it does not appear necessary to appeal to coherent mechanisms \\citep{krishnanwiita90,benfordlesch98,begelmanetal05} or to proton synchrotron radiation \\citep{kardashev00} to understand these objects. Testable predictions of the theory are a hard radio to infra-red spectrum and gamma-ray emission in the MeV to GeV range."
    },
    "0512/astro-ph0512059_arXiv.txt": {
        "abstract": "{The sky region containing the soft gamma-ray repeater \\sgr\\ has been observed three times with XMM-Newton in February and September 2004. \\sgr\\ has been detected with an absorbed flux of $\\sim$9$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ (2-10 keV). For a distance of 11 kpc, this corresponds to a luminosity of $\\sim$3$\\times$10$^{33}$ erg s$^{-1}$, the smallest ever observed for a Soft Gamma Repeater and possibly related to the long period of inactivity of this source. The observed flux is smaller than that seen with Chandra in 2001-2003, suggesting that the source was still fading and had not yet reached a steady quiescent level. The spectrum is equally well fit by a steep power law (photon index $\\sim$3.2) or by a blackbody with temperature kT$\\sim$0.8 keV. We also report on the INTEGRAL transient \\igr\\ that lies at $\\sim$10$'$ from \\sgr . It was detected only in September 2004 with a luminosity of $\\sim4\\times10^{33}$ erg s$^{-1}$ (for d=7 kpc), while in February 2004 it was at least a factor 10 fainter. ",
        "introduction": "\\label{sect:intro}   \\sgr\\ is one of the four confirmed Soft Gamma-ray Repeaters (SGRs) that are currently known. According to the widely  accepted magnetar model (Duncan \\& Thompson 1992, Thompson \\& Duncan 1995), these sources are isolated neutron stars in which the high-energy emission is powered by ultra-strong magnetic fields (B$\\sim$10$^{14}$-10$^{15}$~G). The distinctive characteristic of SGRs is the emission, during sporadic periods of activity, of short bursts ($<$ 1 s) of hard X--rays with super-Eddington peak luminosity  L$\\sim$10$^{40}$--10$^{41}$ erg s$^{-1}$. Persistent (i.e. non-bursting) emission is also observed from SGRs in  the soft X--ray range ($<$10 keV), with typical luminosity of $\\sim$10$^{35}$ erg s$^{-1}$. Periodic pulsations at several seconds, reflecting the neutron star rotation, are observed in three SGRs. Occasionally, SGRs also emit energetic (giant) flares with luminosity from $\\sim$10$^{43}$ erg s$^{-1}$ up to 10$^{47}$ erg s$^{-1}$.  For a review of the properties of these sources see Woods \\& Thompson (2004).  It is interesting to study the relation between the properties of the persistent X--ray emission and the level of SGR bursting and flaring activity. In fact the persistent soft X--rays are thought to consist, at least in part, of thermal emission from the neutron star surface, whose properties, e.g. temperature and magnetization, can be influenced by the largely non-thermal phenomena responsible for the bursts.  The 1-10 keV luminosity of the two SGRs which have been more active in recent years, SGR 1806--20 and SGR 1900+14, displayed only moderate (factor $\\sim$2) long term variations around average values of $\\sim5\\times10^{35}$d$^2_{15}$ erg s$^{-1}$ and $\\sim1\\times10^{35}$d$^2_{10}$ erg s$^{-1}$ respectively\\footnote{we indicate with d$_N$ the distance in units of N kpc} (e.g., Mereghetti et al. 2005, Woods et al. 2001). The SGR in the Large Magellanic Cloud, SGR 0526--66, has a similar luminosity of $\\sim10^{36}$d$^2_{55}$ erg s$^{-1}$ (Kulkarni et al. 2003), despite no bursts have been detected from this source since 1984\\footnote{some bursts might have been missed in the time interval 1985--1991 due to the lack of suitable detectors in operation (see Woods \\& Thompson 2004)}. The only difference with respect to the two more active SGRs mentioned above is that its spectrum is much softer, requiring a power law photon index larger than $\\sim$3 (Kulkarni et al. 2003).     \\sgr\\,  was discovered  in 1998, when more than 100 bursts were observed with different satellites (CGRO, Woods et al. 1999; Ulysses, Hurley et al. 1999; Wind, Mazets et al. 1999;  RXTE, Smith et al. 1999; BeppoSAX, Feroci et al. 1998). No other bursts from this source have been reported to date. Its soft X-ray counterpart was identified with BeppoSAX in 1998 at a luminosity level of $\\sim$10$^{35}$d$^2_{11}$  erg s$^{-1}$ (Woods et al. 1999). Subsequent observations, carried out over a time span of five years with BeppoSAX, ASCA and Chandra, showed a monotonic decrease in the luminosity, interpreted  as evidence for cooling of the neutron star surface after the deep crustal heating that occurred during the 1998 period of SGR activity (Kouveliotou et al. 2003). The latest Chandra observation (March 2003) yielded an X-ray  flux consistent with that measured in September 1999, suggesting that the luminosity of \\sgr\\ settled at its ''quiescent'' level of $\\sim$4$\\times$10$^{33}$d$^2_{11}$ erg s$^{-1}$. To further study the luminosity evolution of \\sgr\\ we observed it with XMM-Newton in September 2004. The SGR was also serendipitously detected in two other XMM-Newton observations pointed on the transient IGR J16358--4726 (Patel et al. 2004), which lies at an angular distance of $\\sim$10$'$. We report here also the results of these serendipitous detections, as well as a reanalysis of the BeppoSAX observations. For completeness, we report also the results of the three XMM-Newton observations for IGR J16358--4726. ",
        "conclusions": "The light curve of \\sgr\\ based on data from the different satellites is shown in Fig.~2, where the 2-10 keV flux values correspond to the fits with the same absorption in all the observations (N$_H$=9$\\times10^{22}$ cm$^{-2}$). The long term decrease in luminosity is clear, but, owing to the source spectral variations,  the detailed shape of the decay is different for the observed (upper panel) and unabsorbed (lower panel) flux. Kouveliotou et al. (2003) fitted the decay of the unabsorbed flux using a model involving a deep crustal heating following the 1998 bursting activity and requiring a massive neutron star (M$>$1.5 $M{_\\odot}$). In particular, they pointed out that this model could well explain the plateau between days 400 and 800, but noticed that the March 2003 Chandra observation could not be explained in this framework, suggesting that the source reached a steady low level luminosity. According to our reanalysis of the BeppoSAX data the evidence for a plateau between days 400 and 800 is not so compelling. In fact all the BeppoSAX and ASCA points, before the rapid decline seen with Chandra in September 2001, are well fit by a power law decay, F(t)$\\propto$(t-t$_0$)$^{-\\delta}$. Fixing t$_0$ at the time of the discovery outburst, we obtain $\\delta$=0.6.   If one considers the observed fluxes, the Chandra and XMM-Newton data suggest that \\sgr\\ has continued to fade also after September 2001. There is evidence that the spectrum softened between the two Chandra observations (Kouveliotou et al. 2003). The photon index measured with XMM-Newton is consistent with that of the last Chandra observation but, due to the large uncertainties,  also a further softening cannot be excluded. This apparent fading is not necessarily related to a variation of the source overall luminosity, as clearly indicated by the fluxes corrected for the absorption plotted in the lower panel of Fig.~2.  The XMM-Newton data of September 2004 imply a luminosity $\\sim$3.5$\\times$10$^{33}$d$^2_{11}$ erg s$^{-1}$. This is the lowest luminosity observed from a SGR. The fact that \\sgr\\ has not emitted bursts during the last $\\sim$6 years suggests that a luminosity below 10$^{34}$ erg s$^{-1}$ might be typical of ''quiescent'' SGRs. This simple interpretation is possibly contradicted by the two following considerations. First, the SGR in the Large Magellanic Cloud,  SGR 0526--66,  has a higher luminosity (10$^{36}$ erg s$^{-1}$), but has not shown signs of strong bursting activity in the last 15 years. However, faint bursts, like those recently observed from SGR 1806--20 with INTEGRAL (G\\\"{o}tz et al. 2004) might have passed undetected in SGR 0526--66 due to its larger distance and location in a less frequently monitored sky region. Second, most Anomalous X-ray Pulsars (AXPs, see, e.g. Mereghetti et al. 2002 for a review), which are also generally thought to be magnetars, have nearly steady luminosity larger than 10$^{35}$ erg s$^{-1}$. Although bursts have been observed in three of them (1E 1048.1--5937, Gavriil et al. 2002; 1E 2259$+$586, Kaspi et al. 2003; XTE J1810--197, Woods et al. 2005), there are a few AXPs which have not shown any bursting activity and yet are relatively luminous X-ray sources.  The low luminosity and soft spectrum of \\sgr\\ seen with XMM-Newton are quite similar to the values measured in archival data of the AXP XTE J1810--197 (Gotthelf et al. 2004) obtained before its discovery as a bright transient source in January 2003 (Ibrahim et al. 2004). Based on the currently available data, the two sources seem to behave in a similar way. Further observations of XTE J1810--197 will establish if and how a steady quiescent level is attained.  We finally note that another soft repeater, SGR 1900+14, has possibly been quiescent in the last three years. To our knowledge, the last reported burst activity from this source occurred in November 2002 (Hurley et al. 2002). Therefore it will be interesting to see whether also in SGR 1900+14 the X-ray luminosity will evolve toward a low state similar to that observed for \\sgr ."
    },
    "0512/astro-ph0512084_arXiv.txt": {
        "abstract": "New boron abundances for seven main-sequence B-type stars are determined from HST STIS spectroscopy around the \\ion{B}{3} 2066 \\AA\\ line.  Boron abundances provide a unique and critical test of stellar evolution models that include rotational mixing since boron is destroyed in the surface layers of stars through shallow mixing long before other elements are mixed from the stellar interior through deep mixing.  The stars in this study are all on or near the main-sequence and are members of young Galactic clusters.  They show no evidence of mixing with gas from H-burning layers from their CNO abundances.  Boron abundances range from 12+log(B/H)$\\le$1.0 to 2.2.  The boron abundances are compared to the published values of the stellar nitrogen abundances (all have 12+log(N/H) $\\le$7.8) and to their host cluster ages (4 to 16 Myr) to investigate the predictions from models of massive star evolution with rotational mixing effects.    We find that the variations in boron and nitrogen are generally within the range of the predictions from the stellar evolution models with rotation (where predictions for models with rotation rates from 0 to 450 \\kms\\ and $\\mu$-barriers are examined), especially given their age and mass ranges.  Three stars (out of 34 B-type stars with detailed boron abundance determinations), deviate from the model predictions, including HD\\,36591, HD\\,205021, and HD\\,30836. The first two of these stars have much larger boron depletions than are predicted for their spectroscopic masses and very young ages, even adopting the highest rotation rates from the model predictions.   HD\\,36591 also shows no significant nitrogen enhancement, as uniquely predicted by the rotating stellar evolution models. HD\\,205021, however, has a small nitrogen enrichment which could be explained by stellar rotation or mass transfer since it is in a binary system. The spectroscopic mass for the third star, HD\\,30836, is marginally lower than expected given the rotating model predictions for its age and boron abundance.   This star also has no significant nitrogen enhancement, thus even though it is in a binary system it does not show the nitrogen enrichment expected if it has undergone mass transfer. Therefore, the results from these three stars suggest that rotational mixing could be more efficient than currently modelled at the highest rotation rates. ",
        "introduction": "Rotation is recognized as an important physical component in understanding the evolution of massive stars and yet is a theoretically challenging problem. Rotation affects the lifetimes, chemical yields, stellar evolution tracks, and the properties of supernova and compact remnants (Heger $\\&$ Langer 2000; Maeder $\\&$ Meynet 2000, 2005).   The new rotating stellar evolution models also address long standing problems such as the origin of the B[e] and WNL/Ofpe (slash) stars, the distribution of red to blue supergiants on the HR diagram, and the unseen post main-sequence gap predicted in all standard stellar evolution scenarios.  With rotation, it is also possible to explain the variations in the surface helium, carbon, nitrogen, and oxygen (He and CNO) abundances in OB stars on and near the main sequence (e.g., Gies \\& Lambert 1992, Herrero \\etal\\ 1992, Cunha \\& Lambert 1994, Dennisenkov 1994, Lyubimkov 1996, Lyubimkov \\etal\\ 2004). Several of these observations have been used to constrain the various mixing prescriptions used in the new models of massive star evolution with rotation. Additional observations are now necessary to test the model predictions and provide new constraints for the transport of angular momentum and chemical species in rotating massive stars.  One avenue of observational testing is to determine the helium and CNO abundances in massive stars in young star clusters in the Galaxy and Magellanic Clouds (Evans \\etal\\ 2005).  This tests the metallicity effects predicted by the models.  Another line of research includes the determination of the light element abundances at the surface of massive stars.  This tests the earliest stages of rotational mixing (timescales and mixing efficiencies), {\\it before} hot gas from the stellar interior is observable at the surface (Fliegner \\etal\\ 1996, Venn \\etal\\ 2002).  The abundances of the light elements, lithium, beryllium, and boron (LiBeB) are known to be sensitive to rotational mixing in both low and high mass stars. LiBeB is destroyed on exposure to protons at temperatures too low for H-burning to have occurred ($\\le$6x10$^6$~K), therefore even shallow mixing at the stellar surface induced by rotation can lead to LiBeB depletions.  In low mass stars, variations in the surface abundances of all three elements have been traced to rotational mixing (e.g., Boesgaard $\\&$ Lavery 1986; Pinsonneault 1997). In high mass stars, only boron\\footnote{There are no readily accessible spectral lines of Li in hot stars, and the resonance lines of \\ion{Be}{2} near 3130\\AA\\ are predicted to be very weak in B-type stars, and they occur at the atmospheric cutoff, thus these spectra would be difficult to access from ground based telescopes.   Currently, there are no published observations of the \\ion{Be}{2} lines in B-stars.} has been available.   Spectroscopy of the \\ion{B}{3} feature at $\\lambda$2066 using the International Ultraviolet Explorer (IUE) archived spectra or the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS) or Goddard High Resolution Spectrograph (GHRS) have made boron abundance determinations possible in B-type stars. Significant variations in the boron abundances in hot, massive stars have been observed (e.g., Proffitt \\& Quigley 2001, Venn \\etal\\ 2002).  Boron depletions in B-type stars are predicted to be associated with nitrogen enhancements.  This is because mechanisms that can deplete boron at the surface of a star will also mix the surface with CN-cycled gas, i.e., gas from H-burning layers where the CN-cycle has converted carbon into nitrogen. This is true whether boron is destroyed through rotational mixing (where hot CN-cycled gas from the interior is mixed to the surface) or through binary mass transfer (where CN-cycled gas is deposited on the surface of the star; Wellstein 2001).   Simple initial abundance variations and mass loss from B-type stars are ruled out (see Venn \\etal\\ 2002).   However, there is one {\\it unique} characteristic of the rotational mixing scenario that makes for an interesting and exciting new observational constraint. Boron depletion occurs {\\it before} nitrogen enhancement.  This is because rotational mixing taps the surface boron-free layers first, then subsequent deeper mixing can penetrate layers where H-burning has occured via the CN-cycle, converting carbon into nitrogen.  Thus, the initial phases of rotational mixing in stars is revealed by depletions of boron, followed only {\\it later} by nitrogen-enrichment and carbon-depletion, a different signature from binary mass transfer.  Some B-type stars with boron depletions, but no nitrogen enrichments, were found by Venn \\etal\\ (2002) and intrepreted as concrete and unambiguous evidence for rotational mixing on the main sequence.  However, two stars with masses of 12-13~$M_{\\sun}$ showed uncharacteristically large boron depletions.  These two stars required models with higher masses (20 M$_\\odot$) and the highest rotation rates ($\\sim$450~\\kms) to reproduce the boron depletions (from Heger \\& Langer 2000 models).  Since very few stars are expected to rotate at these high speeds, and because of the mass discrepancy, this suggested that rotational mixing may be even more efficient than currently predicted.  Presently, 32 solar neighborhood (see Table~9 in Venn \\etal\\ 2002) and two stars in the Small Magellanic Cloud (Brooks \\etal\\ 2002) B-type stars have boron abundance determinations.  Of these, only nine have been determined from high quality HST STIS or GHRS spectra of the relatively unblended \\ion{B}{3} 2066 line. In this paper, we present new boron abundances for seven additional B-type stars from HST STIS spectroscopy.   These stars were selected from the group of B-type stars examined with International Ultraviolet Explorer (IUE) that may have low boron abundances (Proffitt \\& Quigley 2001). They are also all in young galactic clusters (for age estimates), and they show no nitrogen enrichments (12+log(N/H) $\\le$7.8). Therefore, the determination of boron abundances in these stars provides a unique observational test for studying the earliest phases of massive star rotational mixing effects. ",
        "conclusions": "Detailed abundances of boron have been determined in a careful selection of seven sharp-lined B-stars in young clusters to test massive star evolution scenarios that include rotational mixing effects.   The seven stars in this analysis show moderate-to-severe depletions of boron (1/4 to 1/50 solar!) {\\it without} significant nitrogen enrichments.   This is a unique prediction of the massive star evolution models that include rotational mixing (Heger $\\&$ Langer 2000), because shallow mixing will destroy surface boron abundances before deep mixing bring CN-cycled gas to the surface. Only three stars deviate from the model predictions, HD\\,30836, HD\\,36591, and HD\\,205021.   In all three cases, their spectroscopic masses are smaller than predicted from the rotating evolution models (i.e., the masses required to explain their low boron abundances and young ages). These three stars appear to indicate that rotational mixing is more efficient than currently modelled at the highest rotation rates. These results are consistent with previously published results (Venn \\etal\\ 2002), but approximately doubles the sample size of stars with both moderate and severe boron depletions."
    },
    "0512/astro-ph0512421_arXiv.txt": {
        "abstract": "Properly apodized pupils can deliver point spread functions (PSFs) free of Airy rings, and are suitable for high dynamical range imaging of extrasolar terrestrial planets (ETPs). To reach this goal, classical pupil apodization (CPA) unfortunately requires most of the light gathered by the telescope to be absorbed, resulting in poor throughput and low angular resolution. Phase-induced amplitude apodization (PIAA) of the telescope pupil \\cite{Guyon2003} combines the advantages of classical pupil apodization (particularly low sensitivity to low order aberrations) with full throughput, no loss of angular resolution and little chromaticity, which makes it, theoretically, an extremely attractive coronagraph for direct imaging of ETPs. The two most challenging aspects of this technique are (1) the difficulty to polish the required optics shapes and (2) diffraction propagation effects which, because of their chromaticity, can decrease the spectral bandwidth of the coronagraph. We show that a properly designed hybrid system combining classical apodization with the PIAA technique can solve both problems simultaneously. For such a system, the optics shapes can be well within today's optics manufacturing capabilities, and the  $10^{-10}$ PSF contrast at $\\approx 1.5 \\lambda/D$ required for efficient imaging of ETPs can be maintained over the whole visible spectrum. This updated design of the PIAA coronagraph  maintains the high performance of the earlier design, since only a small part of the light is lost in the classical apodizer(s). ",
        "introduction": "\\label{sec1}  An optical system capable of extremely high contrast imaging (about $10^{-10}$) at separations  comparable to the telescope's diffraction limit is critical for direct imaging of extrasolar terrestrial planets.   Properly apodized telescope pupils (Nisenson \\& Papaliolios 2001; Kasdin et al. 2003), or designs derived from the classical Lyot coronagraph \\cite{Soummer2003,Kuchner2005} provide the appropriate contrast level. Unfortunately, they suffer from low throughput, ranging from 0.1 to 0.3, and large inner working angles (IWAs), above $3\\lambda/D$. More efficient concepts, capable of near 100\\% throughput and $\\approx\\lambda/D$~IWA exist \\cite{Roddier1997,Baudoz2000,Rouan2000}. They however exhibit a reduced performance for off-axis rays, sufficiently strong to prevent high contrast on nearby partially resolved stars.  A recently proposed alternative to the ``classical'' pupil apodization (refered to as  CPA in this work) is to geometrically remap the entrance pupil of~~the~~telescope~~into~~an~~apodized~~pupil (Guyon  \\begin{landscape} \\begin {figure}[p] \\psfig{figure=f1.eps,width=1.3\\textwidth} \\caption{ Optical layout of a PIAA/CPA hybrid Coronagraph. The system is shown with a focal plane input (a) and output (g), but could also be designed to accept and deliver a collimated beam. Most of the apodization is performed by the 2 aspheric mirrors M1 (c) and M2 (e), which remap the incoming beam into a truncated gaussian-like profile. A second apodization, produced by the classical apodizer (f), removes some of the light in the wings of this profile to produce a spheroidal prolate profile. The opaque focal plane mask (g) efficiently removes the light of the central source, while the rest of the field is fed to a PIAA unit mounted backwards (h) to restore a clean off-axis PSF over a ``wide'' field. In order to minimize unwanted diffraction effects, the apodization profile delivered by the aspheric mirrors is carefully chosen to avoid strong curvature on the M1 mirror. Further mitigation of diffraction effects is obtained by slightly oversizing the entrance beam and apodizing its outer edge (b). Thanks to a constant-curvature extension (d) of the first PIAA mirror, this oversized edge-apodized beam is projected on the second PIAA mirror (e) which therefore acts as the pupil stop in the system.} \\label{fig1} \\end{figure} \\end{landscape}  \\noindent 2003) (this technique is referred to as PIAA, or phase-induced amplitude apodization, in this work). This can be done with two aspheric optics, preferably mirrors: the first aspheric mirror is mostly used to project on the second mirror the desired beam profile, and the second mirror recollimates (or refocuses) the output beam. Mirror shapes can easily be computed by solving a differential equation \\cite{Guyon2003,Traub2003}. Although such a \\hbox{system}~~ corrupts~~~the~~~telescope~~~isoplanaticity \\\\ \\hbox{(the~~unabberated~~field~~of~~view~~is~~only~~about~~} of $5\\lambda/D$ \\footnote{Only the sky related angular scale $\\lambda/D$, measured for the principal ray of the system \\cite{Guyon2005}, is used in this paper.}), a wide field of view can be restored by using the second set of post-coronagraphic PIAA optics \\cite{Guyon2003} which does not affect the coronagraphic performance.  The  PIAA technique combines many of the advantages found separately in other coronagraphs: \\begin{enumerate} \\item{Very high throughput for the planet's light (nearly 100\\%)} \\item{Small inner working angle (slightly larger than $\\lambda/D$).} \\item{Excellent achromaticity if implemented with mirrors (in the geometrical optics approximation).} \\item{Relative insensitivity to pointing errors.} \\end{enumerate} These advantages have been quantified in several studies of the PIAA (Guyon 2003; Traub \\& Vanderbei 2003; Vanderbei \\& Traub 2005). A detailed analysis of a complete PIAA coronagraph (PIAAC) design was recently performed \\cite{Guyon2005}, and the performance of the same design was evaluated for an imaging survey of ETPs with a space telescope \\cite{Martinache2005}. This last study showed that the PIAAC is significantly more efficient than CPAs. A laboratory experiment, performed with lenses, has demonstrated beam apodization and imaging with a PIAA unit \\cite{Galicher2005}.  While these studies showed that the PIAAC is, in theory, very efficient for direct imaging of ETPs, two serious concerns remain unanswered: \\begin{itemize} \\item{{\\bf Optics manufacturability.} In the original PIAA design \\cite{Guyon2003}, the outer edge of the first PIAA mirror is highly curved. This feature, which is essential to obtain the desired apodization, is extremely difficult to polish.} \\item{{\\bf Effects of diffraction propagation.} PIAA units have so far been designed and studied with geometric and Fraunhoffer approximations. As recently shown by Vanderbei (2005), differences between diffraction propagation and geometric/Fraunhoffer optics are not negligible at the $10^{-10}$ contrast level.} \\end{itemize}  In this work, we will address both issues through the study of a PIAA/CPA hybrid system. This new design combines a PIAA unit with a mild ``classical'' apodization of the beam. In \\S\\ref{sec2}, we focus on the optics manufacturability issue, present our hybrid design and explain how it solves this challenge.    In \\S\\ref{sec3} we introduce the diffraction propagation problem in the PIAA apodizer and describe our method of diffraction calculation. The effects of diffraction propagation on the PSF contrast in a poorly designed system and possible solutions of the problem are analyzed in \\S\\ref{sec4}. Lessons learned from \\S\\ref{sec4}  are used to design a much superior hybrid system which is studied in \\S\\ref{sec5}. We give there a broader analysis of the PIAA design tradeoffs and quantify the performance of such systems for direct imaging of ETPs. ",
        "conclusions": "\\label{sec7} Phase-induced amplitude apodization (PIAA) offers full throughput and small IWA, but the required optics shapes are challenging to manufacture and the technique is prone to diffraction-induced chromatic effects. On the other hand, classical apodization coronagraphy is very robust, but suffers from low throughput and large IWA. We have shown in this work that both techniques can be combined in a ``hybrid'' coronagraph design to offer high coronagraphic performance (nearly 90\\% throughput,  $1.5 \\lambda/d$ IWA, low chromaticity) with \"manufacturing-friendly\" optics shapes.  The system presented in \\S\\ref{sec5} achieves $10^{-10}$ contrast at 1.5 $\\lambda/d$ and beyond in a wide spectral band ($d\\lambda/\\lambda \\approx 0.21$) at a small cost in throughput (14\\%) and angular resolution ($7\\%$). Systems with higher throughput can be designed to operate in a smaller bandwidth.  The flexibility of our hybrid design leaves room for further optimization. For example, the roles of the system's two apodizers can be shared to increase throughput. Ultimately, PSF contrast, spectral bandwidth, optics shapes and system throughput would need to be optimized for a particular telescope size and target list.  At the $10^{-10}$ contrast level, small mirror figure errors (wether in OPD or reflectivity) introduce chromatic aberrations in the wavefront \\cite{Shaklan2005} which require the coronagraphic spectral bandwidth to be reduced to 10\\% or less. Our study therefore shows that a PIAA hybrid coronagraph can be designed to not be the dominant source of chromatic aberrations."
    },
    "0512/astro-ph0512617_arXiv.txt": {
        "abstract": " ",
        "introduction": "HH 110 (in L1267, $d=460$ pc \\cite{lopez:Rei91}) and HH 262 (in L1551, $d=140$ pc \\cite{lopez:Graham90}) are two Herbig-Haro (HH) objects that share a peculiar, rather chaotic morphology. In addition, no stellar source powering these jets has been detected at optical or radio wavelengths. Both, previous observations \\cite{lopez:Lopez98} \\cite{lopez:Lopez05} and models \\cite{lopez:Raga02}, suggest that the jet emission reveals an early stage of the interaction between a supersonic outflow and the dense outflow environment. These HHs are thus suitable to search for observational signatures of supersonic outflow/dense environment interaction.  We mapped these HHs with the Integral Field Instrument PMAS (Postdam Multi-Aperture Spectrophotometer) at the 3.5m CAHA telescope, under the PPAK configuration (331 science fibers, of 2\\rlap.''7 each, in an hexagonal grid of $\\sim$ 72'' of diameter). We used the J1200 grating (spectral resolution $\\sim$ 15 km~s$^{-1}$ for H$\\alpha$; wavelength range $\\sim$ 6500-7000 \\AA\\, that includes the emissions from the characteristic red HHs lines: H$\\alpha$, and the [NII] and [SII] doublets). Mosaics from several overlapping pointings (4 for HH 110 and 8 for HH 262) were made in order to cover the whole area of the emission of the HHs. ",
        "conclusions": ""
    },
    "0512/hep-ph0512114_arXiv.txt": {
        "abstract": "\\hskip -5mm  We have re-analyzed the world data on inclusive polarized DIS, in both NLO and LO QCD, including the new HERMES and COMPASS data. The updated NLO polarized densities are given in both the $\\rm \\overline{MS}$ and JET schemes. The impact of the new data on the results is discussed.  \\vskip 1.0cm PACS numbers: 13.60.Hb, 13.88+e, 12.38.-t, 14.20.Dh ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512171_arXiv.txt": {
        "abstract": "037-B327 is of interest because it is both the most luminous and the most highly reddened cluster known in M31. Deep observations with the Advanced Camera for Surveys on the $Hubble$ $Space$ $Telescope$ provide photometric data in the F606W band, and also show that this cluster is crossed by a dust lane. We determined the structural parameters of 037-B327 by fitting the observed surface brightness distribution to a King model with $r_c=0.72\\arcsec(=2.69~\\rm{pc})$, and $r_t=5.87\\arcsec(=21.93~\\rm{pc})$, and a concentration index $c=\\log (r_t/r_c)=0.91$. The surface brightness profile appears to be essentially flat within $0.25\\arcsec$ of the center and shows no signs of core collapse. Although the dust lane affects the photometry, the King model fits the surface brightness profile well except for the regions badly affected by the dust lane. We also calculate the half-light radius $r_h=1.11\\arcsec(=4.15~\\rm{pc})$. Combined with previous photometry, we find that this object falls in the same region of the $M_V$ versus log $R_h$ diagram as do $\\omega$ Centauri, M54 and NGC 2419 in the Milky Way and the massive cluster G1 in M31. All four of these objects have been claimed to be the stripped cores of former dwarf galaxies. This suggests that 037-B327 may also be the stripped core of a former dwarf companion to M31. ",
        "introduction": "It has been speculated that some of the most luminous known globular clusters might be the remnants of tidally stripped dwarf galaxies nuclei \\citep{zinnecker88,freeman93,bassino94}. The study of globular clusters in M31 was initiated by \\citet{hubble32}, who discovered 140 GCs with $m_{pg}\\leq 18$ mag. The continued importance of the study of GCs in this galaxy has been reviewed by \\citet{bh00}. M31 globular cluster B327 (B for `Baade') or Bo037 (Bo for `Bologna', see Battistini 1987), which, in the nomenclature introduced by \\citet{huchra91} will subsequently be referred to as 037-B327. The extremely red color of this object was first noted by \\citet{kronmay60}.  The brightest globular clusters in M31 are more luminous than the giant Galactic cluster $\\omega$ Centauri. Among these are 037-B327 \\citep{bergh68} and G1 \\citep[see details from][]{bk02a}. The latter has been considered as the possible remnant core of a former dwarf galaxy which lost most of its envelope through tidal interactions with M31 \\citep{meylan97,meylan01}. Subsequently \\citet{mackey05} strengthened the \\citet{meylan97} and \\citet{meylan01} conclusion.  In this paper, we have determined the structural parameters of 037-B327 using its deep image obtained with the Advanced Camera for Survey (ACS) on the $Hubble$ $Space$ $Telescope$ $(HST)$. Combined with the previous photometry, we find that this cluster lies in the same region of the log $R_h$ versus $M_V$ diagram as do $\\omega$ Centauri, M54 and NGC 2419 in the Milky Way and G1 in M31. This suggests that 037-B327 may also be the remnant core of a now defunct dwarf companion to the Andromeda galaxy.  \\begin{figure*} \\begin{center} \\centerline{\\includegraphics[angle=0,width=120mm]{f1.eps}} \\caption{The image of GC 037-B327 observed in the F606 (left) and its deconvolved counterpart (right). The central structure is clearly more complex in the deconvolved image. The image size is $7.8\\arcsec\\times8.8\\arcsec$.} \\label{fig1} \\end{center} \\end{figure*} ",
        "conclusions": "In this paper, we determine the structural parameters of 037-B327 that were derived from an F606W image that was obtained with the Advanced Camera for Surveys on the $Hubble$ $Space$ $Telescope$, by fitting between the surface brightness distribution and the King model. Combined with the previous photometry, we find that this object falls in the same region of the $M_V$ versus $R_h$ diagram as do $\\omega$ Centauri, M54 and NGC 2419 in the Milky Way and the massive cluster G1 in M31 on the size (log $R_h$) versus luminosity ($M_V)$ diagram. All four of these objects have been suggested to be the stripped cores of former dwarf galaxies. So, we argue that 037-B327 may also be the core of a former dwarf spheroidal companion to M31. We also compared the images of the F606W and F814W, and did not find any difference in the colors of the brightest incipiently resolved stars, where this term is used in the sense that the image is not clearly resolved into individual stars, but has a mottled or granular appearance, which was employed by \\citet{baade63}."
    },
    "0512/astro-ph0512492_arXiv.txt": {
        "abstract": "$\\beta$ Cet, 31 Com and $\\mu$ Vel represent the main stages through which late-type giants evolve during their lifetime (the Hertzsprung gap (31 Com), the rapid braking zone ($\\mu$ Vel) and the core helium burning ``clump'' phase ($\\beta$ Cet)). An analysis of their high resolution {\\it Chandra} X-ray spectra reveals similar coronal characteristics in terms of both temperature structure and element abundances for the more evolved stars ($\\mu$ Vel and $\\beta$ Cet) with slight differences for the `younger' giant (31 Com). The coronal temperature structure of 31 Com is significantly hotter showing a clear peak while $\\beta$ Cet and $\\mu$ Vel show a plateau. $\\beta$ Cet and $\\mu$ Vel show evidence for a FIP effect in which coronae are depleted in high FIP elements relative to their photospheres by a factor of $\\sim 2$. In contrast, 31~Com is characterized by a lack of FIP effect. In other words, neither depletion nor enhancement relative to stellar photospheric values is found. We conclude that the structural changes during the evolution of late-type giants could be responsible for the observed differences in coronal abundances and temperature structure. In particular, the size of the convection zone coupled with the rotation rate seem obvious choices for playing a key role in determining coronal characteristics. ",
        "introduction": "The solar coronal abundance anomaly commonly known as the ``FIP Effect'' (or First Ionization Potential Effect), in which low first ionization potential (FIP) elements (e.g.; Si, Fe, Mg) appear enhanced by average factors of 3-4, was already { known} by the time the first observational clues to similar abundance anomalies in stellar coronae were uncovered in the 1990's \\citep[e.g.;][]{Feldman92}.  These clues came from low resolution soft X-ray spectra ({\\it GINGA, ASCA, BeppoSAX}), together with moderate resolution extreme ultraviolet (EUV) spectra obtained by the Extreme Ultraviolet Explorer (EUVE). The early stellar studies found evidence for a solar-like FIP Effect in some stars, but for the more active stars the abundances pointed toward metal paucity rather than enhancement \\citep{Drake94,White94}.  The last four years---the beginning of the {\\it Chandra} and {\\it XMM-Newton} era--- have seen early hints of abundance anomalies fleshed out into an interesting array of diverse abundance patterns in which active stars appear to show signs not only of low FIP element depletion, but also of high FIP element enhancements.  Studies of solar-like active stars confirm the suspicions engendered by earlier work \\citep[e.g.;][]{Drake95} that a transition from a metal-depleted to a metal-rich corona occurs as the activity decreases \\citep{Guedel02}; this is now better characterised as a change from an apparently ``inverse FIP effect'' to a FIP effect. \\citet{Audard03} showed a similar transition from an inverse FIP effect to the absence of an obvious FIP effect for RS CVn binary stars with decreasing activity. \\citet{Drake95}, \\citet{Raassen02} and \\citet{Sanz-Forcada04} showed that coronal abundances can also be similar to that of its underlying photosphere. \\citet{Sanz-Forcada04} also presented evidence of possible variation in coronal metal abundance relative to the temperature of the emitting plasma.    For some years we have argued \\citep[e.g.,][ and earlier references therein]{Drake03a} that coronal abundances, when better understood, might provide new and powerful diagnostics of the physical processes underpinning stellar coronae.  The emerging patterns of coronal abundance anomalies are telling us something about the dynamical structure and heating of coronal plasma; the challenge is to learn to read these patterns. Aiming toward this goal, one question that arises is that of fine tuning of abundance patterns: how do the abundance anomalies in the late-type giants compare with those of solar-like and other active stars?  In low-mass (M$<$1.5\\,M$_\\odot$) solar-type stars the magnetic activity is thought to derive from a dynamo powered by convection and driven by stellar spin \\citep[e.g.][]{Parker70}. Thus extremes in activity and rotation are closely related.  Moderate-mass giants (M$\\sim$3.0\\,M$_\\odot$), with A and late B progenitors on the MS having no outer convection zone, cross into the cool half of the H-R diagram during their post-main-sequence phase as yellow giants. The magnetic activity of these stars experiences several stages as they evolve: the Hertzsprung-gap phase of relatively rapid rotation \\citep{Simon89}; the ``rapid braking zone'' \\citep{Gray89} and { the red giant branch (RGB) phase \\citep{Ayres83}. The latter includes the ascent of the RGB, the helium ignition at its tip, and the return to the base of the RGB as core helium burning (``clump'') stars. The rejuvenation of magnetic activity of more evolved ``clump'' giants is not well understood. Infusion of angular momentum at the base of the RGB \\citep{Simon89} and the rejuvenation of the dynamo by planet/brown dwarf accretion \\citep{Siess99} have been suggested as possible causes}. Coronal characteristics and X-ray emission, triggered by the onset of efficient convection occurring while passing through the Hertzprung-gap, may change in response to evolutionary changes of the rotation and internal structure of the star (e.g.; deepening of the convection zone, rapid expansion of the stellar radius and photosphere cool down). Yellow giants are particularly interesting due to the relatively short time scales in which internal structure changes which might affect the corona and the coronal abundances take place \\citep{Ayres98,Gondoin99}.  In this paper, using {\\it Chandra} High Energy Transmission Grating spectrograph (HETGS) observations, we present a comparative analysis of the coronal X-ray spectra of the active late-type giants $\\beta$ Cet, 31 Com and $\\mu$~Vel with particular emphasis on their abundances.  We first describe the three stars and briefly review earlier work (\\S\\ref{s:stars}), then in \\S\\ref{s:obs} we report on the {\\it Chandra} observations and data reduction.  The methods used for a differential emission measure analysis together with results obtained are shown in \\S\\ref{s:anal}. In \\S\\ref{s:results} we discuss our results on the coronal abundances and temperature structure and report our conclusions in \\S\\ref{s:concl}. ",
        "conclusions": "\\label{s:concl}  $\\beta$ Cet, 31 Com and $\\mu$ Vel represent the important stages through which intermediate mass late-type giants evolve during their lifetime (the Hetzsprung gap (31 Com), the rapid braking zone ($\\mu$ Vel) and the core helium burning ``clump'' phase ($\\beta$ Cet)).  As such, a comparison of their coronal properties provides an illuminating glimpse of any fundamental underlying differences in their magnetic dynamos and activity.  Based on an analysis of high resolution {\\it Chandra} X-ray spectra of these stars we draw the following conclusions. \\begin{enumerate}  \\item $\\beta$ Cet and $\\mu$ Vel show coronal temperature structures and element abundances that are remarkably similar.  Element abundances are characterized by a mild FIP-type effect in which the abundances of low FIP elements are enhanced relative to those of high FIP elements by a factor of $\\sim 2$.  While we cannot rule out this result as being a fluke of underlying photospheric, rather than coronal, abundances, the latter scenario is supported by similarities with abundance anomalies seen in low-intermediate activity dwarfs.  \\item The coronal temperature structure of 31 Com differs from those of $\\beta$ Cet and $\\mu$ Vel and exhibits a sharper peak at higher temperatures, as seen earlier by \\citet{Ayres98} based on EUVE spectra. Its corona is significantly hotter, as is evident from comparison of its spectrum with those of $\\beta$ Cet and $\\mu$ Vel: lines formed at hotter temperatures are stronger in 31 Com, and it also has a stronger short wavelength continuum. Element abundances are characterized by a lack of an obvious FIP effect and appear closer to photospheric estimates.   \\item We speculate that structural changes during the evolution of late-type giants are likely responsible for the small observed differences in coronal abundances and temperature structure.  In particular, the size of the convection zone coupled with the rotation rate seem obvious choices for playing a key role in determining coronal characteristics.   \\end{enumerate}"
    },
    "0512/astro-ph0512347_arXiv.txt": {
        "abstract": "We use new deep near-infrared (NIR) and mid-infrared (MIR) observations to analyze the 850$~\\mu$m image of the Great Observatories Origins Deep Survey-North region around the Hubble Deep Field-North. We show that much of the submillimeter background at this wavelength is picked out by sources with $H(AB)$ or 3.6~$\\mu {\\rm m} (AB)<23.25$ (1.8 $\\mu$Jy). These sources contribute an $850~\\mu$m background of $24\\pm2$~Jy~deg$^{-2}$. This is a much higher fraction of the measured background ($31-45$~Jy~deg$^{-2}$) than is found with current 20~cm or $24~\\mu$m samples. Roughly one-half of these NIR-selected sources have spectroscopic identifications, and we can assign robust photometric redshifts to nearly all of the remaining sources using their UV to MIR spectral energy distributions. We use the redshift and spectral type information to show that a large fraction of the $850~\\mu$m background light comes from sources with $z=0-1.5$ and that the sources responsible have intermediate spectral types. Neither the elliptical galaxies, which have no star formation, nor the bluest galaxies, which have little dust, contribute a significant amount of 850~$\\mu$m light, despite the fact that together they comprise approximately half of the galaxies in the sample. The galaxies with intermediate spectral types have a mean flux of $0.40\\pm0.03$~mJy at $850~\\mu$m and $9.1\\pm0.3~\\mu$Jy at 20~cm.  The redshift distribution of the NIR-selected 850~$\\mu$m light lies well below that of the much smaller amount of light traced by the more luminous, radio-selected submillimeter sources. We therefore require a revised star-formation history with a lower star-formation rate at high redshifts. We use a stacking analysis of the 20~cm light in the NIR sample to show that the star-formation history of the total 850~$\\mu$m sample is relatively flat down to $z\\sim 1$ and that half of the total star formation occurs at redshifts $z<1.4$. ",
        "introduction": "\\label{intro}  The integrated extragalactic background light (EBL) is a measure of the history of the luminous energy production of the universe from both star formation and active galactic nuclei (AGNs). Directly emitted light is seen in the UV and optical, whereas dust reradiated energy appears in the far-infrared (FIR) and submillimeter. \\emph{COBE} obtained detailed measurements of the EBL at FIR and submillimeter wavelengths \\citep[e.g.,][]{puget96,fixsen98}, showing that the total radiated emission reprocessed by dust in the FIR/submillimeter is comparable to the total measured optical EBL. However, to proceed further, we also need to know the redshift distribution of the sources contributing to the submillimeter background, and this information has been extremely difficult to obtain.  In the last decade, the submillimeter/millimeter EBL has been resolved into discrete sources by deep surveys with the Submillimeter Common-User Bolometer Array (SCUBA) on the 15~m James Clerk Maxwell Telescope (JCMT) and with the Max-Plank Millimeter Bolometer array on the 30~m IRAM telescope. Blank-field surveys have resolved sources in the $2-20$~mJy range that account for $\\sim20-30$\\% of the 850~$\\mu$m EBL \\citep[e.g.,][]{barger98,hughes98,barger99a,eales99,eales00,eales03,bertoldi00, scott02,webb03,borys03,wang04}. With the help of strong lensing, surveys in cluster fields have resolved sources over the $0.3-2$~mJy range that account for a further $45-65$\\% of the 850~$\\mu$m EBL \\citep{smail97,chapman02,cowie02,knudsen05}. Together these surveys provide a nearly complete resolution of the background at 850~$\\mu$m. The ``typical'' source contributing to the 850~$\\mu$m EBL has a mean flux of about 0.9~mJy and a median flux of about 0.6~mJy \\citep{cowie02}.  However, the redshift follow-up of the submillimeter sources has been very slow. Because of the large beam size ($15\\arcsec$) of SCUBA and the optically-faint nature of the dusty sources, identifying the optical and near-infrared (NIR) counterparts to the submillimeter sources is time consuming \\citep[e.g.,][]{barger99b,ivison00}. To date, the most successful identifications of the submillimeter sources rely on the empirical correlation between the nonthermal radio emission and the thermal dust emission \\citep[e.g.,][]{condon92}. Once the radio counterparts to the submillimeter sources are detected by radio interferometers, the redshifts of the sources can be crudely estimated using the radio-to-submillimeter flux ratios \\citep{carilli99,barger00,hughes02,ivison02,chapman03b} or accurately measured with optical spectroscopy \\citep{chapman03a,chapman05}. The radio-identified sources are mostly bright ($\\gg 2~$mJy) submillimeter sources at $z=1.5-3.5$, with properties similar to the local ultraluminous infrared galaxies (ULIRGs; $L_{\\rm IR}>10^{12}~L_{\\sun}$, where $L_{\\rm IR}$ is the $8-1000~\\mu$m infrared luminosity; see, e.g., \\citealp{sanders96}). We note, however, that because of the $K$-correction and the sensitivity limit in the radio, only $\\sim60\\%$ of the bright submillimeter sources are identified in the radio \\citep{barger00}. It is not known whether the remaining 40\\% are at higher redshifts that simply cannot be reached by current radio telescopes.  Importantly, however, the properties and redshift distribution of the faint submillimeter sources that dominate the submillimeter EBL remain essentially unknown. The absence of any redshift information for more than 90\\% of the 850~$\\mu$m EBL represents a formidable uncertainty in determining the star-formation history, and this is what we aim to resolve in the present paper.  Like the radio emission, the mid-infrared (MIR) emission at $\\gtrsim5~\\mu$m could serve as another proxy to the submillimeter emission, since it also comes from dust. The MIR window has been opened by the Infrared Array Camera (IRAC, \\citealp{fazio04}) and the Multiband Imaging Photometer for \\emph{Spitzer} (MIPS, \\citealp{rieke04}) on the \\emph{Spitzer Space Telescope} \\citep[e.g.,][]{huang04,serjeant04,ivison04,egami04}. MIPS should be sensitive to $z\\lesssim1$ galaxies with infrared luminosities similar to local normal galaxies ($L_{\\rm IR}\\sim 10^{10}~L_{\\sun}$, corresponding to $\\sim0.1$~mJy at 850~$\\mu$m) and to $z\\lesssim3.5$ ULIRGs (i.e., typical of the bright submillimeter sources). Thus, MIPS should be able to detect the radio-identified submillimeter sources at $z\\lesssim3.5$ and to provide a large sample of faint sources that are beyond the confusion limit of current submillimeter telescopes \\citep[e.g.,][]{chary04}. However, as we shall show in this paper, even the extraordinarily deep MIPS data of the Great Observatories Origins Deep Survey-North (GOODS-N) \\emph{Spitzer} Legacy Science Program in the Hubble Deep Field-North (HDF-N) region does not substantially identify the 850~$\\mu$m EBL.  Remarkably, however, the combination of a $J$ or $H$-band sample (selected from images obtained with the new generation of ground-based, wide-field NIR cameras) and the IRAC 3.6~$\\mu$m sample does identify much of the 850~$\\mu$m EBL. We show this using the $H$-band image of the GOODS-N region obtained by \\citet{trouille06}. This result makes sense if the bulk of the sources contributing to the 850~$\\mu$m EBL are actually at lower redshifts and luminosities than those identified at the brighter submillimeter fluxes. Such sources have strong rest-frame optical/NIR counterparts that are picked up in the NIR sample. We utilize the spectroscopic and photometric redshift information on our NIR sample to confirm this result. We find that more than half of the 850~$\\mu$m EBL arises in sources with $z<1.5$ and that the sources that are responsible have intermediate spectral types. Neither the elliptical galaxies, which have no star formation, nor the bluest galaxies, which have little dust, contribute substantially to the 850~$\\mu$m EBL, despite the fact that together they comprise approximately half of the sample.  This result has profound implications for our understanding of the star-formation history, lowering previous estimates of the high-redshift star formation rate densities by factors of at least two. We analyze the star-formation history of our NIR sample using a 20~cm stacking analysis and compare this with the maximum star formation rate density at higher redshifts obtained directly from the submillimeter light. Together these show that the total star formation rate density peaks at a redshift at or just below one and is roughly flat at higher redshifts.  The paper is organized as follows. The submillimeter, NIR, MIR, optical, radio, and X-ray data are described in \\S\\ref{secsample}. The spectroscopic and photometric redshifts are discussed in \\S\\ref{secz}. The use of the NIR, MIR, and radio populations to identify the submillimeter background is discussed in \\S\\ref{secebl}, and the 850~$\\mu$m EBL identified by the NIR-sample is broken down by galaxy flux, color, spectral type, and redshift. The star-formation history is described in \\S\\ref{section_evolution}. Our main results are summarized in \\S\\ref{secsummary}. Throughout the paper, we assume the WMAP cosmology: $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.73$, and $\\Omega_{\\Lambda}=0.27$ \\citep{bennett03}. ",
        "conclusions": "\\label{secsummary}  We have obtained accurate redshifts for the sources in the GOODS-N area using existing spectroscopic redshifts and improved photometric redshifts from NIR and MIR data. The radio-identified bright ($>2$~mJy) SCUBA sources in this area are in the redshift range $z\\sim1-3$ and have a median redshift of $z=2.5$, consistent with previous radio and spectroscopic surveys.  However, we used a stacking analysis to show that much of the 850~$\\mu$m EBL is in fact traced by a NIR sample constructed from sources with fluxes greater than $1.8~\\mu$Jy in either the $H$ or $3.6~\\mu$m bands. We showed that much of this light arises from galaxies with intermediate spectral types at $z<1.5$. Thus, many of the fainter submillimeter sources that give rise to most of the 850~$\\mu$m EBL are at lower redshifts and lower luminosities than the bright submillimeter sources that are detected directly.  Finally, we used a stacking analysis to estimate the average 20~cm EBL produced by the unidentified or intermediate spectral type galaxies in our NIR sample as a function of redshift, from which we determined the SFRD. We found that this SFRD evolves rapidly between $z=0$ and $z=0.8$, after which it becomes approximately flat. Using the submillimeter data directly, we then calculated a submillimeter based SFRD at $z\\sim 1-3$  which agrees closely with the radio based SFRD. In addition, by assuming that all of the submillimeter EBL that is not accounted for by our NIR sample is also at these redshifts, we put an upper bound on the SFRD at $z\\sim 1-3$. Even with this maximum completeness correction, we found consistency with a nearly flat or slowly rising extrapolation of the SFRD from $z\\sim1$. We conclude that the majority of the star formation traced by the submillimeter light comes from redshifts near one rather than at the higher redshifts that have been favored until now.  \\vskip -0.3cm"
    },
    "0512/astro-ph0512037_arXiv.txt": {
        "abstract": "We report on our first results from a mid-infrared spectroscopic study of ISM features in a sample of deeply obscured ULIRG nuclei using the InfraRed Spectrograph (IRS) on the Spitzer Space Telescope. The spectra are extremely rich and complex, revealing absorption features of both amorphous and crystalline silicates, aliphatic hydrocarbons, water ice and gas phase bands of hot CO and warm C$_2$H$_2$, HCN and CO$_2$. PAH emission bands were found to be generally weak and in some cases absent. The features are probing a dense and warm environment in which crystalline silicates and water ice are able to survive but volatile ices, commonly detected in Galactic dense molecular clouds, cannot. If powered largely by star formation, the stellar density and conditions of the gas and dust have to be extreme not to give rise to the commonly detected emission features associated with starburst. ",
        "introduction": "With the launch of the InfraRed Spectrograph (IRS) on the Spitzer Space Telescope, mid-infrared spectroscopists have been handed a powerful tool to study extragalactic objects at high signal-to-noise and over a wide mid-infrared wavelength range, previously only available for the study of Galactic sources and a few nearby galaxies.  The foundations for the Spitzer studies currently underway were laid by the Infrared Space Observatory (ISO), which freed mid-infrared spectroscopists from the confinements of the Earth's atmospheric windows and the atmosphere's glaring foreground emission.  Major extragalactic topics addressed early on in the ISO mission centered around the properties of dusty starbursts and how they evolve (\\cite[e.g. Thornley et al. 2000]{thornley00}), the unification of optically classified type 1 and 2 active galaxies in relation to the properties of the Active Galactic Nucleus (AGN) (\\cite[e.g. Clavel et al. 2000; Laurent et al. 2000]{clavel00,laurent00}), and the dominant power source in Ultra-Luminous Infrared Galaxies (ULIRGs): massive starbursts or AGN activity (\\cite[e.g. Genzel et al. 1998; Tran et al. 2001]{genzel98,tran01})?  After the expiration of ISO, two unusual galaxy spectra provided first mid-infrared insights into the properties of gas and dust in deeply obscured galactic nuclei. The 2--5\\,$\\mu$m spectrum of the nucleus of NGC\\,4945 revealed strong absorptions of water ice (3\\,$\\mu$m), CO$_2$ (4.26\\,$\\mu$m) and a blend of `XCN' and CO ice (4.65\\,$\\mu$m), seen against a continuum obscuring the central massive black hole (\\cite[Spoon et al. 2000]{spoon00}). Groundbased follow-up observations confirmed the 4.65\\,$\\mu$m absorption band to consist of separate components of processed `XCN' and CO ice and warm (35\\,K) CO gas. The detection of processed ices against the nuclear continuum is taken as an indication for the presence of dense star forming molecular clouds towards the nucleus of this active galaxy (\\cite[Spoon et al. 2003]{spoon03}).  The second unusual spectrum is that of the nucleus of NGC\\,4418, originally classified by \\cite{roche86} as a ``very extinguished galaxy'', for its very deep 10\\,$\\mu$m silicate absorption feature. Instead of the commonly detected PAH emission features, the 5.5--8\\,$\\mu$m ISO spectrum of its nucleus is dominated by absorption features of water ice (6\\,$\\mu$m), hydrocarbons (6.85\\,$\\mu$m and 7.25\\,$\\mu$m) and methane ice (7.67\\,$\\mu$m), reminiscent of the line of sight towards our own Galactic Center (\\cite[Spoon et al. 2001]{spoon01}). Supporting groundbased observations indicate that most of the infrared luminosity (L$_{\\rm IR}$\\,=\\,10$^{11}$\\,L$_{\\odot}$) from NGC\\,4418 is produced in a compact nucleus with a radius of less than 80\\,pc (\\cite[Evans et al. 2003]{evans03}). If powered entirely by star formation, the conditions of the nuclear gas and dust within this environment must be exceptional not to give rise to emission features commonly associated with star formation.  \\begin{figure} \\begin{center} \\includegraphics{f1.ps} \\end{center} \\caption{Spitzer IRS low-resolution spectra of ULIRGs sorted by spectral shape. The three spectra at the top are continuum-dominated (AGN-like), the next four are PAH-dominated (starburst-like) and the rest are absorption-dominated (burried nuclei). Vertical lines indicate the positions of the 6.2, 7.7 and 11.2\\,$\\mu$m PAH emission bands} \\label{fig:sed_trends} \\end{figure}  A large scale follow-up study into the presence of 5.5--8\\,$\\mu$m absorption features in ISO galaxy spectra resulted in the finding of 6\\,$\\mu$m water ice absorption in 12 out of 19 ULIRGs, 2 out of 62 AGNs and 4 out of 21 starburst galaxies surveyed. Also, 6.85\\,$\\mu$m hydrocarbon absorption was detected in three galaxies besides NGC\\,4418, all three of which are ULIRGs. The results are consistent with findings from other wavelength ranges that more molecular material is present in ULIRG nuclei than in other galaxy types (\\cite[Spoon et al. 2002]{spoon02}).  In the following sections we present an overview of the first results from an IRS spectroscopic study of the properties of gas and dust in strongly obscured ULIRG nuclei. The spectra were selected from a larger sample of ULIRG spectra obtained as part of the GTO program of the Spitzer IRS team. ",
        "conclusions": "Using the IRS spectrograph on the Spitzer Space Telescope, we have obtained mid-infrared spectra for a large sample of ULIRGs. The spectra show a great diversity in spectral shape, reflecting the diverse nature and merger evolutionary states of the sample.  Especially interesting for astrochemists is the subsample of spectra showing signatures of strong obscuration, betraying the presence of huge amounts of dust and gas in and towards the merger nuclei. The wide spectral coverage of the IRS (5--38\\,$\\mu$m), assisted by redshifts ranging from z=0.02 to z=0.4, allowed us to study the ISM in these sources from rest frame $\\sim$3.8\\,$\\mu$m to 37\\,$\\mu$m, revealing absorption features of both amorphous and crystalline silicates, aliphatic hydrocarbons, water ice and gas phase bands of hot CO and warm C$_2$H$_2$, HCN and CO$_2$. PAH emission bands were found to be generally weak and in some cases absent.  Our analysis of the absorption features is far from complete and requires further comparison to the results from our emission line analysis to obtain a more detailed picture of the energetic processes responsible for these rich and complicated spectra. It is clear, however, at this time that the features are probing a dense and warm environment in which crystalline silicates and water ice are able to survive but volatile ices, commonly detected in Galactic dense molecular clouds, cannot. If powered largely by star formation, the stellar density and conditions of the gas and dust have to be extreme not to give rise to the commonly detected emission features associated with starburst."
    },
    "0512/gr-qc0512070_arXiv.txt": {
        "abstract": "We use the dynamical systems approach to investigate the Bianchi type VIII models with a tilted $\\gamma$-law perfect fluid.  We introduce expansion-normalised variables and investigate the late-time asymptotic behaviour of the models and determine the late-time asymptotic states. For the Bianchi type VIII models the state space is unbounded and consequently, for all non-inflationary perfect fluids, one of the curvature variables grows without bound. Moreover, we show that for fluids stiffer than dust ($1<\\gamma<2$), the fluid will in general tend towards a state of extreme tilt. For dust ($\\gamma=1$), or for fluids less stiff than dust ($0<\\gamma< 1$), we show that the fluid will in the future be asymptotically non-tilted. Furthermore, we show that for all $\\gamma\\geq 1$ the universe evolves  towards a vacuum state but does so rather slowly, $\\rho/H^2\\propto 1/\\ln t$. ",
        "introduction": "Cosmology in the recent years has proven to be an important arena to perform tests of the general theory of relativity. Today, theoretical cosmology gives insights into the possible physical and mathematical properties of the universe while observations help to constrain the various theoretical possibilities. The purpose of this paper is to fill  one of the major gaps in the theoretical understanding of the behaviour of spatially homogeneous (SH) cosmologies \\cite{EM,BS,DS1,DS2}. The aim is to give a description of the evolution of a general spatially homogeneous model with a perfect fluid, in particular, the  SH Bianchi type VIII models with a $\\gamma$-law perfect fluid. Until now, only Bianchi type VIII models where the fluid flow is orthogonal to the surfaces of homogeneity have been studied. More specifically, vacuum Bianchi type VIII models were studied in \\cite{BG,Ringstrom1,Ringstrom2} and a detailed derivation of the asymptotic expansions of Bianchi type VIII vacuum metrics was given in \\cite{Ringstrom3}. The future asymptotic behaviour with a non-tilted perfect fluid was also studied in \\cite{VIII}. Here, we will allow for the fluid flow to be tilted; i.e., where the fluid flow is not orthogonal to the surfaces of homogeneity \\cite{KingEllis}.   Several SH models with tilt have been studied before: type II \\cite{HBWII}, IV \\cite{CH,HHC}, V \\cite{Shikin,Collins,CollinsEllis,HWV,Harnett}, VI$_0$ \\cite{hervik,coleyhervik}, VII$_0$ \\cite{HHLC}, and VII$_h$ \\cite{CH,HHC} (see also \\cite{CH,BHtilted} for a subclass of the type VI$_h$ models). Allowing for a tilted fluid, results in several interesting new phenomena, such as  future limiting curves and attracting tori \\cite{HHC}. Most of the above-mentioned works have been utilizing the theory of dynamical systems in their investigations \\cite{DS1,DS2,BN}; we will do the same by generalising the formalism from the solvable case \\cite{CH} to the semisimple Bianchi type VIII models.   The Bianchi type VIII models are particularly interesting models and are the most general ever-expanding cosmological model of Bianchi type. The state space of the tilted $\\gamma$-law type VIII models is of dimension 8 (compared to 5 in the non-tilted case and 4 in the vacuum case); hence, by studying the type VIII models we can gain insight into the behaviour of a `general' Bianchi model.  Under many physical circumstances there are certain  self-similar solutions that act as attractors for more general solutions of the model. This `self-similarity hypothesis' \\cite{CC1,CC2} is known to be valid for many SH models, however, there are a few notable exceptions. It was pointed out in \\cite{VII0} that for the non-tilted perfect fluid Bianchi type VII$_0$ models one of the curvature variables grows without bound and the models are consequently not asymptotically self-similar. The same is also true for the non-tilted perfect fluid type VIII models \\cite{VIII}. By allowing for a tilted perfect fluid we will see that this property does not change. Neither the tilted type VII$_0$ models (which was shown in \\cite{HHLC}) nor the tilted type VIII models are asymptotically self-similar.  The Bianchi type VIII Lie algebra corresponds to the Lie algebra of the matrix group $SL(2,\\mathbb{R})$ with a connected covering group usually denoted by $\\widetilde{SL(2,\\mathbb{R})}$. This group is one of the eight so-called \\emph{Thurston geometries} (see, e.g. \\cite{thurston97}). The Thurston geometries play an important role in the famous geometrization conjecture for 3-manifolds by Thurston \\cite{thurston}\\footnote{See, e.g. \\cite{Morgan} for a review of the geometrization conjecture and a discussion of a recent attempt of proving it.}. The type VIII model is usually considered to be the connected and simply connected group $\\widetilde{SL(2,\\mathbb{R})}$ (we will do the same here); however, it is worth pointing out that the group  $\\widetilde{SL(2,\\mathbb{R})}$ allows for a discrete subgroup $\\Gamma$ such that the quotient $\\widetilde{SL(2,\\mathbb{R})}/\\Gamma$ is compact. In this regard Barrow and Kodama \\cite{BK1,BK2} pointed out that as the complexity of the quotient $\\widetilde{SL(2,\\mathbb{R})}/\\Gamma$ increases, the number of parameters describing the corresponding Bianchi model grows without bound.  A set of left-invariant one-forms on $\\widetilde{SL(2,\\mathbb{R})}$ can, for example, be given by: \\beq {\\mbold\\omega}^1&=&a\\left(\\d x-\\frac{\\d z}{y}\\right), \\nonumber \\\\ {\\mbold\\omega}^2&=&\\frac by\\left(\\cos x\\ \\d y+\\sin x\\ \\d z\\right), \\nonumber \\\\ {\\mbold\\omega}^3&=&\\frac cy\\left(-\\sin x\\ \\d y+\\cos x\\ \\d z\\right). \\label{VIII1forms}\\eeq These one-forms fulfil the type VIII Lie algebra relations: \\beq \\d {\\mbold\\omega}^i=-\\frac 12C^i_{jk}{\\mbold\\omega}^j\\wedge{\\mbold\\omega}^k, \\quad C^i_{jk}=\\varepsilon_{jkl}n^{li}, \\eeq where $n^{ij}$ is a symmetric matrix with eigenvalues $\\lambda_1<0<\\lambda_2,\\lambda_3$. For the choice (\\ref{VIII1forms}), \\beq (n^{ij})=\\mathrm{diag}\\left(-\\frac{a}{bc},\\frac{b}{ac},\\frac{c}{ab}\\right). \\eeq Using the one-forms (\\ref{VIII1forms}) we can introduce a Bianchi type VIII cosmology by \\beq \\d s^2=-\\d t^2+\\delta_{ab}\\widetilde{\\mbold\\omega}^a\\widetilde{\\mbold\\omega}^b, \\quad \\widetilde{\\mbold\\omega}^a={\\sf e}^a_{~i}(t){\\mbold\\omega}^i. \\eeq Here, $t$ is the proper time of an observer whose world-line is a geodesic orthogonal to the type VIII surfaces of homogeneity (also called the cosmological time). The hypersurface orthogonal vector is ${\\bf n}=\\partial/\\partial t$.  The energy-momentum tensor of the perfect fluid is \\beq T_{\\mu\\nu}=(p+\\rho)u_{\\mu}u_{\\nu}+pg_{\\mu\\nu}, \\eeq where $\\rho$, $p$ and $u^{\\mu}$ are the energy density, pressure and four-velocity of the fluid, respectively. The equation of state will be taken to be \\beq p=(\\gamma-1)\\rho, \\eeq where $\\gamma$ is a constant. This choice includes the important cases of dust ($\\gamma=1$) and radiation ($\\gamma=4/3$). In this paper we will study models where the four-velocity $u^{\\mu}$ of the perfect fluid is not parallel with the normal vector $n^{\\mu}$. Since these models are future geodesically complete \\cite{Rendall}, we will choose the fundamental observers to follow the geodesic congruences defined by the vector field $n^{\\mu}$. This choice avoids the possible singular behaviour that may occur for observers following the fluid flow lines \\cite{CollinsEllis,CHL}. ",
        "conclusions": "\\begin{table} \\centering \\begin{tabular}{|c|c|c|l|} \\hline Invariant &   &  &  \\\\ subspace & Matter & Attractor & Asymptotic tilt \\\\ \\hline \\hline $T(VIII)$ & $2/3<\\gamma<1$ & $\\widetilde{P}_1$ & non-tilted (e)\\\\ & $\\gamma=1$ & $\\widetilde{P}_2$ & non-tilted (p) \\\\ & $1<\\gamma<2$ & $\\widetilde{E}_1$ & extremely tilted (e) \\\\ \\hline $N(VIII)$ & $2/3<\\gamma<1$ & $\\widetilde{P}_1$ & non-tilted (e)\\\\ & $\\gamma=1$ & $\\widetilde{P}_2$ & non-tilted (p) \\\\ & $1<\\gamma<2$ & $\\widetilde{E}_1$ & extremely tilted (e) \\\\ \\hline $F(VIII)$ & $2/3<\\gamma<1$ & $\\widetilde{P}_1$ & non-tilted (e) \\\\ & $1\\leq\\gamma<3/2$ & $\\widetilde{P}_2$ & non-tilted (e) \\\\ & $\\gamma=3/2$ & $\\widetilde{E}_2$ & extremely tilted (p) \\\\ & $3/2<\\gamma<2$ & $\\widetilde{E}_2$ & extremely tilted (e) \\\\ \\hline $T_1(VIII)$ &$2/3<\\gamma<1$ & $\\widetilde{P}_1$ & non-tilted (e)  \\\\ & $\\gamma=1$ & $\\widetilde{P}_2$ & non-tilted (p) \\\\ & $1<\\gamma<2$ & $\\widetilde{E}_1$ & extremely tilted (e) \\\\ \\hline $B(VIII)$ &  $2/3<\\gamma<1$ & $\\widetilde{P}_1$ & no tilt\\\\ & $1\\leq\\gamma<2$ & $\\widetilde{P}_2$ & no tilt\\\\ \\hline \\end{tabular} \\caption{The late-time behaviour of the VIII Bianchi  model with a tilted $\\gamma$-law perfect fluid (see the text for details and references). The comments refer to the late-time asymptotics, and for all cases $\\bar{N}\\rightarrow \\infty$. The case $0<\\gamma<2/3$ is covered by the no-hair theorem (the non-tilted version is given in \\cite{Wald}, the tilted version in \\cite{CH}). The right-most column indicates the asymptotic tilt and whether the tilt velocity, $V$, approaches this state exponentially (e) or power law (p) in terms of the dynamical time $\\tau$.} \\label{tab:outline} \\end{table}  Table \\ref{tab:outline} displays an outline of the late-time asymptotic behaviour of tilted Bianchi type VIII models. As for the late-time behaviour of the Bianchi models, all of the non-type-IX class A models have now been studied. The only remaining class A model, namely the closed type IX model is notoriously difficult and requires a different formalism. Of the class B models, an obvious lacuna is the type VI$_h$ model, all other tilted class B models have been studied.  Regarding the tilted type VIII models we have seen that the extreme Weyl-curvature dominance (in the terminology of \\cite{BHWeyl}) for $4/5<\\gamma<2$ which was found for the non-tilted models, persists into the tilted model. This extreme Weyl-curvature dominance is a signal of the self-similarity breaking that occurs at late times for these models. More explicitly, this Weyl-curvature dominance is a result of an increasingly rapid oscillation that takes place in the shear and the curvature variables. This rapid oscillation is exactly what prevents the type VIII models to be asymptotically self-similar.  Moreover, we have shown that for the type VIII models with fluids stiffer than dust ($1<\\gamma<2$) the tilt becomes asymptotically extreme at late times. The energy density itself asymptotically approaches vacuum but does so rather slowly, \\[ \\Omega\\propto\\tau^{-1}\\sim (\\ln t)^{-1}.\\] We can compare this slow decay with the non-tilted case for which: \\[ \\text{non-tilted VIII:} \\quad \\Omega\\propto \\tau^{-\\frac 12}e^{-3(\\gamma-1)\\tau}\\sim {t^{-2(\\gamma-1)}(\\ln t)^{-\\frac\\gamma 2}}.\\] In particular, this means that for a radiation dominated tilted Bianchi type VIII model the expansion-normalised energy density, $\\Omega$, only decays logarithmically at late times, in terms of the cosmological time, $t$.  In this paper we have discussed the late-time behaviour of Bianchi type VIII cosmologies. Regarding the early-time behaviour, the models undergo Mixmaster dynamics \\cite{Ringstrom1,Jantzen}, which is also the typical behaviour for inhomogeneous models \\cite{UvEWE,Garfinkle}. While the early-time behaviour is similar in inhomogeneous generalisation of Bianchi models, we do not expect this for the late-time behaviour due to the dominance of inhomogeneity at late times (see, e.g., \\cite{DS1,HC})."
    },
    "0512/astro-ph0512201_arXiv.txt": {
        "abstract": "{We calculate the redshift-space power spectrum of the Sloan Digital Sky Survey (SDSS) Data Release 4 (DR4) Luminous Red Galaxy (LRG) sample, finding evidence for a full series of acoustic features down to the scales of $\\sim 0.2\\,h\\,\\mathrm{Mpc}^{-1}$. This corresponds up to the 7th peak in the CMB angular power spectrum. The acoustic scale derived, $(105.4 \\pm 2.3)\\,h^{-1}\\,\\mathrm{Mpc}$, agrees very well with the ``concordance'' model prediction and also with the one determined via the analysis of the spatial two-point correlation function by \\citet{2005ApJ...633..560E}. The models with baryonic features are favored by $3.3 \\sigma$ over their ``smoothed-out'' counterparts without any oscillatory behavior. This is not only an independent confirmation  of \\citet{2005ApJ...633..560E} results made with different methods and software but also, according to our knowledge, the first determination of the power spectrum of the SDSS LRG sample. ",
        "introduction": "In the beginning of $1970$'s it was already realized that acoustic waves in the tightly coupled baryon-photon fluid prior to the epoch of recombination will lead to the characteristic maxima and minima in the post-recombination matter power spectrum. The same mechanism is also responsible for the prominent peak structure in the CMB angular power spectrum \\citep{1970Ap&SS...7....3S,1970ApJ...162..815P,1978SvA....22..523D}. The scale of these features reflects the size of the sound horizon, which itself is fully determined given the physical densities $\\Omega_b h^2$ and $\\Omega_m h^2$. The acoustic horizon can be calibrated using the CMB data, thus turning it into a standard ruler which can be used to carry out classical cosmological tests. For example, if we are able to measure the redshift and angular intervals corresponding to the physically known acoustic scale in the matter power spectrum at a range of redshifts, we can immediately find angular diameter distance $d_{\\rm A}$ and Hubble parameter $H$ as a function of redshift. Having good knowledge of these dependencies allows us to put constraints on the properties of the dark energy. To carry out this project one needs a tracer population of objects whose clustering properties with respect to the underlying matter distribution is reasonably well understood. There have been several works discussing the usage of galaxies \\citep{2003ApJ...594..665B,2003PhRvD..68f3004H,2003PhRvD..68h3504L,2003ApJ...598..720S} and clusters of galaxies \\citep{2003PhRvD..68f3004H,2004ApJ...613...41M,astro-ph/0505441} for this purpose. What is most important is that already currently existing galaxy redshift surveys have lead to the detection of acoustic features in the spatial distribution of galaxies, this way providing clearest support for the standard gravitational instability picture of the cosmic structure formation. In the paper by \\citet{2005ApJ...633..560E} the detection of the acoustic ``bump'' in the two-point redshift-space correlation function of the SDSS \\footnote{http://www.sdss.org/} LRG sample is announced. The discovery of similar features in the power spectrum of 2dF \\footnote{http://www.mso.anu.edu.au/2dFGRS/} galaxies is presented in \\citet{2005MNRAS.362..505C}. These results clearly demonstrate the great promise of the future dedicated galaxy redshift surveys like K.A.O.S.\\footnote{http://www.noao.edu/kaos/} Similarly, useful measurements of the acoustic scale can be hoped by the planned SZ cluster surveys like the ones carried out by the PLANCK Surveyor \\footnote{http://astro.estec.esa.nl/Planck} spacecraft and SPT \\footnote{http://astro.uchicago.edu/spt} \\citep{astro-ph/0505441} and also with a large future photometric redshift surveys \\citep{2005MNRAS.tmp..876B}. For the SZ surveys one needs an additional optical follow-up to get estimates for the cluster redshifts. In this paper we calculate the redshift-space power spectrum of the SDSS LRG sample finding evidence for the acoustic oscillations down to the scales of $\\sim 0.2\\,h\\,\\mathrm{Mpc}^{-1}$, which effectively correspond up to the 7. peak in the CMB angular power spectrum. These scales in the CMB are very strongly damped due to the finite width of the last-scattering surface and also due to the Silk damping \\citep{1968ApJ...151..459S}. This can be seen in Fig. \\ref{fig1} \\footnote{Here instead of the usual multipole number $\\ell$ we have plotted the CMB angular power spectrum against the wavenumber $k$. For the ``concordance'' cosmological model $\\ell = 9990\\,k[h\\,\\mathrm{Mpc}^{-1}]$.} where the CMB data is plotted in a somewhat unusual way to enhance the acoustic features at the high wavenumber damping tail. Also, at those scales the secondary CMB anisotropies (mostly thermal Sunyaev-Zeldovich effect \\citep{1972CoASP...4..173S,1980ARA&A..18..537S}) start to dominate over the primary signal. On the other hand, features in the matter power spectrum, although being small ($\\sim 5\\%$ fluctuations), are preserved by the linear evolution and so opening up the way to probe acoustic phenomena at scales smaller than the ones accessible for the CMB studies.  The paper is structured as follows. In Sec. 2 we describe the dataset to be analyzed. Sec. 3 presents the method of the power spectrum calculation. In Sec. 4 we determine power spectrum errors and covariance matrix. Sec. 5 discusses the convolution effect of the survey window. Analytical model spectra are presented in Sec. 6. The results of the measurement of the acoustic scale are given in Sec. 7. Correlation function analysis is carried out in Sec. 8. In Sec. 9 we compare the measured power spectrum with the published results for the 2dF and SDSS main sample, and finally we conclude with Sec. 10.  \\begin{figure} \\centering \\includegraphics[width=\\plotwd] {Figs/fig01.eps} \\caption{Acoustic oscillations in the CMB (upper panel) and linear matter power spectrum (lower panel) for the ``concordance'' cosmological model. Here, as we have plotted the spectra against spatial wavenumber $k$, we have changed the standard notation of $C_\\ell$ to $C_k$. Due to the $k^3$ factor the first CMB acoustic peak is barely visible. Density fluctuations in matter at smaller scales, being mostly induced by the velocity fields, are out of phase with respect to the fluctuations in the CMB component. Also the fluctuation period is twice as large.} \\label{fig1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\plotwd] {Figs/fig02.eps} \\caption{Comoving number density of galaxies as a function of comoving distance. Smooth solid line shows a cubic spline fit to the number density estimated for 50 discrete radial bins.} \\label{fig2} \\end{figure} ",
        "conclusions": "In this paper we have calculated the redshift-space power spectrum of the SDSS DR4 LRG sample, finding evidence for a series of acoustic features down to the scales of $\\sim 0.2\\,h\\,\\mathrm{Mpc}^{-1}$. It turns out that models with the baryonic oscillations are favored by $3.3 \\sigma$ over their ``smoothed-out'' counterparts without any oscillatory behavior. Using the obtained power spectrum we predict the shape of the spatial two-point correlation function, which agrees very well with the one obtained directly from the data. Also, the directly calculated correlation function is consistent with the results obtained by \\citet{2005ApJ...633..560E}. We have made no attempts to put constraints on the cosmological parameters, rather we have assumed in our analysis the ``concordance'' cosmological model. The derived acoustic scale $(105.4 \\pm 2.3)\\,h^{-1}\\,\\mathrm{Mpc}$ agrees well with the best-fit WMAP ``concordance'' model prediction of $\\simeq 106.5 \\,h^{-1}\\,\\mathrm{Mpc}$.  The existence of the baryonic features in the galaxy power spectrum is very important, allowing one (in principle) to obtain Hubble parameter $H$ and angular diameter distance $d_A$ as a function of redshift, this way opening up a possibility to constrain properties of the dark energy \\citep{2003PhRvD..68f3004H}. The currently existing biggest redshift surveys, which are still quite shallow, do not yet provide enough information to carry out this project fully. On the other hand, it is extremely encouraging that even with the current generation of redshift surveys we are already able to see the traces of acoustic oscillations in the galaxy power spectrum, showing the great promise for the dedicated future surveys like K.A.O.S. We have seen that acoustic features seem to survive at mildly nonlinear scales ($k \\gtrsim 0.1\\,h\\,\\mathrm{Mpc}^{-1}$), which is in agreement with the results of the recent N-body simulations \\citep{2005Natur.435..629S,2005ApJ...633..575S}. In order to fully exploit available information one needs a complete understanding of how nonlinear effects influence these features. Nonlinear bias and redshift space distortions also add extra complications. In general redshift-space distortions, biasing and nonlinear evolution do not create any oscillatory modulation in the power spectrum and so acoustic features should be readily observable. So far there have been only a few works studying these important issues (e.g.  \\citet{2005Natur.435..629S,2005ApJ...633..575S,2005APh....24..334W}) and probably it is fair to say that currently we really do not have a full theoretical description of them. In our paper we have modeled the above mentioned effects using the results from the 2nd order Lagrangian perturbation theory in combination with the Halo Model. Although these models are very successful in capturing many important aspects of the structure formation, one has to keep in mind that they are still approximations.  The bare existence of the baryonic oscillations in the galaxy power spectrum tells us something important about the underlying cosmological model and the mechanism of the structure formation. First, it confirms the generic picture of the gravitational instability theory where the structure in the Universe is believed to be formed by the gravitational amplification of the small perturbations layed down in the early Universe. Under the linear gravitational evolution all the density fluctuation modes evolve independently i.e. all the features in the power spectrum will be preserved. And certainly, we are able to identify features in the low redshift galaxy power spectrum that correspond to the fluctuations seen in the CMB angular power spectrum (which probes redshifts $z \\sim 1100$), providing strong support for the above described standard picture of the structure formation. Actually, we can also probe scales that are unaccessible for the CMB studies due to the strong damping effects and steeply rising influence of the secondary anisotropies, reaching effectively the wavenumbers that correspond to the 6th-7th peak in the CMB angular power spectrum. Second, the ability to observe baryonic features in the low redshift galaxy power spectrum demands rather high baryonic to total matter density ratio. In \\citet{2003A&A...412...35B} it has been shown that it is possible to fit a large body of observational data with an Einstein--de Sitter type model if one adopts low value for the Hubble parameter and relaxes the usual assumptions about the single power law initial spectrum. In the light of the results obtained in our paper these models are certainly disfavored due to the fact that the high dark matter density completely damps the baryonic features. And finally, purely baryonic models are also ruled out since for them the expected acoustic scale would be roughly two times larger than observed here \\footnote{For a clear discussion of this see Daniel Eisenstein's home page http://cmb.as.arizona.edu/ $\\sim$eisenste/acousticpeak/}. So the data seems to demand a weakly interacting nonrelativistic matter component and all the models that try to replace this dark matter component with something else e.g. modifying the laws of gravity might have severe difficulties to fit these new observational constraints."
    },
    "0512/astro-ph0512322_arXiv.txt": {
        "abstract": "I review here the results of the first RV survey for spectroscopic companions to very young brown dwarfs (BDs) and (very) low-mass stars in the ChaI star-forming cloud with UVES at the VLT. This survey studies the binary fraction in an as yet unexplored domain not only in terms of primary masses (substellar regime) and ages (a few Myr) but also in terms of companion masses (sensitive down to planetary masses) and separations ($<$ 1\\,AU). The UVES spectra obtained so far hint at spectroscopic companions of a few Jupiter masses around one BD and around one low-mass star (M4.5) with orbital periods of at least several months. Furthermore, the data indicate a multiplicity fraction consistent with field BDs and stellar binaries for periods $<$100 days. ",
        "introduction": "The multiplicity properties of brown dwarfs (BDs) are key parameters for their formation. For example, embryo-ejection scenarios predict few binaries in only close orbits (see Delgado-Donate, this proceeding \\cite{joe:delgado}), while isolated fragmentation scenarios allow for an abundance of binaries over a wide range of separations. \\begin{figure}[b] \\centering \\includegraphics[height=0.7\\textwidth,angle=0]{joergensF1.eps} \\caption{ \\label{joe:fig2} Radial velocity data for the young BD candidate Cha\\,H$\\alpha$\\,8 (M6.5) recorded with UVES/VLT: significant variability occurring on time scales of months to years hint at a companion at $a>0.2\\,$\\,AU and a mass Msin\\,$i$ of at least 6 M$_\\mathrm{Jup}$ (Joergens 2005 \\cite{joe:j2005}). } \\end{figure} In recent years, numerous low-mass and BD binaries were detected by direct imaging in the field and in young clusters and associations (see Bouy, this proceeding \\cite{joe:bouy}). Based on these observations, it was found that very low-mass stars (VLMSs) and BDs pair less frequently in binary systems than solar-like stars. However, these surveys cannot resolve the inner $\\sim$3 to 10\\,AU (depending on distance) around the objects. Companions orbiting at such close separations may have originated based on a substantially different companion formation mechanism than the one found so far by direct imaging. They can be detected indirectly by spectroscopic Doppler surveys. Precise monitoring of the radial velocities (RVs) of BDs became possible in the last years with the generation of 8--10\\,m class telescopes. Shortly after the high-resolution echelle spectrograph UVES at the VLT saw its first light in October 1999, two programs were started with this instrument in spring 2000 and 2001 in order to systematically search for close companions to BDs and VLMSs in the very young Cha\\,I star-forming region (Joergens \\& Guenther 2001 \\cite{joe:j2001}; Joergens 2005 \\cite{joe:j2005}) and in the field (Guenther \\& Wuchterl 2003 \\cite{joe:gw}). In this article, we present and discuss the current results of the survey in Cha\\,I. ",
        "conclusions": "\\begin{figure}[b] \\centering \\includegraphics[height=.7\\textwidth,angle=0]{joergensF4.eps} \\caption{ \\label{joe:fig3} Radial velocity data for the low-mass star CHXR\\,74 (M4.5) recorded with UVES/VLT: significant variability occurring on time scales of months to years hint at a companion at $a>0.4\\,$\\,AU and a mass Msin\\,$i$ of at least 19 M$_\\mathrm{Jup}$ (Joergens 2005 \\cite{joe:j2005}). } \\end{figure}  The study of the multiplicity properties of BDs for separations of less than $\\sim$3--10\\,AU has been recognized as one of the main observational efforts that are necessary in order to constrain the formation of BDs. This can be done by means of high-resolution spectroscopic surveys. We have presented here the current results of the first RV survey of very young BDs (Joergens 2005 \\cite{joe:j2005}). It exploits the high resolving power and stability of the UVES spectrograph and the large photon collecting area of the VLT. A remarkable feature of this survey is that it is sensitive to planetary mass companions. This is due to a precise RV determination and to the fact that systematic RV errors caused by activity are sufficiently small for the targets to allow for the detection of Jupiter mass planets around them, as shown for the first time by this survey. Thus, very young BDs, at least in Cha\\,I, are suitable targets for the search for close extrasolar planets in contrast to very young stars.  None of the BDs and VLMSs monitored shows signs of BD or planetary companions for separations smaller than 0.1\\,AU. This hints at a small binary fraction and a low frequency of giant planets in this separation range (zero of ten). Within the limited statistics, this result is consistent with the binary frequency found for field BDs/VLMSs (12$\\pm$7\\%, Guenther \\& Wuchterl 2003 \\cite{joe:gw}) and with the frequency of stellar G dwarf binaries (7\\%, \\cite{joe:dm91}) in the same separation range. For some of the Cha\\,I targets also larger separations were probed leading to the detection of two candidates for spectroscopic systems: Both the BD candidate Cha\\,H$\\alpha$\\,8 (M6.5) and the low-mass star CHXR74 (M4.5) exhibit long-term RV variations that were attributed to orbiting companions with several Jupiter masses at minimum. Orbit solutions have to await follow-up observations, however, the data suggest orbital periods of at least several months, i.e. separations of $>$ 0.2\\,AU and 0.4\\,AU, resp.  Direct imaging surveys found a significantly lower frequency of BD binaries with separations $a>$3--10\\,AU compared to solar-like stars \\cite{joe:bouy}. This might be (partly) caused by a shift to smaller separations for lower mass primaries. The first surveys that probe the inner few AU around BDs by spectroscopic means are the one presented here (Joergens \\& Guenther 2001 \\cite{joe:j2001}, Joergens 2005 \\cite{joe:j2005}, fair detection efficiency for $a<$0.4\\,AU \\cite{joe:maxted}, sensitive to M$_{\\mathrm{Jup}}$ planets), and the following ones by other groups: Basri \\& Mart\\'{\\i}n (1999 \\cite{joe:basri}, detection of first spectroscopic BD binary), Reid et al. (2002 \\cite{joe:reid}, single epoch spectra, sensitive to double-lined spectroscopic binaries), Guenther \\& Wuchterl (2003 \\cite{joe:gw}, fair detection efficiency for $a<$0.7\\,AU \\cite{joe:maxted}, sensitive to planetary masses) and Kenyon et al. (2005 \\cite{joe:kenyon}, fair detection efficiency for $a<$0.02\\,AU \\cite{joe:maxted}, sensitive to BD companions). These surveys do not hint at a higher BD binary fraction at $a<$1\\,AU compared to stellar binaries indicating that also the overall binary frequency is lower in the substellar than in the stellar regime. While corrections have to be applied to the observed values because of selection biases (e.g. Burgasser et al. 2003) and sparse sampling of the velocity data (e.g. Maxted \\& Jeffries 2005 \\cite{joe:maxted}), a primary goal is to enlarge and improve the available data set for BD spectroscopic binary studies in terms of sample sizes, phase coverage, and precision of RV data."
    },
    "0512/astro-ph0512608_arXiv.txt": {
        "abstract": "{Random motions can occur in the intergalactic gas of galaxy clusters at all stages of their evolution. Depending on the poorly known value of the Reynolds number, these motions can or cannot become turbulent, but in any case they can generate random magnetic fields via dynamo action. We argue that magnetic fields inferred observationally for the intra\\-cluster medium require dynamo action, and then estimate parameters of random flows and magnetic fields at various stages of the cluster evolution. Polarization in cluster radio halos predicted by the model would be detectable with the SKA. ",
        "introduction": "We witness the epoch when galaxy clusters are being formed. Galaxies and smaller-size structures have generally achieved (quasi-)steady states in their evolution, whereas structures of larger scales keep evolving. There is abundant evidence of formation processes in galaxy clusters, including the merger of large masses comparable to the cluster mass. Relaxed, symmetric galaxy clusters are rare. Thus, most constituents of the clusters keep evolving. Our particular interest here is with motions of the gas and magnetic fields in the intergalactic space of evolving galaxy clusters. The epoch of major mergers is notable for widespread, intense random motions, followed by a period of decaying flows. In a steady state, random motions can be confined to the wakes produced by the cluster galaxies and smaller mass subcluster clumps that continue falling into the cluster; as we argue here, these motions are likely to occupy a small fraction of the total volume.  The high electrical conductivity of the intra\\-cluster gas is not sufficient to ensure the survival of a magnetic field captured by the forming cluster. The reason is that the intra\\-cluster gas is very viscous, and any inhomogeneous magnetic field would drive gas motions which will rapidly decay, thereby converting magnetic energy into heat. Even for a small viscosity, the field would decay by driving decaying MHD turbulence. On the other hand, random gas motions of the intra\\-cluster gas can be efficient generators of (random) magnetic fields via a process known as the fluctuation dynamo.  The origin and properties of magnetic field in the intra\\-cluster plasma cannot be understood without knowing the nature and parameters of the plasma motions. The reason is that the magnetic Reynolds number in the intra\\-cluster gas is large, hence magnetic induction effects are strong for any plausible speed of gas motion. Our model of gas motions and magnetic fields in galaxy clusters, presented and justified in detail by Subramanian et al.\\ (2006), is in close agreement with the available data. It predicts that magnetic field in the intra\\-cluster gas is represented by magnetic sheets of a thickness 20--30 kpc, with magnetic field strength about $2$--$4\\mkG$ within the sheets. These structures fill about 20\\% of the volume of a single turbulent cell. Apart from this intermittent component, magnetic field also has a weaker, widely distributed part.  We propose that, at late stages of cluster evolution, turbulence is confined to the wakes of galaxies and infalling mass clumps, and does not fill the volume. However, the area covering factor of the wakes can be close to unity, so most lines of sight pass through a turbulent region. This can help to reconcile the available evidence of turbulence in such clusters as Perseus with the existence of long filaments in the intra\\-cluster gas that seem to be inconsistent with pervasive turbulence.  Since typical lines of sight through a cluster pass through a small number of turbulent cells, synchrotron emission from galaxy clusters can show a detectable degree of polarization at wavelengths 3--6 cm. ",
        "conclusions": ""
    },
    "0512/astro-ph0512052_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) and their afterglows have been proposed as an excellent probe to study the evolution of cosmic star formation, % the reionization of the intergalactic medium, and the metal enrichment history of the universe \\cite{1997ApJ...486L..71T,1998ApJ...501...15M,2000ApJ...536....1L, 2000ApJ...540..687C}, since the prompt $\\gamma$-ray emission of GRBs should be detectable\\cite{2000ApJ...536....1L} out to distances $z>10$. Hitherto, the highest measured redshift for a GRB has been $z=4.50$\\cite{2000A&A...364L..54A}. Here we report the optical spectrum of the afterglow of GRB 050904 obtained 3.4 days after the burst.  The spectrum shows a clear continuum at the long wavelength end of the spectrum with a sharp cutoff at around 9000 \\AA\\ due to Ly$\\alpha$ absorption  at a redshift of 6.3 with a damping wing. Little flux is present in the waveband shortward of the Ly$\\alpha$ break. A system of absorption lines of heavy elements at redshift $z=6.295\\pm 0.002$ were also detected, yielding a precise measurement of the largest known redshift of a GRB. Analysis of the Si~{\\small II} fine structure lines suggest a dense metal-enriched environment around the GRB progenitor, providing unique information on the properties of the gas in a galaxy when the universe was younger than one billion years. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512578_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[t]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}} \\newcommand{\\TabCapp}[2]{\\begin{center}\\parbox[t]{#1}{\\centerline{ \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable} \\vskip2mm \\centerline{\\footnotesize #2}} \\vskip3mm \\end{center}}  \\newcommand{\\MakeTableSepp}[4]{\\begin{table}[p]\\TabCapp{#2}{#3}\\vspace*{-.7cm} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newfont{\\bb}{ptmbi8t at 12pt} \\newfont{\\bbb}{cmbxti10} \\newfont{\\bbbb}{cmbxti10 at 9pt} \\newcommand{\\uprule}{\\rule{0pt}{2.5ex}} \\newcommand{\\douprule}{\\rule[-2ex]{0pt}{4.5ex}} \\newcommand{\\dorule}{\\rule[-2ex]{0pt}{2ex}} \\def\\thefootnote{\\fnsymbol{footnote}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We use the OGLE-II and OGLE-III data in conjunction with the 2MASS near-infrared (NIR) photometry to identify and study Miras and Semiregular Variables (SRVs) in the Large Magellanic Cloud. We found in total 3221 variables of both types, populating two of the series of NIR period--luminosity ({\\it PL}) sequences. The majority of these objects are double periodic pulsators, with periods belonging to both {\\it PL} ridges. We indicate that in the period -- Wesenheit index plane the oxygen-rich and carbon-rich AGB stars from the NIR {\\it PL} sequences C, C$'$ and D split into well separated ridges. Thus, we discover an effective method of distinguishing between O-rich and C-rich Miras, SRVs and stars with Long Secondary Periods using their $V$ and {\\it I}-band photometry. We present an empirical method of estimating the mean $K_s$ magnitudes of the Long Period Variables using single-epoch $K_s$ measurements and complete light curves in the {\\it I}-band. We utilize these corrected magnitudes to show that the O-rich and C-rich Miras and SRVs follow somewhat different $K_s$-band {\\it PL} relations.}{Stars: AGB and post-AGB -- Stars: late-type -- Stars: oscillations -- Magellanic Clouds} ",
        "introduction": "Our knowledge of the Long Period Variables (LPVs) significantly increased when large microlensing surveys collected large enough amount of photometric data to reliably determine the periods and other parameters of these stars. Wood \\etal (1999) showed five parallel sequences in the period--luminosity ({\\it PL}) plane (labeled A--E), each populated by LPVs of different features. As a luminosity in the {\\it PL} diagram Wood \\etal (1999) used reddening free Wesenheit index, but very similar picture was shown for the near-infrared (NIR) $K$ waveband (Wood 2000). Earlier two of these sequences were known: {\\it PL} relation for Miras (Glass and Lloyd Evans 1981) and the additional sequence for Semiregular Variables (SRVs, Wood and Sebo 1996).  The subsequent papers confirmed and extended these results (Cioni \\etal 2001, Noda \\etal 2002, Lebzelter \\etal 2002, Cioni \\etal 2003). Important progress in this field has been made due to the analysis of the photometric data originated in the Optical Gravitational Lensing Experiment (OGLE) conjuncted with the NIR measurements from various sources (2MASS, SIRIUS, DENIS). The OGLE project is a large scale photometric survey regularly monitoring the densest regions of the sky. Collected long-term photometric data of millions of objects are an ideal material for studying a wide variety of variable stars, also red giants.  Kiss and Bedding (2003) used OGLE and 2MASS data to show that the Wood's sequence B above the Tip of the Red Giant Branch (TRGB) is made up with two sequences. Moreover, below the TRGB there are three sequences which are shifted in $\\log{P}$ relative to the {\\it PL} relations above the TRGB. This was the final proof that the majority of the pulsating variables below the TRGB are the first ascent giants (RGB stars). These results were confirmed by Ita \\etal (2004), who labeled by C$'$ the new discovered sequence between the sequences B and C. We adopt this notation in this paper.  Soszy{\\'n}ski \\etal (2004a) showed the complex structure of the {\\it PL} distribution analyzing the OGLE Small Amplitude Red Giants (OSARGs) -- stars which most frequently can be found in the sequences A and B. When the secondary periodicities of these multi-periodic variables are taken into consideration, one can find that these objects lie in the four relatively narrow {\\it PL} sequences. Moreover, Soszy{\\'n}ski \\etal (2004a) empirically showed that the variables fainter than the TRGB are a mixture of the AGB and RGB stars and both groups follow different {\\it PL} relations. Finally, to complicate this pattern, it was shown that the ridge A below the TRGB consist of three closely spaced parallel sequences, what is distinctly visible in the Petersen diagram of these variables.  The aim of the present study is to show that the structure of the sequences C$'$, C and D is also more complex than it looked at the first glance. We discovered that each of these sequences splits into two ridges in the period -- optical Wesenheit index plane. These division corresponds to the spectral separation into oxygen-rich and carbon-rich AGB stars. We use this feature to separate both spectral types of variables, and show that O-rich and C-rich Miras and SRVs follow somewhat different {\\it PL} relations in the $K_s$ waveband.  The paper is organized as follows. In Section~2 we describe the $I$ and {\\it V}-band observations and the cross-correlation of the optical and NIR data. In Section~3 we present the selection of the Miras and SRVs in our data. Section~4 describes the sample and some features of the {\\it PL} relations are presented. In Section~5 we show period -- optical Wesenheit index diagram on which the O- and C-rich AGB stars are well separated. Section~6 contains the description of the mean $K_s$-band magnitudes estimation. In Section~7 we present {\\it PL(K)} diagram for Miras and SRVs. Our results are discussed and summarized in Sections~8 and~9. ",
        "conclusions": "Now, it is possible to explain the existence of the additional two sequences found in the $\\log{P}$--$W_I$ diagram by Soszy{\\'n}ski \\etal (2004ab): the sequence C$''$ spreading between the sequences C and D, and D$'$ occupied the region of longer periods than the sequence D. Both ridges are formed by C-type stars: sequence C$''$ is populated by Mira-like variables and the sequence D$'$ by C-rich stars with the Long Secondary Periods. The C-type SRVs from the sequence C$'$ overlap in the $\\log{P}$--$W_I$ diagram with the M-type Mira-like variables, so this sequence has not been detected earlier.  A careful investigation of the {\\it PL} diagrams gives us some clues concerning the nature of the different {\\it PL} relations in the $\\log{P}$--$K_s$ plane. The number of the obscured C-rich Miras (with $(J-K_s)_0>2.3$~mag) which do not fit the $K_s$-band {\\it PL} relation, fall very close to the mean $\\log{P}$--$W_{JK}$ relation. That can be expected, because the Wesenheit index is the reddening independent quantity.  The more surprising feature is that the C-rich and O-rich giants seem to obey the same $\\log{P}$--$W_{JK}$ relations in both, C and C$'$ sequences. There are two explanations of that fact. First, since the $J$ waveband spectral range of the C-rich AGB stars is strongly affected by the CN and C$_2$ molecular bands (Cohen \\etal 1981), the {\\it J}-band mean luminosities of these objects are on average fainter than the O-rich stars, and the $(J-K_s)$ colors are larger. Thus, in the $\\log{P}$--$W_{JK}$ plane the C-type variables are located relatively higher, and by a coincidence the period--$W_{JK}$ relations of both spectral types are the same.  The second explanation assumes that if the reddening free $\\log{P}$--$W_{JK}$ relations are the same for O-rich and C-rich giants, than the difference between $K_s$-band {\\it PL} relations of both classes is only an effect of dust obscuration, on average larger in the C-type Miras and SRVs. If the second interpretation is correct, our results paradoxically confirm the idea that un-obscured C- and O-rich Miras fit the same $PL(K)$ relation (Feast \\etal 1989).  A completely different phenomenon must be responsible for the split of sequences C and C$'$ in the period--$W_I$ plane. It is known that {\\it V}-band magnitudes of the M-type LPVs are strongly affected by the titanium oxide (TiO) absorption (Smak 1964, Reid and Goldston 2002). Thus, O-rich stars have larger $(V-I)$ colors than C-rich giants of same $K_s$-band luminosity, and in consequence they lie higher (for the same periods) in the period--$W_I$ diagram.  Using our method of discriminating the O- and C-rich AGB stars we can derive the relative numbers of C- to M-type AGB stars (C/M ratio). The C/M ratio is considered to be an indicator of the mean metallicity of a stellar population as well as a tracer of the history of star formation (Cioni \\etal 2005).  \\MakeTable{c@{\\hspace{10pt}}c@{\\hspace{10pt}}c@{\\hspace{10pt}}c @{\\hspace{10pt}}}{12.5cm}{The C/M ratios for stars from the sequences C, C$'$ and D} {\\hline \\noalign{\\vskip3pt} & Sequence C & Sequence C$'$ & Sequence D \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} All & 1.21 & 0.76 & -- \\\\ Above the TRGB & 1.64 & 0.92 & 0.20 \\\\ \\noalign{\\vskip3pt} \\hline }  The C/M ratios for stars in the sequences C, C$'$, and D are shown in Table~1. Because the sequence D undoubtedly contains a fraction of first ascent giants, we present also the C/M values for stars above the TRGB ($K_s<12.05$~mag). It is clear that the largest relative number of C-rich stars occurs in the Mira-like variables, it is intermediate for stars in the sequence C$'$ and smallest for stars with the Long Secondary Periods. We believe that this is an indicator of the evolutionary status of these groups of AGB stars. The vast majority of variables in the sequence D have secondary periods which falls within the four OSARG {\\it PL} sequences (Soszy{\\'n}ski \\etal 2004a). Thus, our results confirm the fact that the longer the period of the LPVs, the more evolutionary advanced the star is.  It is worth to emphasize that, as can be noticed from Table~1, the C/M ratio depends strongly on the method of the AGB stars selection. Different relative numbers of O- and C-rich stars will be obtained when one considers only Miras, Miras and SRVs, or all AGB stars. Likewise, different C/M ratio will occur when one excludes the stars fainter than the TRGB (the bandpass in which the TRGB is determined also matters) or considers all AGB stars. Therefore, the comparison of the C/M in various environments can be reliable only when the same criteria of the AGB stars selection are fulfilled. This may be the source of large uncertainty in calibration of the C/M \\vs [Fe/H] relationship."
    },
    "0512/astro-ph0512264_arXiv.txt": {
        "abstract": "{We present our recent results from numerical simulations of a magnetized flow in the vicinity of a black hole in the context of the collapsar model for GRBs. The simulations show that after an initial transient, the flow settles into a complex convolution of several distinct, time-dependent flow components including an accretion torus, its corona and outflow, an inflow and an outflow in the polar funnel. We focus on studying the nature and connection between these components, in particular between the inflows and related outflows. We find that rotational and MHD effects launch, accelerate, and sustain the outflows. We also find that an outflow can be formed even when the collapsing envelope has initially a very weak magnetic field and a very small angular momentum. Our main conclusion is that even for a relatively weak initial magnetic field and a slow rotation, a gravitational collapse of a stellar envelope can lead to formation  of a very strong and very fast jet.}    \\addkeyword{Accretion, Accretion disks} \\addkeyword{Gamma rays: bursts} \\addkeyword{Methods: numerical} \\addkeyword{MHD} \\addkeyword{Stars: winds, outflows}   \\begin{document}  \\def\\LSUN{\\rm L_{\\odot}} \\def\\MSUN{\\rm M_{\\odot}} \\def\\RSUN{\\rm R_{\\odot}} \\def\\MSUNYR{\\rm M_{\\odot}\\,yr^{-1}} \\def\\MSUNS{\\rm M_{\\odot}\\,s^{-1}} \\def\\MDOT{\\dot{M}}  \\newbox\\grsign \\setbox\\grsign=\\hbox{$>$} \\newdimen\\grdimen \\grdimen=\\ht\\grsign \\newbox\\simlessbox \\newbox\\simgreatbox \\setbox\\simgreatbox=\\hbox{\\raise.5ex\\hbox{$>$}\\llap {\\lower.5ex\\hbox{$\\sim$}}}\\ht1=\\grdimen\\dp1=0pt \\setbox\\simlessbox=\\hbox{\\raise.5ex\\hbox{$<$}\\llap {\\lower.5ex\\hbox{$\\sim$}}}\\ht2=\\grdimen\\dp2=0pt \\def\\simgreat{\\mathrel{\\copy\\simgreatbox}} \\def\\simless{\\mathrel{\\copy\\simlessbox}} ",
        "introduction": "\\label{sec:intro}  Gamma-ray burts (GRBs) are associated with the huge release of energy in a matter of seconds. The collapsar model is one of most promising scenarios to explain these as well as other properties of GRBs (Woosley 1993; Paczy\\'{n}ski 1998; MacFadyen \\& Woosley 1999; Popham, Woosley \\& Fryer 1999;MacFadyen, Woosley \\& Heger 2001; Proga et al. 2003). In this scenario, the collapsed iron core of a massive star accretes gas at a high rate ($\\sim 1 \\MSUNS$) producing a large neutrino flux, a powerful outflow, and a GRB. Many breakthroughs were made in studying GRBs. For example, the association of long duration GRBs with stellar collapse was firmly confirmed (Hjorth et al. 2003, Stanek et al. 2003). Nevertheless, basic properties of the GRB central engine are quite uncertain. This is because the physical conditions in the central engine are extreme and complex (e.g.,  gravitational field is very strong, the temperature and the  mass and energy densities are very high; large scale supersonic/relativistic flows and small scale turbulent flows are physically connected). Additionally, magnetic fields are most likely very important in determining the properties of the central engine.  Effects of magnetic fields in the context of the collapsar model have been studied by a few groups (e.g., Mizuno, Yamada, Koide, \\& Shibata 2004; Mizuno, et al. 2004; De Villiers, Staff, \\& Ouyed 2005). The main focus of these studies is on 2 and 3 dimensional general relativistic magnetohydrodynamic (GR MHD) simulations of jets launched self-consistently from accretion disks orbiting Schwarzschild or Kerr black holes. These are very important studies as they capture a few of the key elements of the central engine and soon may include more elements such as sophisticated equation of state and neutrino physics.  Here, we present and discuss results from 2.5-dimensional, magnetohydrodynamic (MHD) simulations of the collapsar model using pseudo-Newtonian potential. These simulations (see also Proga et al. 2003) are an extension of the work of Proga \\& Begelman (2003, hereafter PB03) who studied MHD accretion flows onto a black hole (BH). In particular, the collapsar simulations include a realistic equation of state, photodisintegration of bound nuclei and cooling due to neutrino emission. The simulations presented here are also an extension of collapsar simulations by MacFadyen \\& Woosley (1999), as they include very similar neutrino physics and initial conditions but are in the MHD instead hydrodynamical (HD) limit. ",
        "conclusions": "Fully 3 dimensional GR MHD simulations are required to capture many of the effects and instabilities of a magnetized fluid in a rotating star collapsing onto a BH . Neutrino physics, a sophisticated equation of state and self-gravity should also be included. However, such full treatment of the MHD collapsar model is beyond reach of the current numerical codes at least for now.  Here we present results from time-dependent two-dimensional MHD simulations of the collapsar model using pseudo-Newtonian potential. The simulations show that: 1) soon after the rotationally supported torus forms, the magnetic field very quickly starts deviating from purely radial due to MRI and shear. This leads to fast growth of the toroidal magnetic field as field lines wind up due to the torus rotation; 2) The toroidal field dominates over the poloidal field and the gradient of the former drives a torus outflow against supersonically accreting gas through the polar funnel; 3) The torus outflow is Poynting flux-dominated; 4) The torus outflow reaches the outer boundary of the computational domain ($5\\times10^8$~cm) with an expansion velocity of 0.2 c; 5) The torus outflow is in a form of a relatively narrow jet (when the jet breaks through the outer boundary its half opening angle is $5^\\circ$); 6) Most of the energy released during the accretion is in neutrinos, $L_\\nu=2\\times 10^{52}~{\\rm erg~s^{-1}}$. Neutrino driving will increase the outflow energy (e.g., Fryer \\& M\\'{e}sz\\'{a}ros 2003 and references therein), but could also increase the mass loading of the outflow if the energy is deposited in the torus. A comparison of the MHD simulations with their HD counterparts show that a strong outflow breaks through a magnetized star much sooner than through a non-magnetized star.  The above conclusions were reached by Proga et al. (2003) and here we confirmed them using a higher resolution simulation. However, we emphasize the fact that the flow settles into a complex convolution of several distinct, time-dependent flow components and the above mentioned torus and its outflow are just two of them. Other flow components include a torus corona and low $l$ flows. A rotationally supported torus and its corona and outflows were extensively studied in the past and remain a focus of many studies. We stress that in the context of the collapsar model, where very low $l$ and high $l$ fluids are present, the situation is more complex. Therefore, future work, even without neutrino physics or effects of general relativity, is important to explore connection and interaction of all the flow components, and their observational implications.  \\begin{figure*} \\begin{picture}(180,590) \\put(-220,323){\\special{psfile=f1.eps hoffset=200 voffset=-150 hscale=110 vscale=110}}            \\end{picture} \\caption{ Sequence of  logarithmic density (top) and  toroidal magnetic field maps (bottom) overplotted with the direction of the poloidal velocity from run A at times 0.153, 0.161, 0.165, and 0.173 s. The sequence illustrates the early phase of the formation of a rotational supported accretion torus and and of a magnetically driven outflow. The length scale is in units of the BH radius (i.e., $r'=r/R_S$ and $z'=z/R_S$). } \\end{figure*}   \\begin{figure*} \\begin{picture}(280,500) \\put(140,500){\\special{psfile=f2a.eps angle=90 hoffset=280 voffset=-250 hscale=100 vscale=100}} \\put(148,250){\\special{psfile=f2b.eps angle=90 hoffset=280 voffset=-250 hscale=100 vscale=100}} \\end{picture} \\caption{ The time evolution of the mass accretion rate (top left panel), total magnetic energy due to each of the three field components (top right panel), neutrino luminosity (bottom left panel) and area-integrated radial Poynting and kinetic flux in the polar outflow at $r=190~R_S$ (bottom right panel) for run A. Formally, we define the polar outflow as the region where $v_r>0$ and $\\beta<1$. Note the difference in the time range in the panel with the radial fluxes. } \\end{figure*}    \\begin{figure*} \\begin{picture}(180,590) \\put(-220,523){\\special{psfile=f3da.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put(-90,523){\\special{psfile=f3db.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put( 40,523){\\special{psfile=f3dc.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put(180,523){\\special{psfile=f3dd.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put(-220,290){\\special{psfile=f3ba.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put(-90,290){\\special{psfile=f3bb.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put( 40,290){\\special{psfile=f3bc.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\put( 180,290){\\special{psfile=f3bd.eps hoffset=200 voffset=-150 hscale=55 vscale=55}}  \\end{picture} \\caption{ Maps of logarithmic density (top) and  toroidal magnetic field maps (bottom) overplotted with the direction of the poloidal velocity from run A at the end of simulations, i.e. time 0.2815 s. The length range increases from to the right. Note that the accretion torus is relatively small (it spans from 1 to about 20 $R_S$). Nevertheless, this tiny torus generates an outflow and mass and energy that can the dynamics and structure of the collapsing star over a large range of radii along the rotational axis. } \\end{figure*}   \\begin{figure*} \\begin{picture}(180,590) \\put(280,423){\\special{psfile=f4.eps angle =90 hoffset=280 voffset=-250 hscale=80 vscale=80}} \\end{picture} \\caption{Radial profiles of various quantities from our run, time-averaged from $0.2629$  through $0.2818$~s. To construct each plot, we averaged the profiles over angle between $\\theta=86^\\circ$ and $94^\\circ$. The top left panel plots the density (solid line) and temperature (dashed line). The top middle panel plots the gas pressure (solid line) and magnetic pressure. The top right panel plots the rotational, radial, Keplerian, and Alfv${\\acute{\\rm e}}$n velocities (solid, dashed, dot-dashed, and dotted line, respectively), as well as the sound speed (triple-dot dashed line). The bottom left panel plots the angular velocity in units of $2c/R_s$. The bottom middle panel plots the Maxwell stress, $\\alpha_{mag}$, and the Reynolds stress, $\\alpha_{gas}$ (solid and dashed line, respectively). We calculate the Reynolds stress using eq. (15) in PB03 and  show only its amplitude. The bottom right panel plots the radial, latitudinal and toroidal components of the magnetic field (dot-dashed, dashed, and solid line, respectively). The length scale is in units of the BH radius (i.e., $r'=r/R_S$).} \\end{figure*}    \\begin{figure*} \\begin{picture}(180,490) \\put(-220,300){\\special{psfile=f5a.eps hoffset=200 voffset=-150 hscale=102 vscale=102}} \\put(-20,300){\\special{psfile=f5b.eps hoffset=200 voffset=-150 hscale=90 vscale=90}} \\end{picture} \\caption{ A map  of logarithmic density overplotted with the direction of the poloidal velocity (left panel) and a contour map of specific angular momentum, $l$ (right panel) from high resolution run B at $t=0.0582$~s. The specific angular momentum is in units of 2$R_S$c. The minimum of $l$ (contour closest to the rotational z-axis) is 0.2, and the contour levels are equally spaced at intervals of $l = 0.2$. The maximum of $l$ is 1.0 and its contour is plotted using dotted curves, whereas all the other contours are plotted using solid curves. This figure shows an inner most part of the flow when a torus just formed and started to develop an outflow. Note that the latter pushed aside the polar funnel accretion flow only below the equator. This figure was chosen to illustrate the complexity of the inner most part of the flow inside a collapsing magnetized star. The upper case letters, in the right panel, mark seven major components of this complex flow: A -- a very low $l$ polar funnel accretion flow; B -- a highly magnetized outflow generated from a low $l$ inflow C; C -- an inflow with angular momentum so low that the flow would accrete directly onto a BH; D -- a rotationally supported, MHD, turbulent accretion torus; E -- a magnetized torus corona; F -- a highly magnetized outflow generated from the torus; and G -- a polar funnel outflow driven thermally (compare with A). See Proga (2005; in preparation) for more details. } \\end{figure*}  \\begin{figure*} \\begin{picture}(180,590) \\put(40,300){\\special{psfile=f6a.eps angle =90 hoffset=200 voffset=-150 hscale=80 vscale=80}} \\put(300,300){\\special{psfile=f6b.eps angle =90 hoffset=200 voffset=-150 hscale=80 vscale=80}} \\end{picture} \\caption{ Maps of logarithmic density (left panel) and toroidal magnetic field (right panel) overplotted with an example of a streamline corresponding to an inflow/outflow for model~C. The maps are also overplotted with the direction of the poloidal velocity and the direction of the poloidal field (the left  and right panels, respectively). See Proga (2005) for more details. } \\end{figure*}"
    },
    "0512/astro-ph0512028_arXiv.txt": {
        "abstract": "Sh\\,2-188 is an example of strong interaction between a planetary nebula (PN) and the interstellar medium (ISM). It shows a single arc-like structure, consisting of several filaments, which is postulated to be the result of motion through the ISM. We present new H$\\alpha$ images from the Isaac Newton Telescope Photometric H$\\alpha$ Survey of the Northern Galactic Plane (IPHAS) which reveal structure behind the filamentary limb. A faint, thin arc is seen opposite the bright limb, in combination forming a closed ring.  Behind the faint arc a long wide tail is detected, doubling the size of the nebula. The nebula extends 15 arcmin on the sky in total. We have developed a `triple-wind' hydrodynamical model, comprising of the initial `slow' asymptotic giant branch (AGB) wind and the later `fast' stellar wind (the interacting stellar wind model), plus a third wind reflecting the motion through the ISM. Simulations at various velocities of the central star relative to the ISM indicate that a high velocity of 125 \\kms\\ is required to reproduce the observed structure. We find that the bright limb and the tail already formed during the AGB phase, prior to the formation of the PN. The closure of the ring arises from the slow--fast wind interaction. Most of the mass lost on the AGB has been swept downstream, providing a potential explanation of the missing mass problem in PNe. We report a proper motion for the central star of $30\\pm10$ mas\\,yr$^{-1}$ in the direction of the bright limb. Assuming the central star is moving at $125\\pm25$ \\kms, the distance to the nebula is estimated to be $850^{+500}_{-420}$ pc, consistent with a spectroscopic distance to the star. Expansion velocities measured from spectroscopic data of the bright filaments are consistent with velocities measured from the simulation. Sh\\,2-188 is one of the largest PNe known, with an extent of 2.8 pc. The model shows that this size was already set during the AGB phase. ",
        "introduction": "The accepted theory of planetary nebula (PN) formation is the interacting stellar winds model (ISW) \\citep{kwok82,balick87} where a fast wind ($\\sim10^3$ \\kms) from the hot central star of a PN blows into the slow wind ($\\sim10$ \\kms) produced during the preceding asymptotic giant branch (AGB) phase.  The inner regions of the slow wind are compressed into a dense shell and ionised by the energetic UV radiation of the central star. The familiar ring-like appearance of PNe is then observed. Structures in the nebula are normally attributed to asymmetries in the slow wind, related to physical properties of the central star, such as rotation or binarity.  Observations of PNe have shown several cases where the outer shell shows the only departure from symmetry. For those cases, the cause of the asymmetries has been proposed to be an interaction with the interstellar medium (ISM). Interaction of PNe with the ISM was first discussed by \\cite{gurzadyan69}. An early theoretical study by \\cite{smith76} assumed a thin shell approximation and the `snowplough' model of \\cite{oort51}. \\cite{isaacmann79} used the same approximation with higher velocities and ISM densities. Both of these studies concluded similarly that a nebula fades away before any disruption of the nebular shell becomes noticeable.  In contrast, \\cite{borkowski90} found that many PNe with large angular extent show signs of PN--ISM interaction, and that all nebulae containing central stars with a proper motion greater than 0.015 arcsec\\,yr$^{-1}$ do so. \\cite{soker91} hydrodynamically modelled the interaction. The PN shell is first compressed in the direction of motion and then in later stages this part of the shell is significantly decelerated with respect to the central star. Both conclude that the interaction with the ISM becomes dominant when the density of the nebular shell drops below a certain critical limit, of typically $n_{\\rm H} = 40 \\rm cm^{-3}$ for a PN in the Galactic Plane.  \\cite{villaver03} (hereafter referred to as VGM) pointed out the PN--ISM interaction had only been studied by considering the relative movement when the nebular shell had already formed.  They performed 2D hydrodynamic simulations following the full AGB phase followed by the PN phase \\citep{vassiliadis93,vassiliadis94}, with a conservative relative velocity of the central star of 20 \\kms\\ and conservative conditions of the surrounding ISM of $n_{\\rm H} =0.1 \\rm cm^{-3}$.  VGM concluded that interaction provides an adequate mechanism to explain the high rate of observed asymmetries in the external shells of PNe. Further, they conclude that stripping of mass downstream during the AGB phase provides a possible solution to the problem of missing mass in PN whereby only a small fraction of the mass ejected during the AGB phase is inferred to be present during the post-AGB phase. Observational evidence for the effect of the ISM on AGB wind stuctures was found by \\cite{zijlstra02}.  The PN Sh\\,2-188 \\citep{sharpless59} is among the most extreme examples of ISM interaction.  The nebula has a one-sided (semi-circular), filamentary appearance.  It is a large nebula, with a reported 340 arcsec diameter \\citep{acker92}; the currrent paper shows it to be considerably larger. It is located in the Galactic plane at $l = 128^\\circ$, $b = -4^\\circ$.  New data are presented showing the faint back of the shell and an extended H$\\alpha$ tail.  The unusual appearance suggests a high proper motion and makes Sh\\,2-188 an important test case for PN--ISM interaction at high velocity.  We have developed a `triple-wind' model using a initial slow AGB wind, a subsequent fast post-AGB (PN) wind, and adding a third wind reflecting the movement through the ISM into the ISW model. We use a hydrodynamic scheme developed by \\cite{wareing05}, to investigate whether this triple-wind model can reproduce the nebular shape of Sh\\,2-188 without requiring magnetic fields. We support the results of the model with a proper motion study of the candidate central star. ",
        "conclusions": "We have successfully reproduced the morphology and the available kinematic data of Sh\\,2-188 and understood its formation in terms of the `triple wind' model. Following the AGB evolution of the central star has been crucial in fitting the whole structure, in particular the tail behind the nebula which is comprised of purely AGB material. The triple-wind model of Sh\\,2-188 has predicted a velocity of the candidate central star of 125 \\kms. Velocities from spectroscopic data on the bright filaments \\citep{rosado82} are in agreement with velocities measured from the simulation. A proper motion study of the central star has shown it to be moving at $30.0 \\pm 10.0$ mas\\,yr$^{-1}$ in the direction of the head of the bright arc. The combination of these two measurements has resulted in estimates of $D = 850^{+500}_{-420}$ pc, $d \\sim 2.5$ pc and $t_{\\rm PN} = 22,500 \\pm 2,500$ years. These estimates are in agreement with the distance and age estimates of \\cite{saurer95} and \\cite{napiwotzki99,napiwotzki01}. The prediction of ISM density in the vicinty of Sh\\,2-188 is also in agreement with the NE2001 model of \\cite{cordes02}. The triple wind model explains the geometric displacement of the central star and indicates that the faint closure of the bright arc is a transitory structure which evolves downstream of the nebula. The PN--ISM interaction has caused $\\approx \\frac{2}{3}$ of the mass expected in the region of the star to be swept downstream providing a solution to the missing mass phenomenon in PN and a valuable way of mixing ISM and stellar material several pc downstream of the central star.  The success of the triple-wind model to fit this extreme object gives confidence in our ability to fit objects with lower speeds. It is now clear that the outer halo structure of PNe contain the effects of ISM interaction and should not be modelled as stand-alone structures. Further, the ISM interaction is an important method of mixing stellar material back into the ISM. The next generation of telescopes, particularly ALMA, will be able to reveal cool dust structure in the universe and shed light on circumstellar AGB material. Further simulations considering temporal evolution of the stellar winds, magnetic fields and/or gravity may shed light on why we have not reproduced the observed fragmentation in Sh\\,2-188."
    },
    "0512/astro-ph0512502_arXiv.txt": {
        "abstract": "{In this paper we test the  astrometric precision of  VLT/FORS2 observations using a serie of CCD frames taken in Galactic bulge area. A special reduction method based on symmetrization of reference fields was used to reduce the atmospheric image motion. Positional precision of  unsaturated  $R=16$~mag star images at 17~sec exposure and 0.55$\\arcsec$ seeing was found to be equal to 300~$\\mu$as. The total error of observations was decomposed into components. It was shown  that  astrometric error  depends largely on the photon centroiding error of the target (250~$\\mu$as for 16~mag stars) while the image motion is much less (110~$\\mu$as). At galactic latitudes to about 20$\\degr$, precision  for a serie of frames with a 10~min total exposure is estimated to be 30--50~$\\mu$as for  14--16 mag stars providing the images are not overexposed and the filter $R_{\\rm special}$ is used. Error estimates for fields with smaller sky star density are given. We conclude that astrometric observations with large telescopes, under optimal reduction, are never atmospheric limited. The bias caused by differential chromatic refraction and residual chromatism of LADC is considered and expressions valid for correcting color effects in the measured positions are given. ",
        "introduction": "The studies of exoplanets and  microlensing effects are usually restricted to measuring  two physical quantities: radial velocities and/or brightness of stars. Angular measurements are rarely considered due to their poor precision. The extraction of many parameters related  to the studied object,  from a single physical quantity, however, is not always possible and reliable. So, radial velocity data allow to find only the low limit of planet mass. The additional information on the star angular displacements essentially improves an accuracy of exoplanet mass and orbital parameters determination (Pravdo \\& Shaklan \\cite{Pravdo}; Benedict et al. \\cite{Benedict}). This  data are  also useful for microlensing studies of massive compact objects in the galaxy because allow to measure lens mass, its distance and velocity (e.g. Boden et al. \\cite{Boden}). Safizadeh et al. (\\cite{Safizadeh}) argue that a combined use of photometric microlensing and astrometric measurements allows detection of low mass planets, determination of their masses and semi-major axes.   To be of practical importance, an accuracy of astrometric measurements for the discussed purposes should be at least of the order of 10 to 100~$\\mu$as. A long history of ground-based observations testifies however that a real accuracy is much worse, about 1~mas per a night for a 10--20$\\arcmin$ reference frame  (Gatewood \\cite{Gatewood}) and 1~mas/h for a 1$\\arcmin$ double star separation measurements (Han \\cite{Han}). The best 150~$\\mu$as/h precision was achieved by Pravdo \\& Shaklan (\\cite{Pravdo}) in a special serie of observations at a 5-m Palomar telescope. A major factor limiting the accuracy is the image motion caused by atmospheric turbulence and displayed as a random relative change of star image positions in unpredicted directions and at uncertain angle. Therefore development of ultra high-precision astrometric methods of observations from the ground  is normally related to infrared interferometers which have 10~$\\mu$as expected precision (e.g. Frink et al. \\cite{Frink}). A similar accuracy, at acceptable exposure times, was found to be unattainable for filled one-aperture telescopes due to   atmospheric image motion (Lindegren \\cite{Lindegren}).  This conclusion was revised in the previous work (Lazorenko \\& Lazorenko \\cite{Lazorenko}, further Paper I) where  we considered the image motion as a turbulent light phase-related quantity dependent on initial wave-front fluctuations at the telescope entrance pupil, on the process of differential measurements and on the way how these phase fluctuations affect the  measured positions. In the spectral domain, transformation of phase fluctuations into image motion is described by four filters of non-atmospheric origin which correspond to: 1) conversion from the phase to the wave-front gradient, 2) averaging over the entrance pupil, 3) averaging due to a finite exposure and 4) formation of function differences in directions to reference stars. Favorable for the inhibition of  atmospheric image motion spectrum is that circumstance that a filter $Y(q)$ correspondent to the averaging over the entrance pupil is a low-pass filter transparent to about $q \\sim 2/D$ spatial frequency ($D$ is the telescope diameter) while reference field stars form a high-pass filter $Q(q)$. The filter bands are only partially overlapping by their rising (for $Q(q)$) and descending (for $Y(q)$) branches. The shape of these instrumental functions can be adjusted so as to minimize the combined system response $Y(q)Q(q)$. Improvement of the  $Y(q)$ filter shape is achieved by apodization, or applying to the entrance pupil a special covering with a variable light transmission. This measure results in a fast attenuation of the filter transmission at $q > 2/D$.  The function $Q(q)$ shape is adjusted at the reduction phase by setting special weights to the each reference star which virtually reconfigures the reference field into a highly symmetric star group and supresses $Q(q)$ responce at low frequencies.  Attenuation factor of the  system with modified filters  $Y(q)$ and $Q(q)$ depends on the telescope diameter, angular reference field radius $R$, the turbulent layer height $h$ and, being a power function of the ratio $Rh/D$ is especially large for narrow-field mode of observations when $Rh \\ll D$. For that reason, the method is recommended for future large 30--100~m telescopes though for existing 8--10~m instruments the expected precision is also high. Based on theoretical asumptions, we had shown that for 10-m  telescopes  precision of differential referencing to the background stars is, depending on the brightness and number of reference stars, from 10 to 60~$\\mu$as at 10~min exposure.  In the current study we test the  effeciency of this astrometric method application to actual observations at VLT and in the presence of numerous noise sources, such as atmospheric image motion, photon noise in the images of stars, pixelization, complex shape of PSF, background noise, optical aberrations, active optics performance etc. In Sect.6 we discuss another atmospheric effect known as differential chromatic   refraction (DCR) and which causes a relative displacement of star images of different colors (Pravdo \\& Shaklan \\cite{Pravdo}; Monet et al. \\cite {Monet}; Louarn et al. \\cite{Louarn}). At VLT, the DCR effect is strongly reduced by use of a special Longitudial Atmospheric Dispersion Compensator (LADC) which mimics atmospheric refraction but introduces it in the opposite direction (Avila et al. \\cite {Avila}).  \\section {Observational data and computation of centroids} For the test study we used a 4-hour serie of  FORS2 frames obtained by Moutou et al. (\\cite{Moutou}) when observing  OGLE-TR-132b object with the filter $R_{\\rm special}$, at  17~sec average exposure and a good FWHM=0.55$\\arcsec $ seeing. The  HR mode with $2 \\times 2$ pixel binning and 0.125$\\arcsec $ pixel size was used. Only master chip of CCD with a field of view of $3.9 \\times 2.1 \\arcmin $ containing about 400 stars  brighter than 20 mag was considered, and only well exposed stars fainter than 16~mag($R$) were measured. Four frames obtained with $V$ filter provided  color information used for examination of chromatic effects (Sect.6).  A full profile fitting was used for computation of  $ \\bar {x}  $, $\\bar {y}  $ coordinates of stellar centroids. The shape of PSF was fitted by a 12 parameter model of a sum  of three elliptic Gaussians with a common  $ \\bar {x}  $, $ \\bar {y}  $ centre: \\begin{equation} \\label{eq:gauss} \\begin{array}{ll} PSF(x,y)= & I G(x,y)+ (x-\\bar{x})^2 I' G'(x,y) + \\\\ & (y-\\bar{y})^2 I'' G''(x,y) \\end{array} \\end{equation} where $x $, $y $ are pixel coordinates. The main Gaussian $G$ with 5 free parameters ($ \\bar {x}$, $\\bar{y}$, gaussian width parameters $\\sigma_{0x}$, $\\sigma_{0y}$,  a term $\\alpha_0$ specifying orientation of semi-axes)  and a flux $I$ (sixth parameter) contains about 98--99\\% of the total star flux; two auxiliary Gaussians  $G ' $ and $G''$ with fluxes $I ' $ and $I'' $ are oriented along  $x $, $y $ axes. As model parameters, they both include $ \\bar {x}$, $\\bar{y}$,  two width terms $\\sigma_{1x}$, $\\sigma_{1y}$ for  $G ' $ and  similar $\\sigma_{2x}$, $\\sigma_{2y}$ terms for  $G''$. Note that since image motion affects photocenter positions, it is useful to treat $ \\bar {x}$, $\\bar{y}$ in  Eq.(\\ref{eq:gauss}) not as the \"fitting model center\" but as the  {\\it  weighted photocenter} defined by equations $\\int \\int (x-\\bar{x}) PSF(x,y) \\, dx \\, dy = \\int \\int (y-\\bar{y}) PSF(x,y) \\, dx \\, dy =0$. This interpretation emphasizes a necessity to use the PSF models with  first derivates on model parameters (except $ \\bar {x}$, $\\bar{y}$) being  symmetric functions of coordinates $x$, $y$. In this case small errors in the parameter determination induce symmetric bias of the PSF shape which does not affect computation of $ \\bar {x}$, $\\bar{y}$. For that reason, the model (\\ref{eq:gauss}) does not contain odd  coordinate powers.   Transition from the continuous  expression for the number of electrons (\\ref {eq:gauss})  to the  discrete pixel counts requires  numerical integration of  PSF$ (x, y) $ within the pixel limits. To avoid this procedure, we used  analytical approximation of the integral precise to  $ 10 ^ {-4} $.  Solution of Eq.(\\ref{eq:gauss}) was obtained in square 10x10~px windows which contained main light signal and are small enough to minimize overlapping of nearby images. Approximation of the actual PSF   shape by the model (\\ref{eq:gauss}) typically is accurate to  5--10\\%, the photon noise statistics essentially differs from the Poissonian. Deviations of the actual PSF shape from the model have a specific wavelike pattern which is highly correlated for all stars images of a certain frame (especially at short spatial scales) but  completely changes at the next frame. At a high level of a light signal (bright images, central pixels), these deviations strongly exceed random photon noise.  Using initial fitting for each star, we averaged the \"measured PSF - model\" residuals over all star images in a frame to find  systematic part of  these residuals as a function var$PSF(x-\\bar{x},y-\\bar{y})$. To obtain this function estimates at subpixel level, it was approximated by a set of 2D cubic polinomials  defined in $3 \\times 3$~pixel areas with centres displaced at 1~pixel. The function var$PSF(x-\\bar{x},y-\\bar{y})$ was then subtracted from individual star images and a final profile fitting applied. This slightly improved  the precision of the centroiding. One may note that the measured PSFs usually show variations at spatial scales of some hundreds of pixels. Therefore it seems reasonable to obtain more precise var$PSF(x-\\bar{x},y-\\bar{y})$ function shapes by performing averaging of the \"measured PSF - model\" residuals only in the limited areas of a frame. Leading to really better fitting of PSFs, this consideration is however incorrect with regards to  determination of  photocenters because application of variable image shape corrections dependent on position of a star in the frame induces a space-variable bias in  the weighted center positions.  An attempt to assign weights to pixels depending on these residual values failed also, though the quality of PSF approximation according to  $ \\chi^2 $ criterion improved. This criterion seems to be improper for precise centroiding works since inspite of low $\\chi^2 \\sim 10$  fitting quality for bright stars, the actual centroiding precision is near the photon limit (Section 4).   To estimate the centroiding precision for  images with PSFs of  complex shape (\\ref {eq:gauss}), we created random images adding Poisson noise to the model function. At 340~e$^-$/px sky level typical for FORS2, the results are  represented by the expression \\begin{equation} \\label{eq:accur} \\varepsilon = \\frac{ \\mbox{FWHM}}{2.26\\sqrt{I}}(1+ 475 I^{-0.7}) \\end{equation} Here FWHM refers to $x $ or $y $ axis and $I $ is given in electrons. Note that Eq.(\\ref {eq:accur}) is valid for random photon fluctuations only and does not take into account a real photon statistics in central pixels. It was found  (Sect.4.1) that  for bright images  a limit (\\ref {eq:accur}) is not achievable. ",
        "conclusions": "This study is a first practical test  of an astrometric method based on symmetrization of reference fields. We have shown that, in contrast to the common belief, differential astrometric observations with large telescopes (prividing an adequate reduction performed) are not atmospheric limited. At optimal reduction, the photon centroid noise from reference field is exactly equal to the atmospheric noise. Considering extra  cetroiding errors from the target star, the photon noise  component {\\it always} dominates over the image motion (Table 3).  This study allows us to predicate that astrometric precision of observations at VLT is 30--50~$\\mu$as for stars of 14--16~mag at  10~min exposure (a session of observations taking about 40 minutes). This precision is sufficiently  good for determination of precise parallaxes and proper motions of stars, examination of microlensing effects and  exoplanets. Of course, the measured parallaxes and proper motions would be relative. Moreover, since each target star position is measured with reference to its peculiar set of background stars, zero-points of parallaxes and proper motions are different for different stars. This circumstance, however, is not critical for the purposes of exoplanet and (in a minor extent) for microlensing studies which require only relative astrometric data. If necessary, zero-points can be easily reduced to a common system by iterative procedure.  The most essential restriction for VLT is the saturation of bright images. Also, sufficiently large light signal necessary for a good referencing can be obtained if observations are restricted to galactic latitudes 20--30$\\degr$. Atmospheric color effects, generally rather large, can be compensated by relevant calibrations based on  star colors and applying restrictions on hour angles. Long-term systematic errors were not considered in this study and may be a source of additional bias.   Existing filled aperture 8--10~m telescopes are powerful astrometric tools due to effective averaging of atmospheric fluctuations of a phase over the aperture and by providing strong light signals. For larger telescopes, in view of dependences $ \\varepsilon \\sim D ^ {-1} $, $ \\sigma _ {\\rm {rf}} \\sim D ^ {-1} $ and $ \\sigma _ {\\rm{at}} \\sim D ^ {-3/2} $ a further increase of the accuracy is expected . Interesting possibilities are  allowed by apodization of an entrance pupil (Paper I) and, in particular, by use of adaptive optics which improves both centroiding precision and removes low-frequency components of image motion spectrum."
    },
    "0512/astro-ph0512444_arXiv.txt": {
        "abstract": "Measurements of molecular hydrogen (\\H2) column densities are presented for the first six rotational levels ($J$=0 to 5) for \\finalnumberofsources\\ extragalactic targets observed with \\FUSE. All of these have a final signal-to-noise ratio larger than \\snlimit, and are located at galactic latitude \\babs$>$20\\deg. The individual observations were calibrated with the \\FUSE\\ calibration pipeline CalFUSE version 2.1 or higher, and then carefully aligned in velocity. The final velocity shifts for all the \\FUSE\\ segments are listed. \\H2\\ column densities or limits are determined for the 6 lowest rotational ($J$) levels for each \\HI\\ component in the line of sight, using a curve-of-growth approach at low column densities ($\\ltsim$16.5), and Voigt-profile fitting at higher column densities. Detections include \\finaldetections\\ measurements of low-velocity \\H2\\ in the Galactic Disk and lower Halo. Eight sightlines yield non-detections for Galactic \\H2. The measured column densities range from log\\,$N$(\\H2)=14 to log\\,$N$(\\H2)=20. Strong correlations are found between log\\,$N$(\\H2) and \\T01, the excitation temperature of the \\H2, as well as between log\\,$N$(\\H2) and the level population ratios (log\\,($N(J^\\prime)/N(J)$)). The average fraction of nuclei in molecular hydrogen ($f$(\\H2)) in each sightline is calculated; however, because there are many \\HI\\ clouds in each sightline, the physics of the transition from \\HI\\ to \\H2\\ can not be studied. Detections also include \\H2\\ in  \\ivdetections\\ intermediate-velocity clouds in the Galactic Halo (out of \\ivlos\\ IVCs). Molecular hydrogen is seen in one high-velocity cloud (the Leading Arm of the Magellanic Stream), although \\hvlos\\ high-velocity clouds are intersected; this strongly suggests that dust is rare or absent in these objects. Finally, there are five detections of \\H2\\ in external galaxies. ",
        "introduction": "Molecular hydrogen (\\H2) is an important constituent of the interstellar medium (ISM), especially in the denser parts, where the presence of \\H2\\ is connected with star formation. In its ground electronic state \\H2\\ has many absorption lines in the ultraviolet part of the spectrum, between about 900 and 1130~\\AA. Using the {\\it Copernicus} satellite, \\H2\\ was detected in many sightlines toward nearby disk stars (Spitzer \\& Jenkins 1975, Savage et al.\\ 1977), showing column densities between \\dex{14} and \\dex{20}~\\cmm2\\ for sightlines with N(\\HI) below $\\sim$5\\tdex{21}~\\cmm2. At higher values of N(\\HI) a large fraction of the hydrogen becomes molecular. The formation of \\H2\\ on interstellar dust particles is theoretically understood (Hollenbach et al.\\ 1971). The transition from mostly neutral to mostly molecular to gas was worked out by Federman et al.\\ (1979). \\par After {\\it Copernicus} there was a long hiatus in the possibility of measuring interstellar \\H2. The {\\it IMAPS} (Jenkins et al.\\ 1988) and {\\it ORFEUS} (Kr\\\"amer et al.\\ 1990) payloads on the {\\it ASTRO-SPAS} space-shuttle platform provided R=\\dex5 and \\dex4 resolution data in the 900--1100\\,\\AA\\ spectral range, although for only a limited number of targets with V magnitudes $\\ltsim$12. The launch of the {\\it Far Ultraviolet Spectroscopic Explorer} (\\FUSE) in 1999 resulted in much better access to the wavelength region containing the \\H2\\ lines. The properties of the \\FUSE\\ satellite have been described in detail by Moos et al.\\ (2000) and Sahnow et al.\\ (2000). The sensitivity of \\FUSE\\ allows observing extragalactic and stellar background targets with fluxes down to about \\dex{-14}~\\fu\\ (corresponding to a V magnitude of about 16). \\par The molecular hydrogen lines in the FUV are associated with transitions from the ground electronic state, X$^1$$\\Sigma_g^+$ to excited electronic states, B$^1$$\\Sigma_u^+$ (Lyman bands) and C$^1$$\\Pi_u$ (Werner bands). The transitions start from the ground ($v$=0) vibrational state and a range of rotational states, $J$. In practice transitions starting with $J$=0 to $J$=5 are observed. The even states ($J$=0, 2, 4, ...) have a nuclear spin $S$=0 (para-\\H2), while the odd states ($J$=1, 3, 5,...) have nuclear spin $S$=1 (ortho-\\H2). Abgrall et al.\\ (1993a, b) presented detailed tables of the wavenumbers and Einstein-A values for all these transitions. However, these are out of date, and this paper instead use the oscillator strengths, wavelengths and damping constants provided by Abgrall to the FUSE PI Team (courtesy of Ken Sembach). \\par The \\H2\\ molecule provides some diagnostics of the physical conditions in the interstellar medium (see e.g.\\ review by Shull \\& Beckwith 1982). The population ratios of the different rotational states can be described by a Boltzmann exponential ($N(J)\\propto \\exp\\left(-E/kT\\right)$). In practice, the Boltzmann temperatures are often just a convenient shorthand for describing the rotational level population ratios. However, the lowest of these temperatures (\\T01) does provide a measure of the collisional excitation. In diffuse clouds, the higher rotational states are also excited by FUV radiation, and the populations of the higher $J$ levels are governed by radiative excitation, followed by UV fluorescence, infrared radiative cascades and collisional de-excitation. These processes are described in several papers that discuss modeling of the population levels of \\H2\\ in the diffuse ISM (see Browning et al.\\ 2003 and many references therein, most usefully Spitzer \\& Zweibel 1974, Jura 1975a,~b). In somewhat denser clouds the higher $J$-levels may also be collisionally excited, for which Gry et al.\\ (2002) present some evidence. \\par Toward almost all Galactic and extragalactic targets observed with \\FUSE\\ Galactic low-velocity \\H2\\ absorption is detected. Published studies include surveys of \\H2\\ in diffuse clouds in the Milky Way, using disk OB stars (Rachford et al.\\ 2002), in interstellar gas toward selected Seyferts and quasars (Sembach et al.\\ 2001a, b), intermediate- velocity clouds (Richter et al.\\ 2003), high-velocity clouds (Richter et al.\\ 2001), the Magellanic Clouds (Tumlinson et al.\\ 2003) and M\\,33 (Bluhm et al.\\ 2003). \\par As of 31 Dec 2004, \\FUSE\\ has observed almost 400 extragalactic background targets (excluding LMC/SMC stars in this count). For 201 of these the signal-to-noise ratio (S/N) per 20~\\kms\\ resolution element near 1030\\,\\AA\\ is $>$3. However, for datasets with S/N$<$\\snlimit, the \\H2\\ becomes difficult to measure, and therefore this paper only includes the \\finalnumberofsources\\ targets with S/N$>$\\snlimit. These targets are located at high galactic latitudes (\\babs$>$20\\deg), and probe the diffuse \\H2\\ in the Galactic Halo and the Disk near the Sun. They were originally observed for many different purposes, including the original survey of AGN sightlines by the \\FUSE\\ Team aimed at measuring Galactic \\OVI\\ absorption, high-velocity clouds and Galactic deuterium. Other targets come from general observer programs in \\FUSE\\ Cycles 1 through 5, which include studies as diverse as searching for intergalactic \\OVI, measuring metals in Galactic high-velocity clouds and analyzing properties of starburst galaxies. The \\H2\\ in these sightlines is often a contaminant for accomplishing the original aim of the observation. In fact, the current survey was started as a way to understand which features of interest might instead be \\H2\\ and to be able to remove the \\H2\\ lines. This paper therefore will concentrate on presenting the data analysis and the \\H2\\ measurements, and only a minimal attempt is made at interpreting the results and discussing the implications for understanding the physical conditions of the Galactic ISM. \\par In \\Sect\\Sobserve\\ the observations are discussed. Section~\\Smeasure\\ presents the methods used to determine N(\\H2), while \\Sect\\Sresults\\ describes the results. ",
        "conclusions": " \\par (1) \\H2\\ is seen in almost all (\\finaldetections\\ out of \\finalnumberofsources) sightlines toward high-latitude (\\babs$>$20\\deg) extragalactic targets. \\par (2) In the northern galactic hemisphere $N$(\\H2) is generally below \\dex{17}\\,\\cmm2\\ at latitudes above 45\\deg. In the southern hemisphere $N$(\\H2) shown no correlation with latitude, but instead is higher for $l$=40 to 150\\deg. \\par (3) There are strong correlations between $N$(\\H2) and the level population ratios, with the average value of \\T01 ranging from 81~K at log\\,$N$(\\H2)=20 to 219~K at log\\,$N$(\\H2)=14. \\par (4) As is theoretically expected, the level populations of the higher rotational states show a relative increase at lower \\H2\\ column densities, with \\T23 and \\T34 rising to 375~K and 575~K at log\\,$N$(\\H2)=15. \\par (5) An unusual cloud is seen at a velocity of +75~\\kms\\ toward Mrk\\,153. This cloud has narrow lines (FWHM$\\sim$8.5~\\kms), and must be small, as it is detected in a 9\\arcmin\\ 21-cm beam, but not in a 36\\arcmin\\ beam. It is detected in \\OI, \\NI, and \\SiII\\ lines. It may be a circumstellar shell, or it may be a low density (100--1\\,\\cmm3), 0.1--10~pc diameter cloud in the Galactic Halo ($D$=0.1--10~kpc) with subsolar abundances. \\par (6) The high-velocity \\H2\\ absorption toward NGC\\,3783 that was reported by Sembach et al.\\ (2001a) is shown to have an \\H2\\ column density of 18.17, rather than 16.80, because the derived FWHM is 8.2~\\kms, rather than the 20~\\kms\\ reported by Sembach et al.\\ (2001a). This factor 25 difference in column densities only leads to noticeable differences in the \\H2\\ absorption profiles for the weakest $J$=2 and $J$=3 lines, which were reported as non-detections by Sembach et al.\\ (2001a). The implication of the higher value for $N$(\\H2) is that the \\H2\\ in the HVC may be in formation-dissociation equilibrium with \\HI\\ after all, rather than being remnant \\H2\\ formed before the tidal interaction. \\par (7) A comparison of the \\H2\\ in IVCs with the results of Richter et al.\\ (2003) shows that in all but three cases similar results are found. Nine new detections of intermediate-velocity \\H2\\ are reported here. \\par (8) Intermediate-velocity \\H2\\ is detected in about half of the cases where intermediate-velocity \\HI\\ is seen, with the fraction increasing toward higher $N$(\\HI). \\par (9) Except for one sightline through the Magellanic Stream (NGC\\,3783) high-velocity \\H2\\ is not detected in the 19 high-velocity \\HI\\ clouds intersected by the 73 sightlines, indicating there is little or no dust in HVCs. \\par (10) Molecular hydrogen is seen in the disks of five external galaxies: NGC\\,4319, M\\,33, NGC\\,625, NGC\\,5236 (=M\\,83), and NGC\\,5253. The detection in NGC\\,4319 is at a distance of 7~kpc from the center of that galaxy, using the Seyfert Mrk\\,205 as a background target, giving log\\,$N$(\\H2)=15.783, and relatively normal Boltzmann temperatures. In the other four cases the detection is the average of many sightlines against OB stars with relatively low extinction, so the low values of $N$(\\H2) are not representative of the ISM in those galaxies."
    },
    "0512/astro-ph0512358_arXiv.txt": {
        "abstract": "GRB~050911, discovered by the {\\it Swift} Burst Alert Telescope, was not seen 4.6~hr later by the {\\it Swift} X-ray Telescope, making it one of the very few X-ray non-detections of a Gamma-Ray Burst (GRB) afterglow at early times. The $\\gamma$-ray light-curve shows at least three peaks, the first two of which ($\\sim$T$_0-0.8$ and T$_0+0.2~\\rm s$, where T$_0$ is the trigger time) were short, each lasting 0.5~s. This was followed by later emission 10--20~s post-burst. The upper limit on the unabsorbed X-ray flux was 1.7~$\\times 10^{-14}~\\rm erg~cm^{-2}~s^{-1}$ (integrating 46~ks of data taken between 11 and 18 September), indicating that the decay must have been rapid. All but one of the long bursts detected by {\\it Swift} were above this limit at $\\sim$4.6~hr, whereas the afterglows of short bursts became undetectable more rapidly. Deep observations with Gemini also revealed no optical afterglow 12~hr after the burst, down to $r=24.0$ (5$\\sigma$ limit). We speculate that GRB~050911 may have been formed through a compact object (black hole-neutron star) merger, with the later outbursts due to a longer disc lifetime linked to a large mass ratio between the merging objects. Alternatively, the burst may have occured in a low density environment, leading to a weak, or non-existent, forward shock -- the so-called `naked GRB' model. ",
        "introduction": "The bimodality in Gamma-Ray Burst (GRB) durations has long been recognised (e.g. Kouveliotou et al. 1993) with the 90\\% $\\gamma$-ray emission interval (T$_{90}$) peaking around 0.3 and 30~s, with a minimum at two seconds. The short duration bursts also typically exhibit systematically harder emission than the longer ones (Kouveliotou et al. 1993). The revolution which transformed the study of long duration bursts via the identification of afterglows and host galaxies at cosmological redshifts has only just reached the short bursts. The recent discoveries of short burst afterglows in several cases (e.g. GRB~050509B: Gehrels et al. 2005, Bloom et al. 2005; GRB~050709: Covino et al. 2005, Fox et al. 2005b, Hjorth et al. 2005a, Villasenor et al. 2005; GRB~050724: Barthelmy et al. 2005b; GRB~050813: Fox et al. 2005c) and the association of these with host galaxies of various morphological types (including ellipticals) indicate that short GRBs have a different origin from the longer duration bursts. They are typically found at lower redshifts (e.g. Bloom et al. 2005; Berger et al. 2005; Tanvir et al. 2005) and have isotropic energies three orders of magnitude below those of the long GRBs. The definitive lack of detection of a supernova related to GRB~050509B (Hjorth et al. 2005b) supports the difference between long and short bursts. The observations to date are in line with what may be expected from GRBs occuring via compact object mergers [neutron star-neutron star (NS-NS) or black hole-neutron star (BH-NS)].  The two burst populations clearly overlap in the hardness duration parameter space and it is interesting to ask what distinguishes the different classes for the cases where classification simply via $\\rm T_{90}$ is ambiguous. At least one of the short GRBs found by {\\it Swift} -- GRB~050724; Barthelmy et al. 2005b -- has softer emission beyond the expected T$_{90}$ of 2~s. In fact, these data showed that the distinction between long and short bursts is partly instrument-dependent. Using {\\it Swift} data, the $\\gamma$-ray light-curve of that burst was found to consist of an initial hard spike (lasting about 0.25~s), followed by another peak at T$_0+1.1~\\rm s$; this section of the light-curve would have classified GRB~050724 as a short burst. However, faint, softer emission was detected by the  {\\it Swift} Burst Alert Telescope (BAT) out to 140~s after the trigger. Simulations showed that BATSE (the Burst And Transient Source Experiment) would not have detected this softer pulse at more than 0.3$\\sigma$, obtaining T$_{90} \\sim 0.43~\\rm s$. Similar behaviour was seen for GRB~050709 (Fox et al. 2005b) and both Lazzati, Ramirez-Ruiz \\& Ghisellini (2001) and Connaughton (2002) investigated such `tails' in BATSE short burst data. Norris \\& Bonnell (2005) have looked at short bursts with extended emission, finding that they can be differentiated from long bursts by having spectral lags consistent with zero for their initial spike emission.   Huang et al. (2005) discuss GRB~040924. The duration of this burst (T$_{90} \\sim 1.2~\\rm s$) places it in the short category, but the authors conclude that it might belong at the short end of the long GRB distribution. The presence of a supernova signature in this burst strengthens their assertion (Soderberg et al. 2005). Likewise, GRB~000301C (Jensen et al. 2001) had $\\rm T_{90}~=~\\rm 2.0~s$, but all the properties of a long duration burst (e.g., starforming host galaxy, optical and radio detections of the afterglow and high redshift). These observations provide evidence that the long-duration population extends to at least 1-2~s.  GRB~050911 also appears to be a candidate for a burst whose classification via $\\rm T_{90}$ is unclear. Although $\\rm T_{90}$ places it in the long burst category, any later X-ray or optical emission decayed rapidly and/or was extremely faint, characteristics more common for short bursts (see, e.g., Gehrels et al 2005, Fox et al. 2005c). This Letter presents the {\\it Swift} and ground-based observations of the burst (Sections~\\ref{swift} and \\ref{followup}), finding that it was only detected in $\\gamma$-rays. Explanations for the lack of X-ray emission are explored in Section~\\ref{disc}. ",
        "conclusions": "No X-ray afterglow emission was detected for GRB~050911, starting 4.6~hr after the burst trigger. Comparison of the upper limit of the emission with other {\\it Swift}-detected bursts (Nousek et al. 2005) demonstrates that any afterglow was at least an order of magnitude fainter than for any other long burst, with the possible exception of GRB~050421. The behaviour could be due to either a BH-NS merger, or by a collapsar occuring in a region of low density, thus forming a `naked' GRB. Whatever mechanism produced GRB~050911, it was an unusual, X-ray dark, burst.  \\vspace{-10pt}"
    },
    "0512/astro-ph0512391_arXiv.txt": {
        "abstract": "We study how the proportion of star-forming galaxies evolves between $z=0.8$ and $z=0$ as a function of galaxy environment, using the \\oii line in emission as a signature of ongoing star formation. Our high-z dataset comprises 16 clusters, 10 groups and another 250 galaxies in poorer groups and the field at $z=0.4-0.8$ from the ESO Distant Cluster Survey, plus another 9 massive clusters at similar redshifts.  As a local comparison, we use samples of galaxy systems selected from the Sloan Digital Sky Survey at $0.04< z < 0.08$. At high-z most systems follow a broad anticorrelation between the fraction of star-forming galaxies and the system velocity dispersion.  At face value, this suggests that at $z=0.4-0.8$ the mass of the system largely determines the proportion of galaxies with ongoing star formation.  At these redshifts the strength of star formation (as measured by the \\oii equivalent width) in star-forming galaxies is also found to vary systematically with environment.  Sloan clusters have much lower fractions of star-forming galaxies than clusters at $z=0.4-0.8$ and, in contrast with the distant clusters, show a plateau for velocity dispersions $\\ge 550 \\rm \\, km \\, s^{-1}$, where the fraction of galaxies with \\oii emission does not vary systematically with velocity dispersion. We quantify the evolution of the proportion of star-forming galaxies as a function of the system velocity dispersion and find it is strongest in intermediate-mass systems ($\\sigma \\sim 500-600 \\, \\rm km \\, s^{-1}$ at z=0). To understand the origin of the observed trends, we use the Press-Schechter formalism and the Millennium Simulation and show that galaxy star formation histories may be closely related to the growth history of clusters and groups. We consider a scenario in which the population of passive galaxies (those devoid of ongoing star formation at the time they are observed) consists of two components: ``primordial'' passive galaxies whose stars all formed at $z>2.5$ and ``quenched'' galaxies whose star formation has been truncated due to the dense environment at later times. We propose a scheme that is able to account for the observed relations between the star-forming fraction and $\\sigma$ in clusters at high- and low-z. If this scenario is roughly correct, the link between galaxy properties and environment is extremely simple to predict purely from a knowledge of the growth of dark matter structures. ",
        "introduction": "The universe as a whole was more actively forming stars in the past than today (Lilly et al. 1996, Madau, Pozzetti \\& Dickinson 1998, Hopkins 2004, Schiminovich et al. 2005).  Studies of galaxies in clusters, groups and in the general field indicate an increased star formation activity at higher redshifts, in all environments. However, a complete mapping of the average star formation activity with redshift as a function of environment has still not been achieved.  A large number of studies, during the last thirty years, have showed that distant clusters generally contain many star-forming galaxies.  In fact, the first evidence for galaxy evolution in clusters, and for galaxy evolution in general, has been the detection of evolution in the star formation activity of cluster galaxies, as revealed by photometry and spectroscopy.  Historically, the higher incidence of star--forming galaxies in distant clusters compared to nearby clusters was first discovered by photometric studies of the proportion of blue galaxies -- the so--called Butcher--Oemler effect (Butcher \\& Oemler 1978, 1984, Smail et al. 1998, Margoniner \\& de Carvalho 2000, Ellingson et al. 2001, Kodama \\& Bower 2001, Margoniner et al. 2001).  In agreement with the photometric results, spectroscopic studies of distant clusters have found significant populations of emission-line galaxies (Dressler \\& Gunn 1982, 1983, Couch \\& Sharples 1997, Dressler \\& Gunn 1992, Couch et al. 1994, Dressler et al. 1999, Fisher et al. 1998, Postman et al. 1998, 2001, Balogh et al. 1997, 1998, Poggianti et al. 1999, Tran et al. 2005, Demarco et al. 2005, Moran et al. 2005 to name a few).  In contrast, nearby rich clusters (such as Coma) generally are ``known'' to have relatively few emission line galaxies.  Increased star formation activity in distant clusters is also indicated by the emission properties of composite cluster-integrated spectra (Dressler et al. 2004).  In parallel to the cluster studies, the fraction of star-forming galaxies has been found to be higher at $z=0.3-0.5$ than at $z=0$ also in groups (Allington-Smith et al. 1993, Wilman et al. 2005b).  While these observations have qualitatively shown that star--forming galaxies were more common in the past than today, {\\it quantifying} this evolution has proved to be very hard. At any given redshift, the properties of cluster galaxies display a large cluster to cluster variance. Disentangling cosmic evolution from cluster--to--cluster variations in a quantitative fashion has not been possible to date due to the relatively small samples of clusters studied in detail at different redshifts.  This difficulty in measuring how the fraction of star--forming galaxies evolves with redshift as a function of the cluster properties has affected all types of studies, photometric and spectroscopic, both those based on the \\oii line from spectroscopic multislit surveys and $\\rm H\\alpha$ cluster-wide studies (Couch et al. 2001, Finn et al. 2004, 2005, Kodama et al. 2004, Umeda et al. 2004). This might be the reason why a quantitative detection of a clear evolution with redshift in the fraction of star-forming galaxies has been elusive so far (Nakata et al. 2005).  Knowing how galaxy properties depend on cluster and group properties at different redshifts is therefore a necessary condition to assess the amount of evolution with redshift, even before attempting to shed some light on how this evolution depends on environment. General trends were soon discovered by the early studies of nearby clusters, such as the fact that richer, more centrally concentrated, relaxed clusters tend to have proportionally fewer star-forming galaxies than less rich, irregular, unrelaxed clusters. However, an exact portrait of how the star formation activity in galaxies depends on the cluster characteristics is still lacking. For example, apparently contrasting results have been found in the literature regarding the presence (Martinez et al. 2002, Biviano et al. 1997, Zabludoff \\& Mulchaey 1998, Margoniner et al. 2001, Goto et al. 2003) or absence (Smail et al. 1998, Andreon \\& Ettori 1999, Ellingson et al. 2001, Fairley et al. 2002, De Propris et al. 2004, Goto 2005, Wilman et al. 2005a) of a relation between galaxy properties and global cluster/group properties such as velocity dispersion, X-ray luminosity and richness.  In this paper we analyze how the fraction of actively star-forming galaxies varies with environment and redshift, comparing samples of clusters and groups at $z=0.4$ to $0.8$ with samples in the local universe. This study is based on the ESO Distant Cluster Survey, a photometric and spectroscopic survey of distant clusters described in \\S2. Deriving the proportion of actively star-forming galaxies as those with \\oii emission in EDisCS and other high-z samples (\\S3) and comparing it with low redshift samples from the Sloan Digital Sky Survey (\\S4), we present how the fraction of star-forming galaxies evolves between $z=0.4-0.8$ and $z=0$ as a function of the cluster/group velocity dispersion (\\S5). In \\S5.3 we discuss the incidence of \\oii emitters in the poorest groups and the field, and in \\S5.4 we show how the distributions of the equivalent widths of \\oii vary with environment. Galaxy systems that strongly deviate from the trends followed by most groups/clusters are discussed in \\S5.5.  Star formation activity and galaxy Hubble types are compared in \\S5.6. Finally, we propose a possible scenario accounting for the observed trends and discuss its major implications in \\S6.  Throughout the paper, line equivalent widths and cluster velocity dispersions are given in the rest frame. We use $H_{0}=70 \\, \\rm km \\, s^{-1} \\, Mpc^{-1}$, $h=H_0/100$, ${\\Omega}_m = 0.3$ and ${\\Omega}_{\\lambda}=0.7$. ",
        "conclusions": "The results found in this paper provide for the first time a quantitative description of the evolution of the star-forming galaxy population in clusters as a function of redshift and $\\sigma$.  These results highlight the need to study the evolution of the star--forming galaxy fraction as a function of system mass.  Ignoring this dependence can lead to incorrect % conclusions regarding the evolution.  % Low-z surveys lack large numbers of massive clusters with $\\sigma \\ge 1000 \\rm \\, km \\, s^{-1}$ (due to their rarity) while most high-z samples include only the most massive clusters. If the cluster mass dependence is not taken into account, % $1000 \\rm \\, km \\, s^{-1}$ high-z clusters  end up being compared with $ 500-700 \\rm \\, km \\, s^{-1}$ low-z clusters, making the evolution harder to detect.  Our results also provide a likely explanation of why it has been so difficult to observe trends with cluster mass/$\\sigma$. The way the data are distributed in Fig.~\\ref{main} already shows that finding the general trends of star-forming fraction with system mass requires: a) a large number of clusters, with data as homogeneous as possible; b) a sample covering a wide range of cluster masses; c) a high quality spectroscopic dataset, both in terms of the number of spectra per cluster and of the quality of the spectra themselves; and d) reliable and thoroughly controlled $\\sigma$'s when using velocity dispersion as a proxy for mass.  The most crucial requirement is the range of cluster/group masses that needs to be explored. For example, from Fig.~\\ref{main} it is evident that sampling at high-z only half of the range in $\\sigma$ (only systems with $\\sigma >$ or $< 700 \\rm \\, km \\, s^{-1}$) would result in the trend being buried in the scatter and unrecognizable. At low z, no trend can be observed when including only systems with $\\sigma > 500 \\rm \\, km \\, s^{-1}$.  This might explain at least some of the contrasting results that have been found in the literature regarding the presence or absence of a relation between galaxy properties and system ``mass'' as determined from velocity dispersion, X-ray luminosity or other global cluster properties.  For example, the relatively limited range in cluster mass explored by X-ray selected samples might be the reason why several works could not find a trend of blue fraction with cluster X-ray luminosity (Smail et al. 1998, Andreon \\& Ettori 1999, Ellingson et al. 2001, Fairley et al. 2002). Only sampling the whole mass range, from groups to massive clusters, trends at high-- and low--z become recognizable.  In this respect, the fact that the mass distribution of EDisCs clusters differs significantly from the distribution of X-ray selected samples at high redshift, extending to much lower masses (Clowe et al. 2005), is an advantage.  Moreover, the range of masses sampled by EDisCS, when evolved to z=0, matches significantly better the mass distribution of nearby clusters than X-ray high-z samples do, as the latter only contain the progenitors of the highest mass tip of the low redshift cluster mass distribution. It is also true that very low-mass/low-richness non-centrally concentrated systems could be under-represented in our sample, given the selection method of EDisCS. This type of systems are generally those with the highest incidence of star-forming galaxies. Hence, though our selection criteria could not be responsible for the \\oii -- $\\sigma$ trend observed, they might influence the observed density of points in the \\oii -- $\\sigma$ diagram. If anything, there should be more high-\\oii low mass groups and the \\oii -- $\\sigma$ relation would then be even stronger than we have observed.   \\subsection{A possible scenario for the trends of the star formation activity as a function of environment}  Understanding the origin of the trends of star formation with velocity dispersion would represent a significant step forward towards comprehending the link between galaxy evolution and environment. If galaxy properties depend on the mass of the system where they reside or have resided during their evolution, there should be a connection between the trends observed and the way cosmological structures have grown in mass with redshift.  In this section we investigate whether the trends in star formation activity correspond in some way to the growth history of structure. To quantify the evolution of cosmological structures, we adopt two different approaches. Using a Press-Schechter formalism (hereafter, PS; Bower 1991, Lacey \\& Cole 1993), we analyze the growth history of systems of different mass. In particular, we study what fraction of the system mass was already in massive structures at different redshifts.  In addition, we use the Millennium Simulation (Springel et al. 2005) to study the growth history in terms of number fraction of galaxies instead of mass, to assess what fraction of the galaxies in systems of a given mass were already in massive structures at different redshifts.  This was computed by populating dark matter haloes with galaxies by means of semi-analytic models (De Lucia et al. 2005, Croton et al. 2005).  We have chosen to employ both approaches because it is important to examine the results both in terms of mass and of number of galaxies. The evolution of the mass of cosmological structures is totally independent of assumptions regarding galaxy formation and evolution, therefore it is not affected by all the uncertainties inherent to these assumptions.  At the same time, it is important to ascertain whether the evolution in the number of galaxies follows the mass evolution, given that the former quantity is the one that is directly observed.  In the following we will name ``clusters'' systems with masses $>10^{14}$ $M_{\\odot}$ and ``groups'' systems with masses between $3 \\times 10^{12} < M <10^{14}$ $M_{\\odot}$. These mass limits approximately correspond to the velocity dispersion limits we have adopted in this paper for defining clusters and groups ($>$ and $< 400 \\, \\rm km \\, s^{-1}$, respectively).  In the comparison between observations and theory we are guided by four considerations:  1) So far, we have focused on the fraction of star-forming galaxies $f_{\\oii}$. At each epoch and in each environment, the fraction of galaxies {\\it with no} ongoing star formation is simply (1-$f_{\\oii}$), and we will refer to these as ``passive galaxies''. Observational studies of clusters suggest that there may be two distinct families of passive galaxies.  \\begin{itemize} \\item The first family is composed of galaxies whose stars all formed at very high redshift ($z>2$) over a short timescale, that have been observed in clusters up to and beyond $z=1$ (Bower et al. 1992, Ellis et al. 1997, Barger et al. 1998, Kodama et al. 1998, van Dokkum et al. 2000,2001, Blakeslee et al. 2003, De Lucia et al. 2004, Barrientos et al., 2004, Holden et al. 2005). This family is largely % composed of luminous cluster ellipticals (e.g. Ellis et al. 1997). \\item The second family corresponds to passive galaxies that have had a more extended period of star formation activity (with a longer star formation timescale). Star formation in these galaxies has been quenched when they were accreted into the dense environment (Dressler \\& Gunn 1983, 1992, Couch \\& Sharples 1987, Balogh et al. 1997, Poggianti et al. 1999, Ellingson et al. 2001, Kodama \\& Bower 2001, van Dokkum \\& Franx 2001, Tran et al. 2003, Poggianti et al. 2004, Wilman et al. 2005b). Most of these galaxies are spirals up to at least 1 Gyr after the star formation is quenched (Dressler et al. 1999, Poggianti et al. 1999, Tran et al. 2003). The passive nature of these galaxies is considered to be a consequence of the interaction with their environment. \\end{itemize} In the following we will refer to these two families as ``primordial passive galaxies'' and ``quenched galaxies'', respectively. As the growth of cosmological structures proceeds and clusters and groups accrete more galaxies, we should expect the relative proportion of the two types of passive galaxies to change. In those environments that efficiently quench star formation, quenched galaxies should progressively become a larger part of the passive population (going to lower redshifts) while primordial passive galaxies should dominate the passive population in systems at high redshift.  2) Observationally, primordial passive galaxies are preferentially located in the densest, more massive structures at all redshifts.  At the epoch when they formed their stars ($z \\ge 2.5$), essentially no system more massive than $10^{14}$ $M_{\\odot}$ existed according to current hierarchical theories. The most massive structures at $z=2.5$ had masses similar to those of systems that at low redshift we would call ``groups'' ($> 3 \\times 10^{12}$).  Thus, when they had just completed their star formation, primordial passive galaxies were in systems of masses comparable to groups today.   3) Considering cluster crossing times (typically 1 Gyr), timescales associated with the various physical processes that might lead to the truncation of the star formation activity (e.g. harassment, ram pressure, strangulation, mergers - 1-2 Gyr at most) and the spectrophometric timescale for the evolution of the \\oii signature ($\\sim 5 \\times 10^7$ yr), a few Gyr should be sufficient for suppressing the \\oii emission in most quenched galaxies.  We will then consider a timescale of 3 Gyr as a reasonable upper limit for the time required to totally extinguish star formation in newly accreted galaxies.  4) The existence of a break-point ($\\sim 550 \\rm \\, km \\, s^{-1}$) in the \\oii -- $\\sigma$ relation observed at low redshift, above which essentially every system has a low \\oii fraction regardless of its mass, suggests that systems above this mass are highly efficient at quenching star formation in galaxies falling into them. A system with a velocity dispersion around $\\sim 500 \\rm \\, km \\, s^{-1}$ at $z=0$ approximately corresponds to a system that 3 Gyr ago (see point (3)) had a mass $\\sim 1-2 \\times 10^{14}$ $M_{\\odot}$ (see Table~4). As a working hypothesis it is then natural to adopt $10^{14}$ $M_{\\odot}$ as the reference mass for efficiently quenching star formation, i.e. the mass above which the quenching is a widespread phenomenon affecting sooner or later (within 3 Gyr according to point 3) all accreted galaxies.  \\begin{figure*}[t] \\vspace{-12truecm} \\centerline{\\includegraphics[width=2.5\\columnwidth]{f15.eps}} \\vspace{-2truecm} \\caption{Clusters at high redshift. The solid line is the line drawn in the \\oii -- $\\sigma$ diagram observed at high redshift. Left. Dots represent the {\\it fraction of galaxies} in 90 haloes selected at $z=0.6$ in the MS simulation that were in haloes with mass $< 3 \\times 10^{12}$ at z=2.5. The dotted line is the prediction of the average {\\it fraction of mass} in haloes $< 3 \\times 10^{12}$ at z=2.5 derived from the Press-Schecter formalism. Middle. Empty dots represent the {\\it fraction of galaxies} in the MS simulation that were in haloes with mass $ < 10^{14}$ $\\sim 3$ Gyr prior to $z=0.6$, thus at z=1.3. The dotted line is the prediction of the {\\it fraction of mass} in haloes $< 10^{14}$ at $z=1.3$ derived from the PS. Right. Empty dots represent the {\\it fraction of galaxies} in the MS simulation that were in haloes with mass $ < 10^{14}$ $\\sim 3$ Gyr prior to $z=0.8$, thus at z=1.75. The dotted line is the prediction of the {\\it fraction of mass} in haloes $< 10^{14}$ at $z=1.75$ derived from the PS. \\label{mr2}} \\end{figure*}    We first consider the family of primordial passive galaxies. In Fig.~\\ref{mr2} we compare the \\oii observations at $z=0.4-0.8$ with the theoretical expectations for the growth history.  The solid line (in both panels) is the line drawn in the \\oii -- $\\sigma$ diagram observed at $z=0.4-0.8$ (Fig.~\\ref{main}).  In the left panel, the dotted line is the fraction of mass of systems at $z=0.6$ that was in systems with $M_{sys}<3 \\times 10^{12}$ $M_{\\odot}$ at z=2.5 as derived from the Press-Schechter formalism, averaged over 100 haloes. Solid dots represent the fraction of galaxies within $R_{200}$ and with $M_V$ limits as for EDisCS that were in systems with $M_{sys}<3 \\times 10^{12}$ $M_{\\odot}$ at z=2.5 obtained for 90 haloes in the Millennium Simulation. Haloes were in this case selected at $z=0.6$ from the MS, similarly to how it was done for haloes at $z=0$ in \\S5.2.   Both the PS results and the MS upper envelope trace remarkably well the \\oii -- $\\sigma$ relation observed at high redshift. The scatter of the MS points illustrates that for systems of any given $\\sigma$ at $z \\sim 0.6$ there is a range in the fraction of galaxies that were already in groups at $z=2.5$.  This scatter indeed resembles the scatter of the datapoints in the observed \\oii--$\\sigma$ diagram of distant clusters (see the left panel of Fig.~\\ref{main}). This figure shows that {\\it the fraction of passive galaxies observed in $z=0.4-0.8$ clusters of a given $\\sigma$/mass is comparable with the fraction of its galaxies (or its mass) that was already in dense environments (=groups) at $z = 2.5$.}  We tentatively identify the latter with the population of primordial passive galaxies as described in point 2) above.  Identifying primordial passive galaxies with galaxies already in groups at $z=2.5$ implies that the great majority of galaxies belonging to environments more massive then $3 \\times 10^{12} \\, M_{\\odot}$ at $z=2.5$ completed their star formation activity at high redshift, and, {\\it vice versa}, that those galaxies that completed their star formation at high redshift are mostly galaxies that were in environments more massive than $3 \\times 10^{12} \\, M_{\\odot}$ at $z=2.5$.  We note that among the 90 haloes extracted from the MS there are also a few ``outliers'' located in the lower left region of Fig.~\\ref{mr2}.  This means that a large fraction of their galaxies resided in haloes of masses $M_{sys}> 3 \\times 10^{12}$ at $z=2.5$. The comparison between Fig.~\\ref{mr2} and Fig.~\\ref{main} then suggests that \\oii outliers at high redshift might represent systems that had a high fraction of their mass already in groups at $z=2.5$ and did not accrete a large population of ``field'' galaxies between $z=2.5$ and $z \\sim 0.6$.  In the middle and right hand panels of Fig.~\\ref{mr2}, we contrast the observed fractions of passive galaxies with the trends expected for quenched galaxies. Open circles in the middle panel show the fraction of galaxies in $z=0.6$ systems in the MS that were in haloes with mass $ < 10^{14}$ $\\sim 3$ Gyr prior to $z=0.6$, thus at $z=1.3$ (see points 3) and 4) above).  For most systems with $\\sigma < 700 \\rm \\, km \\, s^{-1}$ the fraction of galaxies that experienced the cluster environment for at least 3 Gyr is zero, and the predicted trend is inconsistent with the observational results.  {In these systems, the passive galaxy population is not consistent with the fraction of galaxies/mass that was already in systems more massive than $10^{14} \\, M_{\\odot}$ 3 Gyr prior to $z=0.6.$} In the most massive systems ($\\sigma > 700 \\rm \\, km \\, s^{-1}$), the middle panel of Fig.~\\ref{mr2} shows that there is already a considerable fraction of galaxies at $z=0.6$ that have resided in a clusters at least since $z=1.3$. We find that only $\\sim 50$\\% of these galaxies were in groups at z=2.5, therefore there is a non-negligible proportion of galaxies that have experienced the cluster environment (=have been quenched according to point 4)) without being ``primordial'' passive galaxies. This suggests that, while the lower mass systems at $z=0.6$ contain essentially no quenched galaxies, in more massive systems at these redshifts quenched galaxies can already account for more than 1/3 of the passive population.  In the right panel of Fig.~\\ref{mr2}, open circles show the fraction of galaxies in MS systems that were in haloes with mass $ < 10^{14}$ $\\sim 3$ Gyr prior to $z=0.8$ (the highest redshift in the sample we study here), thus at $z=1.75$. Few galaxies have experienced prolonged exposure to cluster-like environments since $z=1.75$. The comparison of the middle and right panel of Fig.~\\ref{mr2} shows that between $z=1.75$ and $z=1.3$ in massive systems there is a dramatic change in the fraction of mass/galaxies that have experienced environments more massive than $10^{14} \\, M_{\\odot}$, hence $z \\sim 1.5$ is an important epoch for the build-up of clusters and the beginning of the quenching process.     Turning to low redshifts, in Fig.~\\ref{mr1} we compare the trends observed in Sloan clusters (solid broken line) with the fraction of mass and galaxies in massive environments at previous redshifts. The left panel shows the fractions of mass and number of galaxies in systems at $z=0$ that were in systems with $M_{sys}<3 \\times 10^{12}$ $M_{\\odot}$ at z=2.5 from the PS and MS (dotted line and filled dots), respectively. For the MS, only galaxies within $R_{200}$ and with $M_V$ limit as for Sloan are considered.  {\\it In contrast to the high-z clusters, the fraction of passive galaxies in systems at z=0 does not agree well with the fraction of galaxies residing in groups already at high redshift. While 80\\% of galaxies in massive systems at z=0 are passive, only 20\\% were in groups at z=2.5.}  The right panel in Fig.~\\ref{mr1} shows the fractions of mass and galaxies that were in systems of mass $M_{sys}< 10^{14}$ $M_{\\odot}$ 3 Gyr prior to the observations (corresponding to z=0.28), from the PS and MS (dotted line and empty dots), respectively. In this case, the agreement between the observations of Sloan clusters (solid broken line) and the PS and MS results is remarkable. This shows that {\\it the observed fraction of passive galaxies in systems with $\\sigma > 500 \\, \\rm km \\, s^{-1}$ at z=0 is compatible with the fraction of galaxies that have resided in a {\\it cluster} ($M_{sys}> 10^{14}$ $M_{\\odot}$) for at least 3 Gyr, and therefore have had the time to have their star formation switched off} (see points 3) and 4) above). The passive population in clusters at z=0 amounts to about 80\\% of all galaxies\\footnote{Within $R_{200}$ and for magnitudes brighter than $M_V=-19.8$.}, {\\it of which} 20\\% (left panel in Fig.~\\ref{mr1}) are ``primordial'' passive galaxies that have evolved passively since $z=2.5$ and 60\\% are galaxies which are ``quenched'' at lower redshift.\\footnote{This is the case because we find that $>90$\\% of the galaxies that were in groups (haloes with masses $M_{sys}> 3 \\times 10^{12}$) at z=2.5 end up being in clusters (haloes with $M_{sys}> 10^{14}$) at z=0.28.  Thus, the population of galaxies in clusters at z=0.28 (right panel of Fig.~\\ref{mr1}) essentially {\\it contains} the galaxy population that was in groups at z=2.5 (left panel of Fig.~\\ref{mr1}).} Also in this case the scatter in the growth history of MS haloes is similar to the observed scatter in the fraction of star-forming galaxies (compare the right panels of Fig.~\\ref{mr1} and Fig.~\\ref{main}). The scatter observed in the \\oii -- $\\sigma$ relation at all redshifts probably simply reflects the scatter in the growth histories of systems of any given mass.  According to this discussion, while the passive galaxy populations of the distant clusters are predominantly composed of primordial passive galaxies, the populations of lower redshift clusters are dominated by quenched galaxies.  We considered whether it is possible to obtain an agreement between the fraction of quenched galaxies and the high-z observations by choosing a lower reference mass for quenching star formation.  However, the reference mass of $ \\sim 10^{14}$ $M_{\\odot}$ is set by the mass (3 Gyr ago) of a system with a velocity dispersion at $z=0$ corresponding to the break observed at $\\sim 500 \\rm \\, km \\, s^{-1}$ in Sloan clusters.  If the minimum mass of a system efficiently quenching star formation were much lower, such as for example $3 \\times 10^{12}$ $M_{\\odot}$, the fraction of passive galaxies would be too high (and the fraction of star-forming galaxies too low) compared to the low-z observations, as shown by the long dashed line in the right panel of Fig.~\\ref{mr1}. Thus, under the assumption that physical processes operate at $z=0.6$ as they do at $z=0$ and adopting the same quenching reference mass at all redshifts, both the primordial and the quenched channels are required to simultaneously match the observed trends at high and low redshift.  \\begin{figure*}[t] \\vspace{-10truecm} \\centerline{\\includegraphics[width=2.2\\columnwidth]{f16.eps}} \\vspace{-1.7truecm} \\caption{Clusters at low redshift. Comparison between the \\oii -- $\\sigma$ relation observed at low redshift (solid broken line) and results from the Millennium Simulation (dots) and from the Press-Schecter formalism (dotted lines). Left. The solid broken line traces the \\oii -- $\\sigma$ relation observed at low redshift. Dots represent the {\\it fraction of galaxies} in 90 haloes in the MS simulation that were in haloes with mass $< 3 \\times 10^{12}$ at z=2.5. The dotted line is the prediction of the {\\it fraction of mass} in haloes $< 3 \\times 10^{12}$ at z=2.5 for a system of a given $\\sigma$ at z=0.0 derived from the Press-Schecter formalism. Right. The solid broken line is repeated from the left panel. Empty dots represent the {\\it fraction of galaxies} in 90 haloes from the MS simulation that were in haloes with mass $< 10^{14}$ $\\sim 3$ Gyr prior to observations, thus at z=0.28. The short dashed line is the prediction of the {\\it fraction of mass} in haloes $< 10^{14}$ at z=0.28 for a system of a given $\\sigma$ at z=0.0 derived from the Press-Schecter formalism. The long dashed line is the prediction of the {\\it fraction of mass} in haloes $< 3 \\times 10^{12}$ at z=0.28 for a system of a given $\\sigma$ at z=0.0 derived from the Press-Schecter formalism. \\label{mr1}} \\end{figure*}  A key point to note from this discussion is that the behaviour of the \\oii fraction with $\\sigma$ at low redshift appears to rule out the possibility that the group environment universally quenches star formation.  If the quenching was a widespread phenomenon in ``groups'' (systems with masses significantly lower than $ \\sim 10^{14}$ $M_{\\odot}$), then all groups and clusters at low redshift should contain much lower fractions of star-forming galaxies than is observed. We note that this does not exclude that star formation might be quenched in {\\it some} galaxies in groups or all galaxies in some of the groups, but a truncation affecting all galaxies within 3 Gyr from infall into systems with masses $ \\ll 10^{14}$ $M_{\\odot}$ cannot be reconciled with the observations. Conversely, the $10^{14}$ $M_{\\odot}$ reference mass indicates that the quenching of the star formation is not limited to very massive clusters, but is highly efficient also in low-mass clusters.  Adopting the scenario depicted above as our working hypothesis, we can address the questions raised in \\S5.2:  a) Why does the proportion of passive/star-forming galaxies correlate/anticorrelate (with a large scatter) with the velocity dispersion of the system for the majority of clusters at z=0.4-0.8?  Our previous discussion shows that the observed fractions of passive galaxies at $z=0.4-0.8$ roughly agree with the fraction of mass/galaxies that were already in groups at $z=2.5$. Primordial passive galaxies make up most of the passive population observed at $z \\sim 0.6$, but in systems more massive than $700 \\rm \\, km \\, s^{-1}$ the proportion of quenched galaxies is already significant.  The anticorrelation observed arises because more massive systems tend to have a higher fraction of their mass/galaxies that were already in groups at z=2.5, and massive systems also have a significant population of quenched galaxies.    b) Why does the proportion of passive/star-forming galaxies evolve with redshift in the way observed? In other words, why is there a Butcher-Oemler effect?  At any redshift, the star-forming population is made up of galaxies that were not in groups at $z>2.5$ {\\it and} were not in clusters in the last few Gyrs. In this scenario the proportion of star-forming galaxies varies with redshift because the proportion that was in groups at $z>2.5$ and the proportion in clusters during the last 3 Gyr change according to the growth history, the sum of the two growing towards lower redshifts.   c) Why is there no clear trend with $\\sigma$ at $z=0$ for systems more massive than $500 \\rm km \\, s^{-1}$?  In clusters with $\\sigma > 500 \\rm km \\, s^{-1}$ at $z=0$, about 80\\% of all galaxies are passive and have resided in clusters for at least 3 Gyr. Of these, 20\\% are primordial passive galaxies that formed in groups at $z>2.5$ and 60\\% are quenched galaxies.  At $z=0$, both the proportion of galaxies that were in groups at $z=2.5$ and the proportion of galaxies that were quenched is flat as a function of the system mass, as shown by Fig.~\\ref{mr1}, and this gives rise to the observed plateau at $\\sigma > 500 \\rm km \\, s^{-1}$. In systems less massive than $400-500 \\rm \\, km \\, s^{-1}$, that are not as efficient as more massive systems in quenching star formation in galaxies infalling into them, the passive population could in some cases still largely coincide with the population of primordial passive galaxies formed at $z>2.5$. However, if all the passive population in low $\\sigma$ systems originated as primordial passive galaxies, systems at low-z on average would have {\\it higher} starforming fractions than similar systems at high-z (compare the left panels of Fig.~\\ref{mr2} and  Fig.~\\ref{mr1}). Hence, it is probable that either the same process active in clusters (but with a lower efficiency) and/or other mechanisms are at work suppressing star formation in some of the galaxies in groups, or in some of the groups. As discussed previously, if the most important factor were density instead of mass, the large scatter of the \\oii fractions at low $\\sigma$ could be due to variations of density for haloes of similarly low masses. As discussed in the next section, the existence of S0 galaxies in groups may be suggesting that star formation is indeed truncated also in groups under certain circumstances.   The consistency between the observations and the theoretical scheme outlined above does not constitute a definite proof of the validity of this scenario; this should be further tested by additional observations, especially at redshifts even higher than those considered here.  It is suggestive, however, to find that the observed star formation trends follow both qualitatively and quantitatively the growth history of structure.  If the scenario we have proposed above approximates the real situation, its implications are far-reaching, as discussed in the next section.   \\subsection{Implications}  In the scenario outlined in the previous section there are two channels that ``produce'' passive galaxies in dense environments. ``Primordial'' passive galaxies form all of their stars at $z>2$ and it is reasonable to largely identify them with ellipticals, while ``quenched'' galaxies have their star formation truncated at much later times, when infalling into an environment that can cause a truncation in the star formation activity, and we tentatively identify them with the population of spirals evolving into S0s. Each one of these two channels seems to correspond to a different {\\it typical mass} of the system. While primordial passive galaxies are related to systems with masses typical of groups at $z>2.5$, quenched galaxies appear to be a universal phenomenon in clusters, i.e. systems with masses $M_{sys}> 10^{14}$ $M_{\\odot}$.  We also note that the galaxy mass and luminosity distributions of primordial passive galaxies and quenched galaxies are expected to differ. Since the star formation activity in galaxies proceeds in a downsizing fashion, both in clusters and in the field (Cowie et al. 1996, Smail et al. 1998, Kodama \\& Bower 2001, Poggianti et al. 2001, Gavazzi et al. 2002, Kauffmann et al. 2003, De Lucia et al. 2004, Poggianti et al. 2004), galaxies terminating their star formation at higher redshift (e.g. primordial passive galaxies) will be on average more massive/luminous than galaxies with a more protracted star formation activity that are quenched at later epochs when they are accreted in the dense environment.  As a consequence, quenched galaxies will be on average less massive/fainter than primordial passive galaxies.  Our results show that galaxy properties could be directly linked with the growth history of DM structure: as shown in Fig.~\\ref{mr2} and ~\\ref{mr1}, the history of the mass of structures is reflected in the star-forming fraction we observe. This suggests that, even without using galaxy formation and evolution models, we can use our knowledge of the growth of structure to explain the trends of galaxy properties in clusters. We have only used two pieces of information, namely how much mass/how many galaxies in a system of a given mass at a given redshift were in dense regions at $z>2.5$, and how much mass/how many galaxies experienced the cluster environment for at least a few Gyrs. If this extremely simple, double-channel picture is generally correct, it represents a very powerful recipe for interpreting the environmental trends observed.  If this scenario approximates the real situation, it can also serve as a key to understand the evolution of galaxy morphologies.  We have seen that the observed environmental trends of galaxy properties originate in two ways. The proportion of ``primordial'' passive galaxies tends to increase with system mass in high-z systems. Systems with proportionally more massive seeds at $z>2.5$ formed more ``primordial'' passive galaxies (mostly ellipticals). In massive systems, other galaxies are added (as S0s) to the passive population as time goes by. These are galaxies that would have continued forming stars had they not been acquired by the dense environment that has switched off their star formation activity.  From a morphological point of view, it is reasonable to associate the ellipticals with the primordial component and {\\it some} of the S0 galaxies with the component quenched at $z<1$.\\footnote{As discussed in many previous studies, probably not all S0 galaxies originate from the quenching of spirals at $z<1$ (e.g. Dressler et al.  1997). The 0-20\\% of S0s observed in clusters at $z=0.4-0.8$ might have originated from spirals evolving into S0s at $z>1$, or by some other mechanism. The existence of S0s in groups (e.g. Hickson, Kindl \\& Auman 1989) shows that this type of galaxies can be produced also in systems less massive than $400-500 \\rm \\, km \\, s^{-1}$.} The fact that in some low-z clusters ellipticals have been found to have only old stellar populations, while a significant fraction of the S0s show signs of a more recent star formation activity is consistent with this scenario (Kuntscher \\& Davies 1998, Poggianti et al. 2001, Smail et al. 2001, Terlevich et al. 2001, Thomas 2002).  This would also explain why some local clusters are dominated by S0 galaxies and some others by ellipticals. Oemler (1974) suggested that elliptical-rich and S0-rich clusters are not two evolutionary stages in cluster evolution, but intrinsically different types of clusters in which the abundance of ellipticals was established at high redshifts. This suggestion was supported by the findings of Fasano et al. (2000) that clusters at $z\\sim 0.1-0.2$ have a low (high) S0/E number ratio if they display (lack) a strong concentration of elliptical galaxies towards the cluster center.  In the scenario we outline above, elliptical-rich clusters would be those with the highest incidence of primordial passive galaxies, and S0-rich clusters those in which quenched galaxies represent a dominant portion of the passive population.  Let us compare spectroscopic and morphological evolution in more detail. The average fraction of ellipticals in clusters at z=0 from Dressler (1980) is $\\sim 20$\\% (see Fasano et al. 2000 or Desai et al. 2006).  In agreement with this , at z=0 the fraction of galaxies in haloes with masses $> 3 \\times 10^{12} \\, \\rm M_{\\odot}$ at $z>2.5$, is $\\sim 20$\\% for the majority of systems with $\\sigma > 300 \\, \\rm km \\, s^{-1}$ (left panel of Fig.~\\ref{mr1}).  In total, the early-type population (ellipticals+S0s) reaches $\\sim 80$\\% at z=0, in agreement with the fraction of passive (non-star-forming) galaxies observed at z=0 (right panel in Fig.~\\ref{main}) and with the fraction of ``passive'' galaxies (primordial+quenched) of haloes at z=0 that were in haloes $> 10^{14} \\, M_{\\odot}$ since z=0.28. The observed late-type fraction at all redshifts is in rough agreement with both the observed fraction of star-forming galaxies and the predicted fraction of galaxies that were not in groups at $z=2.5$ {\\it and} did not reside in a cluster for at least a few Gyrs. All these three quantities are roughly equal to $\\sim 20$\\% in clusters at z=0 (compare Fig.~3 in Desai et al. (2006), the right panel of Fig.~\\ref{main} and Fig.~\\ref{mr1}). All of these three quantities also show a trend with $\\sigma$ at high-z (compare Fig.~7 in Desai et al. (2006), the left panel of Fig.~\\ref{main} and Fig.~\\ref{mr2}).  There is, however, one inconsistency when grossly identifying the population of late-type galaxies with the population of star-forming galaxies. While there is no clear trend in the star-forming fraction with $\\sigma$ at z=0 above $500 \\, \\rm km \\, s^{-1}$ (Fig.~\\ref{main}), the percentage of spiral galaxies in nearby clusters has been shown to anticorrelate with the X-ray luminosity in clusters of $\\sigma \\sim 700-1000 \\rm \\, km \\, s^{-1}$ (Bahcall 1977, Edge \\& Stewart 1991). Although the available morphological studies in X-ray clusters at $z=0$ were not done in a similar way to ours (for selection of members, radial coverage, morphological classifications etc.) and notwithstanding the fact that passive spirals could play an even more important role at low than at high z according to the quenching scenario, this remains an unsolved issue.  Interestingly, significant morphological evolution seems to have taken place in clusters for a large number of galaxies only at $z \\le 0.4$. In fact, as discussed by Desai et al. (2006), the S0 and spirals fractions appear to flatten out at $z>0.45$. A tentative explanation for this behaviour can be found in the scenario proposed.  At redshifts higher than $\\sim 0.6$, the population of passive galaxies is generally dominated by primordial passive galaxies (mostly ellipticals, but also the few S0 galaxies present at high redshift). This is supported by the fact that in clusters at redshifts $z \\sim 0.6$ (and even more so at $z=0.8$) the passive fraction can largely be accounted for by the fraction of galaxies that were in dense environments at $z>2.5$.  Only at $z<0.6$ does the quenched galaxy population become a dominant part of the passive population. Given the delay between the truncation of the star formation and the morphological evolution (Poggianti et al. 1999), this might translate into a morphological evolution observable for a large number of galaxies only at $z<0.4$.\\footnote{Passive spirals may be galaxies that are caught in the transition phase of this transformation. Moreover, this might also explain why the morphology-density relation does not evolve much (except in the very highest density bin) between z=0.5 and z=1 (Smith et al. 2005), see below.  At z=1, the MD (and SFD) relation observed is mostly the ``primordial'' relation as established at very high redshift.}  Thus, the epoch where we can observe the quenching of star formation for a significant fraction of galaxies in clusters is only at $z \\le 0.8$, while the epoch where morphological transformations have taken place for a significant fraction of the cluster galaxies is only at $z \\le 0.4$. The redshift range in which these transformations are observable is due to how the relative infall rate % (fraction of system mass/galaxies) from low-mass/low-density regions onto clusters/groups changes with redshift, as shown in Fig.~\\ref{mr2} and Fig.~\\ref{mr1}. Interestingly, $z \\sim 0.6$ seems to be a special epoch also for the evolution of quasars in rich environment (Yee \\& Ellingson 1993).  So far, we have considered the existence of a trend of the \\oii fraction with $\\sigma$ as a sign of a relation between the star formation of galaxies and the ``global environment'' (mass of the system) in which galaxies reside. However, it is possible that this is a secondary relation induced by the fact that mass and density are closely linked. In fact, density at early and later times might be the driving factor.  The main galaxy properties (star formation activity and morphology) are observed to vary with the ``local'' environment in a systematic way.  The most emblematic way to describe these systematic variations is the morphology-density relation (MD), that is the observed correlation between the frequency of the various Hubble types and the local galaxy density, normally defined as the projected number density of galaxies within an area including its closest 10 neighbours. In clusters in the local Universe, the existence of this relation has been known for a long time: ellipticals are frequent in high density regions, while the fraction of spirals is high in low density regions (Oemler 1974, Dressler 1980).  An MD relation qualitatively similar to the one observed in the local Universe has been observed up to z=1 (as it is logical to expect, galaxy properties correlate with environment at all redshifts), but this relation is {\\it quantitatively} strongly evolving between z=0 and z=0.5: in distant clusters the frequency of S0 galaxies is lower, and the frequency of spirals higher, at all densities (Dressler et al. 1997).  Interestingly, first results at $z=0.7-1.3$ seem to indicate that between z=0.5 and z=1 what changes in the MD relation is only the occurrence of early-type galaxies in the very highest density regions (Smith et al.  2005), and that the frequency of ellipticals at any given local density is the same at z=1 and at z=0 (Postman et al. 2005).  In parallel to the MD relation, there is a star formation-density relation (SFD). For a very long time it has been known that in the nearby Universe also the average star formation activity correlates with the local density: in higher density regions, the mean star formation rate per galaxy is lower.  This is not surprising, given the existence of the MD relation: the highest density regions have proportionally more early-type galaxies devoid of current star formation. The correlation between mean SF and local density extends to very low local densities, comparable to those found at the virial radius of clusters, and such a correlation exists also outside of clusters (e.g. Lewis et al. 2002, Gomez et al. 2003, Kauffmann et al. 2004).  Again, this seems to parallel the fact that an MD relation is probably existing in all environments, and it has been observed in clusters of all types (Dressler et al. 1980), groups (Postman \\& Geller 1984) and cluster outskirts (Treu et al. 2003) -- though the MD relation is {\\it not quantitatively the same} in all environments, being different in concentrated vs. irregular clusters, and high- vs. low-$L_X$ clusters (Dressler 1980, Balogh et al. 2002).   Rephrasing our picture in terms of density instead of mass, both the MD and SFD relations should have a ``primordial'' component and an ``evolved'' component, and both of these components should depend on the environment, but in a different way.  In this scenario, the MD relation and the SFD relation are {\\it established} at very high redshift at the moment the first stars formed in galaxies, and they exist due to the close link between the initial star formation activity of galaxies and the ``primordial'' local density of their environment (ellipticals formed and have always resided in the highest density regions of the Universe).  Thus it could be the ``primordial local density'' at very high redshift that determines the properties (star formation history, morphology - and probably mass, see Steidel et al. 2005) of galaxies formed in that region. Primordial local density and primordial mass of the cluster seed are probably closely related, and the relation we observe with the fraction of mass in massive environments at $z>2.5$ could reflect a relation between the primordial local density and the type of galaxy formed in that region.\\footnote{Outliers in the \\oii--mass($\\sigma$) relation (groups with low \\oii fraction and low masses) could be systems with high primordial local density that have grown in mass much less than the average system with similar primordial density.}  Therefore, the {\\it origin} of the MD and SF relations should be ``primordial'', in the sense that a relation between galaxy properties and environment must have been in place at $z>3$. In fact, a morphological and star formation segregation is an outcome of CDM simulations of large scale structure and semianalytic models because the local density of galaxies and DM is related to the epoch of initial collapse (Bower et al. 1991, Kauffmann 1995a, 1995b, Kauffmann et al. 1999, Benson et al. 2001, Diaferio et al. 2001, Springel et al. 2001): the most massive structures at any epoch are the earliest to collapse. A morphological segregation is built-in at a very fundamental level in hierarchical theories of galaxy formation.  \\begin{table} \\begin{center} \\caption{KS-test probabilities for the color distributions of spectroscopic and photometric catalogs of EDisCS fields to be indistinguishable.\\label{tbl6}} \\begin{tabular}{lrrrrrrrr} \\tableline\\tableline Cluster & ${p}_{corre}^{KS}$ & ${p}_{uncor}^{KS}$ \\\\ \\tableline Cl\\,1018     & 0.368 & 0.250 \\\\ Cl\\,1037     & 0.846 & 0.996 \\\\ Cl\\,1040     & 0.845 & 0.845 \\\\ Cl\\,1054-11  & 0.364 & 0.363 \\\\ Cl\\,1054-12  & 0.246 & 0.158 \\\\ Cl\\,1059     & 0.371 & 0.251 \\\\ Cl\\,1103     & 0.515 & 0.515 \\\\ Cl\\,1119     & 0.251 & 0.164 \\\\ Cl\\,1138     & 1.000 & 1.000 \\\\ Cl\\,1202     & 0.018 & 0.058 \\\\ Cl\\,1216     & 0.245 & 0.245 \\\\ Cl\\,1227     & 0.686 & 0.685 \\\\ Cl\\,1232     & 0.249 & 0.162 \\\\ Cl\\,1301     & 0.685 & 0.515 \\\\ Cl\\,1353     & 0.249 & 0.250 \\\\ Cl\\,1354     & 0.246 & 0.247 \\\\ Cl\\,1411     & 0.518 & 0.368 \\\\ Cl\\,1420     & 0.518 & 0.250 \\\\ \\tableline \\end{tabular}   \\end{center} \\end{table}   However, in addition to this, the MD and SF relations {\\it evolve} with redshift in a way that {\\it depends on environment}. In those environments that are effective in quenching star formation, galaxies coming from lower mass/density environments are transformed by environmental effects when they enter the denser region.  In fact, all models so far have failed to reproduce the S0 population (which, it is worth remembering, represents $>$40\\% of the galaxies in some rich clusters at z=0), recognizing that additional processes seem to be required (Diaferio et al. 2001, Springel et al. 2001, Okamoto \\& Nagashima 2001, 2003, Benson et al. in prep.).  Unfortunately, ``trends with environment'' have often been confused with ``environmental effects'', where the latter is used as a synonym for a physical mechanism switching off star formation in infalling galaxies. Thus, for example, the fact that star formation trends exist down to very low local densities and outside of clusters has often been interpreted in the sense that also such low density environments must somehow ``suppress'' star formation in galaxies. This is not necessarily the case, as discussed at length above.  A trend with environment could be ``imprinted'' very early on simply due for example to the amount of galaxies with a short star formation timescale that were able to form at high redshift in that region.  To fully comprehend why galaxy properties depend on environment in the way it is observed, it is necessary to disentangle high-z ``imprinting'' of the initial conditions from ``proper'' environmental effects acting on galaxies when they experience a dense environment for the first time.  The dependence of galaxy properties on environment does not necessarily arise from a ``suppression'' of star formation: depending on the density/mass of the environment, the relative importance of ``primordial'' and ``quenched'' passive galaxies can vary significantly.   Two main challenges remain at this point. Observationally, the physical mechanism responsible for quenching the star formation still needs to be identified. The characteristic mass of $M_{sys}> 10^{14}$ $M_{\\odot}$ (500 $\\rm km \\, s^{-1}$) suggested by this work may help in discriminating among the various processes, but still does not uniquely pick out a culprit. Our knowledge of how the efficiency of the various physical mechanisms proposed (e.g. harassment, ram pressure, strangulation) depends on the mass of the system is still too poor to draw solid conclusions and discriminate between them. From a theoretical point of view, one of the most useful pieces of information that can come from state-of-the-art simulations is the link between mass and density at primordial and successive times, as well as the relation between the density experienced by a galaxy at different epochs, to assess whether the relations observed with mass are simply the mirror of relations with density."
    },
    "0512/astro-ph0512162.txt": {
        "abstract": "Recent High Energy Stereoscopic System (HESS) observations show that microquasars in high-mass systems are sources of very high energy $\\gamma$-rays.  A leptonic jet model for microquasar $\\gamma$-ray emission is developed.  Using the head-on approximation for the Compton cross section and taking into account angular effects from the star's orbital motion, we derive expressions to calculate the spectrum of $\\g$ rays when nonthermal jet electrons Compton-scatter photons of the stellar radiation field. The spectrum of Compton-scattered accretion-disk radiation is also derived by approximating the accretion disk as a point source of radiation located behind the jet. Numerical results are compared with simpler expressions obtained using $\\delta$-function approximations for the cross sections, from which beaming factors are derived.  Calculations are presented for power-law distributions of nonthermal electrons that are assumed to be isotropically distributed in the comoving jet frame, and applied to $\\gamma$-ray observations of LS 5039. We conclude that (1) the TeV emission measured with HESS cannot result only from Compton-scattered stellar radiation (CSSR), but could be synchrotron self-Compton (SSC) emission or a combination of CSSR and SSC; (2) fitting both the HESS data and the EGRET data claimed to be associated with LS 5039 requires a very improbable leptonic model with a very hard electron energy distribution.  Because the $\\gamma$ rays would be variable in a leptonic jet model, the data sets are unlikely to be representative of a simultaneously measured $\\gamma$-ray spectrum. We therefore attribute EGRET $\\gamma$ rays primarily to CSSR emission, and HESS $\\gamma$ rays to SSC emission.  Detection of periodic modulation of the TeV emission from LS 5039 would favor a leptonic SSC or cascade hadron origin of the emission in the inner jet, whereas stochastic variability alone would support a more extended leptonic model.  The puzzle of the EGRET $\\gamma$ rays from LS 5039 will be quickly solved with GLAST. ",
        "introduction": "X-ray binaries with jets, or microquasars, are common in our Galaxy, with $\\approx 16$ now known \\citep[for a recent review see, e.g.,][]{par05}. About one-third are high-mass X-ray binaries (HMXBs), including Cygnus X-1, Cygnus X-3, LS 5039, and LSI+61$^\\circ$303, and the remainder are low-mass X-ray binaries (LMXBs), including GRS 1915+105, GRO J1655-40, Sco X-1, and 1E 1740.7-2942.  The compact companions are a mixture of black holes and neutron stars, and the radio activity of the microquasars is about equally divided into persistent and transient behaviors.  Recent observations \\citep{aha05} made with the High Energy Stereoscopic System (HESS) show that the high-mass microquasar LS 5039 is a source of very high energy (VHE) $\\gamma$ rays in the $\\approx 200$ GeV -- 10 TeV range, confirming its earlier tentative identification with the EGRET source 3EG J1824-1514 \\citep{par00}. A second high-mass microquasar system, LSI+61$^\\circ$303 (V615 Cas), is associated with the COS-B source 2CG 135+01 \\citep{her77,gt78} and the EGRET source 3EG J0241+6103 \\citep{kni97}, but is too far north for observations with HESS. The EGRET source 3EG J1824-1514 associated with LS 5039 shows marginal evidence for variability \\citep{tor03}.  In contrast, the EGRET light curve of 3EG J0241+6103, the counterpart to LSI+61$^\\circ$303, is strongly variable \\citep{tav98}.  Moreover, \\citet{mas04} performed a timing analysis of the 3EG J0241+6103 data and found a most probable period of $27.4\\pm 7.2$ days, compared to its 26.5 day orbital period.   Evidence for stochastic and periodic variability of these sources at $\\gamma$-ray energies would argue in favor of a leptonic microquasar jet model similar to blazar jet models \\citep[for recent reviews of microquasar models, see][]{rom05,fm04}, especially if the X-ray and $\\gamma$-ray emissions display correlated variability \\citep{gbd05}. The importance of Compton scattering of external photons to produce gamma-rays in blazar jets was first considered by \\citet{bs87} and \\citet{mk89} and later, in view of the {\\it Compton Observatory} discoveries, by \\citet{dsm92} and \\citet{sbr94}. A microquasar jet model differs importantly from a blazar jet model through the addition of the stellar radiation field from the high-mass star and the periodic orbital modulation of the binary stellar system \\citep{gak02,krm02}. Although \\citet{cas05} claim low significance periodic variability when folding the HESS data for LS 5039 with its orbital period, the EGRET data showed no compelling evidence for either stochastic or periodic variability \\citep{par00}. The X-rays from LS 5039 are, however, moderately variable. {\\it RXTE} observations in the 3 -- 30 keV range may show periodic variability correlated with the periastron passage of LS 5039, so it is unceratin whether the X-rays are associated with accretion disk or the jet \\citep{bos05}.  In addition to stellar and accretion-disk emissions, microquasar emission from the jet will produce a variable multiwavelength continuum consisting of radio/IR and jet X-ray \\citep{mff01} synchrotron radiation.  Nonthermal $\\gamma$-ray emission is likely to originate from synchrotron self-Compton (SSC) \\citep{aa99} and external Compton (EC) processes \\citep{lb96,gak02} by these same jet electrons.  The bright high-mass star makes an important contribution to the external radiation field in HMXBs, whereas the accretion disk is the dominant external photon source in LMXBs \\citep{gre05}. Compton-scattering leptonic jet models of the $\\gamma$-ray emission from LS 5039 and LSI+61$^\\circ$303 are presented by \\citet{bp04a,bp04b}.  In this paper, we perform a Compton-scattering analysis of the jet $\\gamma$-ray emission from HMXB microquasars for a leptonic jet model, focusing on Compton-scattering effects from the azimuthal variations of the stellar radiation field using parameters inferred from observations \\citep{mcs01,cas05} of LS 5039, which has a period of $3.90603\\pm 0.00017$ days.  We assume that the twin jets of the microquasar are oriented normal to the orbital plane of the compact object and star; geometrical complications of precessing jets are not considered here.  The orbital variations of the bright O or B stars introduce interesting kinematic variations that appear in the $\\gamma$-ray emission spectrum if the $\\gamma$ rays are due to stellar photons that are Compton-scattered by nonthermal jet electrons, including variations of peak $\\nu F_\\nu$ photon energy and inferences of the locations of the $\\gamma$-ray emission site.  This emission is also subject to the effects of $\\gamma\\gamma$ absorption \\citep{bd05,dub05}, although this effect is not included in the calculations shown here.  Angle-dependent effects on Compton-scattered jet radiation are treated in Section 2.  Approximations made in the derivation are clearly enumerated, so that they can be relaxed in more detailed numerical treatments. In particular, the analysis employs a fixed electron distribution. Spectral calculations are presented in Section 3 using parameters appropriate to LS 5039. Difficulties to fit the combined EGRET and HESS spectra of LS 5039, if assumed to be simultaneously radiated, are discussed in Section 4. Implications for establishing the nature of microquasar $\\gamma$-ray emission from LS 5039 from further HESS and upcoming GLAST observations are also considered. A summary of the results is given in Section 5. ",
        "conclusions": ""
    },
    "0512/astro-ph0512385_arXiv.txt": {
        "abstract": "Recent observational studies of $\\omega$ Centauri by {\\it Hubble Space Telescope} have discovered a double main sequence in the color magnitude diagrams (CMDs) of its stellar populations. The stellar population with the blue main sequence (bMS) is observationally suggested to have a helium abundance much larger, by $\\Delta Y\\sim 0.12$, than that of the red main sequence (rMS). By using somewhat idealized models in which stars of the bMS are formed from gas ejected from those of the rMS, we quantitatively investigate whether the helium overabundance of the bMS can result from  self-enrichment from massive AGB stars, from mass loss of very massive young stars, or from type II supernovae within $\\omega$ Cen. We show that as long as the helium enrichment is due to ejecta from the rMS formed earlier than the bMS, none of the above three enrichment scenarios can explain the observed properties of the bMS self-consistently for reasonable IMFs. The common,  serious problem in all cases is that the observed number fraction of the bMS can not be explained without assuming unusually top-heavy IMFs. This failure of the self-enrichment scenarios implies  that most of the helium-enriched gas necessary for the formation of the bMS originated from other external sources. We thus suggest a new scenario that most of the second generation of stars (i.e., the bMS) in $\\omega$ Cen could be formed from gas ejected from field stellar populations that surrounded $\\omega$ Cen when it  was a nucleus of an ancient dwarf galaxy. ",
        "introduction": "One of the most remarkable results of recent observational studies of $\\omega$ Cen is that it shows a double main sequence (DMS) in the color magnitude diagrams (CMDs) of its stellar populations (e.g., Anderson 1997; Bedin et al. 2004). Bedin et al. (2004) proposed four possible scenarios for the origin of the DMS: (1) Theoretical isochrone models or data calibration are in error, (2) Stars on the bluer main sequence (bMS) of the DMS are a super metal-poor ([Fe/H]$ \\ll -2.0$) population, (3) The bMS represents a very helium-rich ($Y\\ge0.3$) population, and (4) The bMS represents a background stellar population  about $1-2$ kpc behind $\\omega$ Cen. A number of recent investigations have suggested that the above scenario (3) (hereafter referred to as the ``helium pollution scenario'' and abbreviated to ``HEPS'' for convenience) is the most  promising among the four (e.g., Lee et al. 2005; Norris 2004; Piotto et al. 2005).   In the HEPS, there were two epochs of major star formation in the history of $\\omega$ Cen. The first generation of stars  was  formed from proto-globular cluster (GC) cloud(s) to become stars on the red main sequence (rMS) whereas the second generation of stars was formed from gas ejected from the rMS. One of the key questions related to  the HEPS is whether and how the observationally suggested large helium enrichment ($\\Delta Y \\sim 0.12$) for the bMS can be obtained from the rMS with normal $Y$ ($ \\sim 0.24$) in the star formation history of $\\omega$ Cen (e.g., Norris 2004). There are three  candidates for the helium overabundance (Norris 2004; D'Antona et al. 2005; Piotto et al. 2005): (1) AGB stars with the masses larger than $6 {\\rm M}_{\\odot}$, (2) stellar winds associated with massive stars during their early evolutionary phases, and (3) type II supernovae (SNe II). It is however unclear which of the above three is the most reasonable and realistic in the HEPS given a  lack of extensive theoretical studies on the three candidates.    The purpose of this Letter is to investigate the three possibilities in a quantitative manner and thereby discuss which is the most promising as the cause of the bMS (and the DMS) in the context of the HEPS. In this investigation, (1) the possible $Y$ value observationally suggested for $\\omega$ Cen (Norris 2004; Piotto et al. 2005) and (2) the observed  number fraction of the bMS (Bedin et al. 2004) are used to constrain the best possible initial mass function (IMF) of forming stars in the HEPS. We do not intend to discuss extensively the observed abundance pattern of the bMS and the rMS in this paper, because chemical yield tables for AGB stars with helium-rich ejecta ($Y > 0.35$) are not currently available. We will discuss this point in our forthcoming papers (Bekki \\& Norris 2005; BN). Previous theoretical studies demonstrated that if GCs lose more than 50\\% of their initial masses, they will  disintegrate (e.g., Geyer \\& Burkert 2001). We also use this result as a constraint for globular cluster IMF in the HEPS. We demonstrate that none of the above three candidates can explain the above constraint (2) without the modeled $\\omega$ Cen disintegrating. ",
        "conclusions": "We have shown that the observed bMS fraction of $\\omega$ Cen can not be  simply explained by any HEPS in which the bMS was formed from ejecta of the rMS. We accordingly have suggested an ``external pollution'' scenario in which the bMS was formed from gas that was initially within the central region of $\\omega$ Cen's host galaxy. The question yet to be answered in this scenario is how star formation could  proceed within  the rMS when gas was transferred to the central region. It would be possible that the bMS was formed outside yet close to the rMS as a star cluster and then merged with the rMS in the central region of $\\omega$ Cen's host galaxy."
    },
    "0512/astro-ph0512450_arXiv.txt": {
        "abstract": "{ The recently discovered substellar companion to GQ Lup possibly represents a direct test of current planet formation theories. We examine the possible formation scenarios for the companion to GQ Lup assuming it is a $\\sim$2 M$_{Jup}$ object.  We determine that GQ Lup B most likely was scattered into a large, eccentric orbit by an interaction with another planet in the inner system. If this is the case, several directly observable predictions can be made, including the presence of a more massive, secondary companion that could be detected through astrometry, radial velocity measurements, or scuplting in GQ Lup's circumstellar disk. This scenario requires a highly eccentric orbit for the companion already detected. These predictions can be tested within the next decade or so.  Additionally, we look at scenarios of formation if the companion is a brown dwarf.  One possible formation scenario may involve an interaction between a brown dwarf binary and GQ Lup.  We look for evidence of any brown dwarfs that have been ejected from the GQ Lup system by searching the 2MASS all-sky survey. ",
        "introduction": "The recent discovery of a substellar companion to GQ Lup has started an exciting phase of extrasolar planet studies, where direct images of planetary companions can inform substellar spectral models as well as models of planet formation.  The inferred age and distance to the companion, as well as its absolute photometry, is consistent with a lower limit of $\\sim$ 2 M$_{Jup}$ \\citep{neuhauser05}.  The models that were used to determine the lower limit are based on those of \\citet{wuchterl03}.  GQ Lup itself is a 0.7\\Msun\\ star with an age of $\\sim$1 Myr at a distance of 140 pc within the Lupus 1 cloud \\citep{neuhauser05,tachihara96,wichmann98,knude98}.  Since GQ Lup is a young star, it still posseses an unresolved circumstellar disk that nevertheless shows up as a strong mid- and far-infrared excess \\citep{hughes94,weintraub90}.  GQ Lup resides in a relatively isolated star forming region within the Lupus complex of clouds, a string of four dark clouds that holds several small young stellar associations \\citep{hughes94}.  The companion, or GQ Lup B, was observed at a projected separation of 0.7\\arcsec\\ from its host star and has had its common proper motion confirmed.  At a distance of 140~pc, the projected separation is 98~AU.  It has absolute K$_{s}$ and L$^\\prime$ magnitudes of  13.1 and 11.7, corresponding to a luminosity of $\\sim$10$^{-2}$ L$_{\\odot}$ \\citep{neuhauser05}.  Additionally, low resolution spectra of the companion identify it as a $\\sim$L1 type object with a T$_{eff}$ of $\\sim$2000 K.  Based on GQ Lup B's spectral type, low specific gravity, and infrared photometry, \\citet{neuhauser05} concluded that GQ Lup B is most likely a planetary object.  The large separation between GQ Lup and its companion provides a unique test for the formation of planetary systems, since this system does not resemble the Solar System or indeed any of the other known extrasolar planetary systems.  GQ Lup B's origin can provide useful insights into how giant planets form.  In this paper, we endeavor to sketch out the possible origin of GQ Lup B and to determine observational signatures that test whether this formation scenario is correct.  In Section \\ref{s1} we determine the most likely scenario for the formation of GQ Lup B, while in Section \\ref{s2} we determine the observational signatures of this scenario.  Finally in Section \\ref{s3} we present our conclusions. ",
        "conclusions": "\\label{s3} We have shown that the most probable scenarios for the formation of GQ Lup B.  If it is a planet, it was most likely formed in the inner system and ejected outwards.  Such a scenario will require the presence of a second planet in an orbit that will be between 2.5 AU and 10 AU, with significant eccentricity and an observational signature that may be detectable in the near future through radial velocity measurements or astrometry.  If GQ Lup B is a very massive planet or brown dwarf, it is less clear how it formed.  If it formed through an interaction between a brown dwarf binary and GQ Lup, then a nearby brown dwarf should have proper motion consistent with an ejection event.  Finally, it is interesting to speculate on the implications the confirmation of GQ Lup B's planetary status has for the frequency of planets in wide orbits.  If GQ Lup B is a common occurence, we would expect many other discoveries to have already been made.  However, we can get an idea of the upper limit of this frequency by looking at a recent survey for substellar objects in wide orbits \\citep{mccarthy04}.  For objects with masses $\\sim$5 M$_{Jup}$, this survey has found 0/42 stars with planets at separations $>$75~AU or an 80\\% probability that less than 4\\% of stars had planets at those separations.  Assuming that GQ Lup B is 2 M$_{Jup}$, the upper limit to the occurrence of such objects $>$ 75~AU would then be 2.5 times greater assuming the M$^{-1}$ probability distribution of radial velocity surveys holds for this population of planets.  However, more detailed observational work would need to be done to better constrain this frequency.  Since the uncertainty in GQ Lup B's mass is large and the upper limit for the mass is not planetary, the possiblity remains that the object is a brown dwarf.  Our results hold for any widely separated planet discovered.  For example, the planets that are postulated to be present around Formalhut and HR 4796A would imply the presence of another planet that is most likely more massive \\citep{kalas05,wyatt99}. This second planet should be in a closer orbit that is eccentric as well.  Another example could be $\\epsilon$ Eridani, with a confirmed planet at a semi-major axis of $\\sim$3.4~AU and an 0.1 M$_{Jup}$ companion postulated at a semi-major axis of 40~AU based on the sculpting of $\\epsilon$ Eridani's dust disk \\citep{quillen02}."
    },
    "0512/astro-ph0512516_arXiv.txt": {
        "abstract": "The H atoms inside minihalos (i.e. halos with virial temperatures $T_{\\rm vir} \\le 10^{4} {\\rm K}$, in the mass range roughly from $10^{4} M_{\\odot}$ to $10^{8} M_{\\odot}$) during the cosmic dark ages in a $\\Lambda$CDM universe produce a redshifted background of collisionally-pumped 21-cm line radiation which can be seen in emission relative to the cosmic microwave background (CMB). Previously, we used semi-analytical calculations of the 21-cm signal from individual halos of different mass and redshift and the evolving mass function of minihalos to predict the mean brightness temperature of this 21-cm background and its angular fluctuations. Here we use high-resolution cosmological N-body and hydrodynamic simulations of structure formation at high redshift ($z\\gtrsim 8$) to compute the mean brightness temperature of this background from both minihalos and the intergalactic medium (IGM) prior to the onset of Ly$\\alpha$ radiative pumping. We find that the 21-cm signal  from gas in collapsed, virialized minihalos dominates over that from the diffuse shocked gas in the IGM. ",
        "introduction": "\\label{sec:Introduction}  One of the most promising means by which to observe the high redshift universe in the cosmic ``dark ages'' is through the 21-cm wavelength hyperfine transition of the neutral hydrogen that is abundant prior to reionization \\citep[e.g.][]{1990MNRAS.247..510S,1993MNRAS.265..101S}. Motivated by the prospect of new radio telescopes that will be able to observe such a signal, several specific observational techniques have been proposed. Among these are the angular fluctuations on the sky (e.g. \\citealt*{1997ApJ...475..429M};  \\citealt{2000ApJ...528..597T}; \\citealt{2002ApJ...572L.123I} -- ISFM hereafter; \\citealt{2003ApJ...596....1C, 2003MNRAS.341...81I}; \\citealt*{2004ApJ...608..622Z}; \\citealt*{2004MNRAS.347..187F}), features in the frequency spectrum of the signal averaged over a substantial patch of the sky \\citep{1999A&A...345..380S,2004ApJ...608..611G} and studies of absorption features in the spectra of bright, high-redshift radio sources (\\citealt*{2002ApJ...577...22C}; \\citealt{2002ApJ...579....1F,2003AIPC..666...85M}).  For most of these techniques, with the exception of foreground absorption against bright radio sources, the 21-cm signal must be distinguished from the CMB, which is only possible if the 21-cm level population corresponds to a spin temperature $T_{\\rm S}$, which differs from the temperature, $T_{\\rm CMB}$, of the CMB. Since radiative excitation and stimulated emission of this transition by CMB photons tends to drive the value of $T_{\\rm S}$ toward $T_{\\rm CMB}$, some competing mechanism must exist to decouple $T_{\\rm S}$ from $T_{\\rm CMB}$. There are two main physical mechanisms by which the spin temperature is decoupled from the CMB temperature: ``Ly$\\alpha$ pumping,'' or absorption of radiation with a wavelength in the Ly$\\alpha$ transition, followed by decay into one of the hyperfine levels of the ground state (the ``Field-Wouthuysen effect'' -- \\citealt{1952AJ.....57R..31W, 1959ApJ...129..536F}), and spin exchange during collisions between neutral hydrogen atoms \\citep{1956ApJ...124..542P}. The efficiency of Ly$\\alpha$ pumping depends upon the intensity of the UV radiation field at the Ly$\\alpha$ transition, whereas the efficiency of collisional coupling depends upon the local gas density and temperature.  For $z\\gtrsim 150$, these mechanisms are ineffective at decoupling $T_{\\rm S}$ from $T_{\\rm CMB}$, since the kinetic temperature of the gas, $T_{\\rm K}$, is coupled to $T_{\\rm CMB}$ by inverse Compton scattering, and sources of UV radiation have not yet formed to initiate Ly$\\alpha$ pumping. At $z\\lesssim 150$, however, $T_{\\rm K}$ drops below $T_{\\rm CMB}$ and, for $z\\gtrsim 20$, gas at the mean density is sufficiently dense for collisions to couple $T_{\\rm S}$ to $T_{\\rm K}$. During the ``dark ages,'' therefore, when there is no Ly$\\alpha$ pumping, the mean 21-cm signal against the CMB will be zero at $z\\gtrsim 150$, then in absorption at  $20 \\lesssim z\\lesssim 150$.  At lower redshift, collisions become negligible for gas at or below the cosmic mean density, and such gas becomes invisible until its spin temperature is again decoupled from the CMB by Ly$\\alpha$ pumping due to an early UV background from the first stars and quasars. Even though collisional decoupling is ineffective for $z\\lesssim 20$ for gas at the {\\it mean} density, gas in overdense and/or heated regions can still be collisionally-decoupled. In particular, the gas density within ``minihalos'' -- virialized halos of dark and baryonic matter with masses $10^{4} \\lesssim M \\lesssim 10^{8}M_{\\sun}$ and virial temperatures $T<10^{4}K$ which are too low to collisionally ionize their H atoms  -- is sufficiently high so as to decouple its gas spin temperature from the CMB, with $T_{\\rm S} > T_{\\rm CMB}$ in general, causing it to appear in emission (ISFM). ISFM predicted the mean and angular fluctuations of the corresponding 21-cm signal by a semi-analytical calculation which integrated the equation of transfer through individual minihalos of different mass at different redshifts ($z > 6$) and summed these individual halo contributions over the evolving statistical distribution of minihalo masses in the $\\Lambda$CDM universe. \\citet{2003MNRAS.341...81I} extended these results to include non-linear biasing effects. These authors concluded that the fluctuations in intensity across the sky created by minihalos were likely to be observable by the next generation of low-frequency radio telescopes. Such observations could confirm the basic CDM paradigm and constrain the shape and amplitude of the power spectrum at much smaller scales than previously possible.  Since then, \\citet{2004ApJ...611..642F} have suggested that shocked, overdense gas in the diffuse IGM (prior to the onset of Ly$\\alpha$ radiative pumping) is also capable of producing a 21-cm emission signal and that this IGM contribution to the mean signal will dominate over that from gas inside minihalos. Their conclusion is based on an extension of the Press-Schechter approximation (\\citealt{1974ApJ...187..425P}) that is used to determine the fraction of baryons in the diffuse IGM that are hot and dense enough to produce a 21-cm emission signal. We will address this question here.  In order to quantify these effects, we have computed the 21-cm signal both from minihalos and the IGM at $z\\gtrsim 8$ for the first time using high-resolution cosmological N-body and hydrodynamic simulations of structure formation. We have assumed a flat, $\\Lambda$CDM cosmology with matter density parameter  $\\Omega_{m}=0.27$, cosmological constant $\\Omega_{\\Lambda}=0.73$, baryon density $\\Omega_b=0.043$, Hubble constant $H={\\rm 70\\, km\\,s^{-1}Mpc^{-1}}$, linearly-extrapolated $\\sigma_{8h^{-1}}=0.9$ and the ``untilted'' Harrison-Zel'dovich primordial power spectrum.  In this paper, we present detailed, high-resolution gas and N-body simulations which predict the 21-cm signal at $z>6$ due to collisional decoupling from the CMB before the UV background is strong enough to make decoupling due to Ly$\\alpha$ pumping important. Because the Ly$\\alpha$ pumping efficiency is expected to fluctuate strongly until enough sources form to make the efficiency uniform \\citep[e.g.][]{2004ApJ...609..474B}, the results presented here will also be relevant for isolated patches of the universe during reionization itself, which would depend upon the location and abundance of the first sources of UV radiation. Within such regions, we focus on properly resolving the gasdynamics of structure formation at small scales through the use of high resolution gasdynamic and N-body simulations. We test the semi-analytical prediction of the halo model of ISFM for the contribution to the mean signal from gas in minihalos, and investigate the extent to which IGM gas may provide a non-negligible contribution to the total fluctuating signal, as suggested by \\citet{2004ApJ...611..642F}.  These results were first summarized in \\citet{2006NewAR..50..179A}. Here we shall describe our calculations in full and present our results in more detail.  The outline of this paper is as follows. In \\S~\\ref{sec:Calculation} we summarize the basic physics of the 21-cm emission and absorption and the analytical model of ISFM. We also describe our cosmological simulations and their initial conditions, and our method for obtaining the 21-cm signal from our simulations. In \\S~\\ref{sec:Result} we present our results. Our conclusions are summarized in \\S~\\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion}  \\begin{figure}[t] \\plotone{f8.eps} \\caption{ Numerical resolution convergence results. Relative contributions of minihalos and diffuse IGM gas to the total 21-cm background. The top panel shows the results obtained directly from simulations (C1: triangle, long-dashed; C2: square, short-dashed; C3: pentagon, dotted; C4: circle, solid). The bottom panel shows the results which were semi-analytically refined (\\S~\\ref{sub:Improvements}; point- and line-types follow those of the top panel). } \\label{fig-conv_rel} \\end{figure}  We have run a set of cosmological N-body and hydrodynamic simulations of the evolution of dark matter and baryonic gas at high redshift ($6<z<100$). With the assumption that radiative feedback effects from the first light sources are negligible, we calculated the mean differential brightness temperature of the redshifted 21-cm background at each redshift. The mean global signal is in absorption against the CMB above $z\\sim20$ and in overall emission below $z\\sim18$. At $z>20$, the density fluctuations of the IGM gas are largely linear, and their absorption signal is well approximated by the one that results from assuming uniform gas at the mean adiabatically-cooled IGM temperature. At $z<20$, nonlinear structures become common, both minihalos and clumpy, hot, mildly nonlinear IGM, resulting in an overall emission at 21-cm with differential brightness temperature of order a few mK.  By identifying the halos in our simulations, we were able to separate and compare the relative contributions of the halos and the IGM gas to the total signal. We find that the emission from minihalos dominates over that from the IGM outside minihalos, for $z\\lesssim 20$. In particular, the emission from minihalos contributes about $70\\,-\\,75\\,\\%$ of the total emission signal at $z<~17$, peaking at 100\\% at $z\\approx18$, and balancing the absorption by the IGM gas at $z\\approx20$. In contrast, the absorption by cold IGM gas dominates the total signal for $z>20$.  These results appear to contradict the suggestion by \\citet{2004ApJ...611..642F}, that the 21-cm emission signal would be dominated by the contribution of shocked gas in the diffuse IGM. They used the Press-Schechter formalism to estimate the fraction of the IGM outside of minihalos, which is shock-heated, by adopting a spherical infall model for the growth of density fluctuations and assuming that all gas inside the turn-around radius is shock-heated. This method is apparently not accurate enough to describe the filamentary nature of structure formation in the IGM.  On the other hand, our results are consistent with the analytical estimates of the mean 21-cm emission signal from minihalos by ISFM. This indicates that the statistical prediction of the collapsed and virialized regions identified as minihalos by the Press-Schechter formalism (or its refinement in terms of the ST formula), with virial temperatures $T < 10^4 {\\rm K}$, with halos characterized individually by the TIS model, is a reasonably good approximation for the mean 21-cm signal for minihalos at all redshifts and a good estimator even for the total mean signal including both minihalos and the diffuse IGM at $z \\lesssim 20$. This encourages us to believe that the angular and spectral fluctuations in the 21-cm background predicted by ISFM based on that model will also be borne out by future simulations involving a much larger volume than was simulated here. The current simulation volume is too small to be used to calculate the fluctuations in the 21-cm background because current plans for radio surveys to measure this background involve beams which will sample much larger angular scales ($>$ arcminutes) than are subtended by our current box ($\\Delta \\theta_{\\rm box} \\sim 0.2' (1+z)_{20}^{0.2}$, where $(1+z)_{z'}=(1+z)/(1+z')$) and bandwidths ($\\sim $ MHz) which are too large to resolve the depth of our simulation box in redshift-space (i.e. $\\Delta \\nu_{\\rm box} \\sim 40 {\\rm kHz} (1+z)_{10}^{-1/2} \\left[L/(0.5 h^{-1} {\\rm Mpc}) \\right] $ ). According to ISFM and \\citet{2003MNRAS.341...81I}, the fluctuations in the 21-cm background from minihalos are significantly enhanced by the fact that minihalos are biased relative to the total matter density fluctuations. A larger simulation volume than ours will also be necessary to sample this minihalo bias in a statistically meaningful way. This bias is likely to affect the minihalo contribution to the 21-cm background fluctuations substantially more than it does the diffuse IGM contribution, thereby boosting the relative importance of minihalos over the IGM even above the ratio of their contributions to the mean signal.  We have considered the limit in which only collisional pumping is available to decouple the spin temperature from that of the CMB, and sources of radiative pumping have not yet emerged to compete with this process. The possibility exists, however, that an X-ray background built up as sources formed inside some halos, which heated the IGM while only partially ionizing it (e.g. \\citealt{2003MNRAS.346..456O}). This heating might have boosted the kinetic temperature of the IGM and enhanced the effect of collisional pumping there (e.g. \\citealt{2004ApJ...602....1C})\\footnote{Recently, \\citet*{2006ApJ...637L...1K} considered the X-rays emitted by an early miniquasar, finding that such an X-ray source can heat the IGM to as much as a few thousand degrees Kelvin without ionizing it. This boosts the 21-cm signal from collisionally-decoupled gas in the diffuse IGM significantly. Their calculations neglect the ionizing UV radiation which might also be released by the miniquasar and its stellar progenitor, as well as the Ly$\\alpha$ pumping they might contribute.}. Such X-ray heating would also have raised the minimum mass of minihalos which formed thereafter, filled with their fair share of neutral H atoms. When stellar sources began to form and build up the UV background at energies below the Lyman limit of hydrogen, Ly$\\alpha$ pumping could then have radiatively coupled $T_{\\rm S}$ to $T_{\\rm K}$, as well. The same sources presumably emitted UV radiation above the H Lyman limit, too, which ionized both the IGM and the minihalos within the HII regions surrounding these sources (e.g. \\citealt*{2004MNRAS.348..753S}; \\citealt*{ISR05}; \\citealt*{ISS05}). Such HII regions would have created holes in the 21-cm background, which then originated only in the remaining neutral regions. Minihalos could have lost the neutral hydrogen gas responsible for their 21-cm emission, not only by ``outside-in'' photoionization by an external source, but also by ``inside-out'' photoionization by internal Pop III star formation (e.g. \\citealt{2004ApJ...613..631K}; \\citealt*{2006ApJ...639..621A}). The $\\rm H_2$ formation required for minihalos to form stars, however, is likely to have been suppressed easily by photodissociation in the Lyman-Warner bands by the background of UV radiation created by the very first minihalos which formed stars, when the ionizing radiation background was still much too low to cause reionization (\\citealt*{2000ApJ...534...11H}). In that case, most minihalos would have remained intact until they were ionized from without. In the future, we plan to improve upon the current calculation by incorporating this more complicated physics. We also intend to run simulations with larger simulation boxes. This would allow us to  predict the 21-cm fluctuation signal (e.g. ISFM) and determine whether the relative contribution of minihalos to the total signal, which we find to be about 70 -- 75 \\% at $z \\lesssim 20$ for the mean signal, varies as the mean signal fluctuates."
    },
    "0512/hep-ph0512090_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Dark Matter (DM) problem calls for new physics beyond the Standard Model (SM). Its simplest interpretation consists in assuming that DM is the thermal relic of a new stable neutral particle with mass $M\\sim T_{0}^{1/2}G_{\\rm N}^{-1/4} \\sim\\TeV$ where $T_{0}\\sim 3\\,{\\rm K}$ is  the present temperature of the universe and $G_{\\rm N}$ is the Newton constant. Attempts to address the Higgs mass hierarchy problem typically introduce a rich amount of new physics at the weak scale, including DM candidates; supersymmetry is widely considered as the most promising proposal~\\cite{DM}. However (i) no new physics appeared so far at collider experiments: the simplest solutions to the hierarchy problem start needing uncomfortably high fine-tunings of the their unknown parameters~\\cite{FT}; (ii) the presence of a number of unknown parameters (e.g.\\ all sparticle masses) obscures the phenomenology of the DM candidates; (iii) the stability of the DM candidates is the result of extra features introduced by hand (e.g.\\  matter parity).  We here explore an opposite, minimalistic approach: focussing on the Dark Matter problem, we add to the Standard Model (SM) extra multiplets ${\\cal X}+\\hbox{h.c.}$ with minimal spin, isospin and hypercharge quantum numbers, and search for the assignments that provide most or all of the following properties: \\begin{enumerate}  \\item The lightest component is automatically stable on cosmological time-scales.   \\item The only renormalizable interactions of ${\\cal X}$ to other SM particles are of gauge type, such that new physics is determined by one new parameter:  the tree-level mass $M$ of the Minimal Dark Matter (MDM) multiplet.  \\item  Quantum corrections generate a mass splitting $\\Delta M$ such that the lightest component of $\\cal X$ is neutral. We compute the value of $M$ for which the thermal relic abundance equals the measured DM abundance.   \\item The DM candidate is still allowed by DM searches. \\end{enumerate} In section~\\ref{list} we list the possible candidates. In section~\\ref{splitting} we compute the mass splitting. In section~\\ref{Omega} we compute the thermal relic  abundance of ${\\cal X}$ and equate it to the observed DM abundance, inferring the DM mass $M$. In section~\\ref{DMexp} we discuss signals and constraints from DM experiments. In section~\\ref{coll} we discuss collider signals. Section~\\ref{concl} contains our conclusions and a summary of the results.    \\begin{table}[t] $$\\begin{array}{|ccc|ccccc|}\\hline \\multicolumn{3}{|c|}{\\hbox{Quantum numbers}} &\\hbox{DM can}&\\!\\!\\!\\hbox{DM mass}&m_{{\\rm DM}^\\pm} - m_{\\rm DM}\\!\\!\\!& \\hbox{Events at LHC} & \\hbox{$\\sigma_{\\rm SI}$ in} \\\\ \\SU(2)_L &\\! {\\rm U}(1)_Y\\! &\\hbox{Spin} & \\hbox{decay into} & \\hbox{in TeV} &\\hbox{in MeV} & \\hbox{$\\int {\\cal L}\\,dt = $100/fb} & \\hbox{$10^{-45}\\,{\\rm cm}^2$}\\\\ \\hline \\hline 2 & 1/2 & 0 & EL & 0.54 \\pm 0.01 & 350 & 320\\div510 & 0.2\\\\ 2 & 1/2 & 1/2 & EH & 1.1 \\pm 0.03 & 341 & 160\\div330 & 0.2\\\\ \\hline 3 & 0 & 0 & HH^* & 2.0 \\pm 0.05 & 166 & 0.2\\div1.0 & 1.3\\\\ 3 & 0 & 1/2 & LH & 2.4 \\pm 0.06 & 166 & 0.8\\div4.0 & 1.3\\\\ 3 & 1 & 0 & HH,LL & 1.6 \\pm 0.04 & 540 & 3.0\\div10 & 1.7\\\\ 3 & 1 & 1/2 & LH & 1.8 \\pm 0.05 & 525 & 27\\div90 & 1.7\\\\ \\hline 4 & 1/2 & 0 & HHH^* & 2.4 \\pm 0.06 & 353 & 0.10\\div0.6 & 1.6\\\\ 4 & 1/2 & 1/2 & (LHH^*) & 2.4 \\pm 0.06 & 347 & 5.3\\div25 & 1.6\\\\ 4 & 3/2 & 0 & HHH & 2.9 \\pm 0.07 & 729 & 0.01\\div0.10 & 7.5\\\\ 4 & 3/2 & 1/2 & (LHH) & 2.6 \\pm 0.07 & 712 & 1.7\\div9.5 & 7.5\\\\ \\hline 5 & 0 & 0 & (HHH^*H^*) & 5.0 \\pm 0.1 & 166 & \\ll1 & 12\\\\ 5 & 0 & 1/2 & - & 4.4 \\pm 0.1 & 166 & \\ll1 & 12\\\\ \\hline 7 & 0 & 0 & - & 8.5 \\pm 0.2 & 166 & \\ll1 & 46\\\\  \\hline \\end{array}$$ \\caption{\\em\\label{tab:1} {\\bf Summary of the main properties of Minimal DM candidates}. Quantum numbers are listed in the first 3 columns; candidates with $Y\\neq 0$ are allowed by direct DM searches only if appropriate non-minimalities are introduced. The 4th column indicates dangerous decay modes, that need to be suppressed (see sec.~\\ref{list} for discussion). The 5th column gives the DM mass such that the thermal relic abundance equals the observed DM abundance (section~\\ref{Omega}). The 6th column gives the loop-induced mass splitting between neutral and charged DM components (section~\\ref{splitting}); for scalar candidates a coupling with the Higgs can give a small extra contribution, that we neglect. The 7th column gives the $3\\sigma$ range for the number of events expected at LHC (section~\\ref{coll}). The last column gives the spin-independent cross section, assuming a sample vale $f=1/3$ for the uncertain nuclear matrix elements (section~\\ref{DMexp}). } \\end{table} ",
        "conclusions": "We extended the Standard Model by adding a spin-0 or spin-1/2 $n$-tuplet of $\\SU(2)_L$ with hypercharge $Y$ that only has gauge interactions and mass $M$. Some multiplets contain neutral components, that are potential Dark Matter (DM)  candidates. \\begin{itemize} \\item Multiplets with $Y \\neq 0$ are already excluded by direct DM searches. They can be resurrected by introducing non-minimal mechanisms that prevent  $Z$-mediated DM/nuclei coupling, e.g.\\ by appropriately mixing their neutral components with a  singlet. \\item Multiplets with $Y=0$ and odd $n=\\{3,5,7,\\ldots\\}$  contain allowed DM candidates. \\begin{itemize} \\item For $n=3$ one needs to impose DM stability by hand. \\item For $n\\ge 5$ the stability is instead automatically guaranteed by renormalizability, much alike proton stability. \\end{itemize}\\end{itemize} The set of interesting candidates is bounded by $n\\circa{<}7$ in order to avoid Landau poles in $\\alpha_2$.  Gauge interactions are spontaneously broken and thereby induce a non-trivial and peculiar Minimal DM (MDM) phenomenology, fully computable in terms of a single unknown parameter: the DM mass $M$. Electroweak breaking effects induce a mass splitting $\\Delta M \\sim \\alpha M_W$ among the components of any given multiplet, making the neutral component lighter than the charged components. Assuming that only one MDM multiplet is present, its mass $M$ is determined by the request that  its relic thermal abundance equals the observed DM abundance. Co-annihilations  play a crucial r\\^ole, giving $M\\sim $ few TeV. Since $M\\gg M_W$, various MDM properties depend dominantly only on the MDM gauge charge, while the microscopic MDM properties (such as their spin) become irrelevant.    \\medskip  The simplest fully successful MDM candidate  is a fermionic $\\SU(2)_L$ quintuplet with mass $M\\approx 4.4\\TeV$. MDM candidates are listed in table~\\ref{tab:1}: some are fully successful (automatically stable and consistent with DM searches), others require a stabilization mechanism (e.g.\\ the wino-like candidate) or a way to elude the bounds from direct DM searches (we list only those with $n\\le 4$), or both (e.g.\\ the Higgsino-like candidate). If multiple MDM multiplets exist, all their masses become lighter than in table~\\ref{tab:1}: e.g.\\ $42\\%$ lighter in presence of 3 identical families of a single multiplet, significantly increasing the number of events expected at LHC.  MDM multiplets contain charged components, slightly heavier than the neutral DM component. For $Y=0$ the charged $\\DM^\\pm$ is $\\Delta M = 166\\MeV$ heavier and has a life-time $\\tau =44\\,{\\rm cm}/(n^2-1)$, giving a clean displaced-vertex signature at colliders. The branching ratios are predicted to be $\\hbox{BR}(\\DM^\\pm \\to \\pi^\\pm\\DM^0)= 97.7\\%$, $\\hbox{BR}(\\DM^\\pm\\to\\DM^0 e^\\pm\\nubarnu_e) = 2.05\\%$, $\\hbox{BR}(\\DM^\\pm\\to\\DM^0\\mu^\\pm\\nubarnu_\\mu) = 0.25\\%$. We computed the event rate at LHC, finding that LHC cannot probe all MDM multiplets, being more sensitive to the ones with lower $n$ (and, if multiple multiplets are present, to the ones that would give subdominant contributions to the DM density).  On the contrary direct DM searches are more sensitive to higher $n$ MDM multiplets (and to ones that dominate the DM density). Indeed one-loop diagrams generate a spin-independent MDM/nucleus cross section parameter $\\sigma_{\\rm SI}\\sim 10^{-44} (n/5)^4 \\,{\\rm cm}^2 $ (up to a QCD uncertainty of about one order of magnitude). As illustrated in fig.\\fig{direct}, this is within the sensitivity of future experiments.  DM DM annihilations in the galactic halo into vector bosons can be resonantly enhanced, giving indirect DM signals, only for DM masses close to the values listed in eq.\\eq{resonant}. CC production of $\\DM^\\pm$ occurs at tree level with a much larger cross section, that can exceed $10^{-34}\\,{\\rm cm}^2$, but in ordinary situations it is forbidden kinematically. We also discussed prospects for attempting CC direct DM detection by accelerating an intense nuclear beam.   \\appendix  \\footnotesize  \\paragraph"
    },
    "0512/astro-ph0512046_arXiv.txt": {
        "abstract": "The rotation curve for the IV galactic quadrant, within the solar circle, is derived from the Columbia University - U. de Chile CO(J=1$\\to$0) survey of molecular gas. A new sampling, four times denser in longitude than in our previous analysis, is used to compute kinematical parameters that require derivatives w/r to galactocentric radius;  the angular velocity $\\Omega(R)$, the epicyclic frequency $\\kappa(R)$,  and the parameters $A(R)$ and $B(R)$ describing, respectively, gas shear and vorticity. The face-on surface density of molecular gas is computed from the CO data in galactocentric radial bins for the subcentral vicinity, the same spectral region used to derive the rotation curve, where the two-fold ambiguity in kinematical distances is minimum.  The rate of massive star formation per unit area is derived, for the same radial bins, from the luminosity of IRAS point-like sources with FIR colors of UC H{\\small II} regions detected in the CS(J=2$\\to$1) line. Massive star formation occurs preferentially in three regions of high molecular gas density, coincident with lines of sight tangent to spiral arms. The molecular gas motion in these arms resembles that of a solid body, characterized by constant angular velocity and by low shear and vorticity. The formation of massive stars in the arms follows the Schmidt law, $\\Sigma_{MSFR} \\propto [\\Sigma_{gas}]^n$, with an index of $n = 1.2 \\pm 0.2$.   Our results suggest that the large scale kinematics, through shear, regulate global star formation in the Galactic disk. ",
        "introduction": "The rotation curve, describing the circular speed of rotating material as a function of galactocentric radius, is a fundamental tool for the study of the kinematics of our Galaxy. It is best  derived, because of interstellar extinction, from observations of atomic and molecular gas in radio and mm wavelengths. The derivation involves determining the {\\it terminal velocity}, or maximum absolute radial velocity relative to the Sun, toward lines of sight that sample the Galaxy within the solar circle (quadrants I and IV). Such terminal velocities correspond, assuming pure circular motion, to the tangent points to circumferences around the galactic center, named  {\\it subcentral points}. These points subtend a circumference that connects the solar position with the galactic center. A detailed analysis of the rotation curve can reveal important physical characteristics of the rotating material, such as  the amount of shear and vorticity at each galactocentric radius. These physical quantities regulate the gravitational stability of a differentially rotating gaseous disk and, consequently, the large scale distribution and properties of star formation in the galactic disk.  The first derivation of the rotation curve for the IV galactic quadrant that made use of the CO(J=1$\\to$0) line - the best tracer of molecular hydrogen in the interstellar medium - was presented by Alvarez, May, \\& Bronfman (1990). The spectral data used to determine the terminal velocities were taken from the Columbia - U. de Chile surveys \\citep{ grabelsky87, bronf89}, which have a sampling interval of 0$^{\\circ}$.125 (roughly the beam size). However, the terminal velocities in \\citet{alvarez90} were measured only every 0$^{\\circ}$.5 in galactic longitude, due to difficulties involved in the visual examination of a very large number of spectra. A new derivation of the rotation curve, that uses a computer search code to examine all the available spectra ( $\\approx$15000), is presented here. The disk kinematic characteristics in the IV galactic quadrant are analyzed in detail, from this new rotation curve. These characteristics, as a function of galactocentric radius, are compared with the molecular gas density and with the local rate of massive star formation.  A proper derivation of the spiral pattern of our Galaxy requires knowledge of the distances to the adopted tracers. These distances are also required to compute the masses and luminosities of such tracers. For the gas, kinematical distances can be obtained from radial velocity data of radio line observations, adopting a rotation curve, under the assumption of pure circular motions. For clouds within the solar circle, however, there is a two-fold ambiguity in the kinematic distance, that is difficult to circumvent and has to be resolved in a case-by-case basis. But in the vicinity of the subcentral points such ambiguity is minimal, since at the subcentral points themselves the kinematic distances are univocally defined.  It is worth noting that large scale streaming motions in spiral arms, with amplitude of $\\sim$10 km/s, which produce deviations from pure rotation, have been observed in a number of regions of the Galaxy (Burton et al. 1988).  Streaming motions of such amplitude may introduce uncertainties of up to 5\\% in the estimation of galactocentric radii when the streaming is along the line of sight. In such unfavorable case, the corresponding  uncertainties in the estimated distances, for the section of the Galaxy analyzed here, may go from of 0.6 kpc to 1.7 kpc. In any case, for objects beyond $\\sim$\\,3\\,kpc from the Sun, because of optical extinction, kinematical distances are usually the only ones available.  Massive stars are formed within aggregates of molecular gas and dust of 10$^5$-10$^6$ solar masses, about 50-100 pc in size, which are commonly known as giant molecular clouds, or GMCs for short. The association between OB stars and the interstellar medium has been established through optical, infrared, and  CO observations of GMCs close enough to be largely unaffected by extinction (Orion, Carina, etc).  The physical conditions in GMCs control their rates of OB star formation, and are one of the main agents that regulate the evolution of the galactic disk \\citep{evans99}.  There is a close relationship between the galactic spiral structure and the formation of GMCs and, hence, with the formation rate of OB stars \\citep{dame86, solomon86}.  Therefore, the GMCs and the regions of OB star formation provide a very good tool to trace the spiral arm pattern of a galaxy.  An early description of the Milky Way spiral arm pattern was given by \\citet{gg76}, who observed the H109$\\alpha$ line emitted in H{\\small II} regions associated with young massive stars.  A four arm spiral pattern for the southern Milky Way was later proposed by \\citet{cyh87}, using a larger observational database of hydrogen recombination lines (H109$\\alpha$ \\& H110$\\alpha$). The four arm spiral pattern is in general agreement with that obtained from  H{\\small I} and CO large scale observations of the Galaxy \\citep{rob83, grabelsky87, bronf88, alvarez90, valle02}.  Star formation is likely to occur  in regions where the gas in the Galactic disk is unstable to the growth of gravitational perturbations.  In a classical paper, \\citet{schm59} introduced the parametrization of the volume density of star formation and the volume density of gas, relating them through a power law; such parametrization, known as \"Schmidt Law'', has been studied observationally \\citep{kennic89, wong02} and explained on theoretical grounds \\citep{toomre64, tan00}. A study of the gas stability in the galactic disk must include (a) comparison of the gas density with a critical value above which the gaseous aggregates undergo gravitational collapse \\citep{toomre64, kennic89} and (b) examination of the gas shear rate, that governs the process of destruction of  molecular clouds (e.g. Kenney, Carlstrom, \\& Young 1993; Wong \\& Blitz 2002), presumably through the injection of turbulent motions  \\citep{maclow04}.  The link between massive star formation and kinematical conditions in disks has been studied mostly for external spiral galaxies \\citep{aalto99, wong02, bossier03}, where the spatial resolution that can be achieved by the observations is not as good as for the Milky Way. The main goal of the present paper is, therefore, to accurately describe the spiral arm structure in the {\\it subcentral vicinity} of our Galaxy, focusing on the molecular gas kinematics, density, and on the rate of massive star formation, with the hope of contributing to the understanding of the formation and evolution of disk galaxies in general. The analysis is carried out  for the IV galactic quadrant, where the spiral structure is more evident \\citep{bronf88} than in the I quadrant. Preliminary work has been presented by \\citet{cys83} and, more recently by \\citet{aluna01}.  Section (\\S 2) describes the observational datasets used, the most complete available in their kind. These data are used in section (\\S 3) to derive the rotation curve and analyze the relation between molecular gas kinematics, molecular gas surface density, and massive star formation rate.  The validity of Schmidt Law for the Milky Way is analyzed in section (\\S 4), and a summary of the results is given in section \\S 5. ",
        "conclusions": "This work analyzes the correlation between molecular gas kinematical properties, molecular gas surface density,  and rate of massive star formation in the IV galactic quadrant, using the most complete data bases available.  The analysis  is carried out for galactocentric radial bins 0.5 kpc wide, a compromise to avoid arm-interarm confusion while having good statistics. The data used are restricted to the subcentral vicinity, to avoid the two-fold distance ambiguity within the solar circle. The main conclusions are:  $\\bullet$ The rotation curve obtained is similar to that presented by Alvarez et al. (1990).  Since the sampling in longitude is 4 times denser,  however, it can be used to calculate, as a function of galactocentric radius, kinematical parameters that require radial derivatives.  $\\bullet$ The angular velocity, $\\Omega(R)$, the epicyclic frequency, $\\kappa(R)$, and the parameter $A(R)$ describing the gas shear, tend to decrease with galactocentric radius; the parameter $B(R)$, describing the gas vorticity, tends to grow. The values derived for Oort's constants $A_0$ and $B_0$, at R = $R_0$, are consistent with those recommended by the IAU.  $\\bullet$ Shear and vorticity have relative  minima, and the epicyclic frequency relative maxima, at radii 0.39, 0.47 and 0.73 $R/R_0$, coincident with the known positions of spiral arm tangent regions.  Near these radii the kinematics are characteristic of solid body rotation: $A(R)$, proportional to the radial derivative of the angular velocity, tends to zero, so the angular velocity is roughly constant. The relative maxima in epicyclic frequency are consistent with such scenario since $\\kappa$ is $\\sqrt{2}$ times higher for solid body rotation than for a flat rotation curve.  $\\bullet$ Differential rotation and shear are weaker for the spiral arm regions than for the interarm regions. The relative importance of tidal shear w/r to gravitation in the stability of the gas for spiral arm regions, where the rotation curve resembles that of a solid body, is half than for spiral arms, where the rotation curve is nearly flat.  $\\bullet$ Massive star formation occurs in regions of high molecular gas density, roughly coincident with the three lines of sight tangent to spiral arms.  In these arms the formation of massive stars follows the Schmidt law, $\\Sigma_{MSFR} \\propto [\\Sigma_{gas}]^n$, with an index of $n = 1.2 \\pm 0.2$. While this law is characteristic of spiral galactic disks, here it applies to much smaller spatial scale.  A modified version of Schmidt law, which modulates the gas density by the angular velocity, $\\Sigma_{MSFR} \\propto\\Sigma_{gas}\\, \\Omega$, describes better the behavior of the gas at this scale, suggesting that the kinematics, through shear, regulate global star formation in the Galactic disk."
    },
    "0512/astro-ph0512336_arXiv.txt": {
        "abstract": "We present a model which predicts inflation without the presence of inflaton fields, based on the $\\epsilon R^2$  and Starobinsky models. It links the above models to the reheating epoch with conformally coupled massive particles created at the end of inflation. In the original Starobinsky model, the reheating era was created by massless non-conformally coupled particles. We assume here that non-conformal coupling to gravitation does not exist. In the $\\epsilon R^2$ model, inflation is produced by the gravitational Lagrangian to which a term $\\epsilon R^2$ is added, where $\\epsilon$ is a constant and $R$ is the Ricci scalar. Inflation is created by vacuum fluctuations in the Starobinsky model. Both models have the same late-inflation time-dependence, which is described by a characteristic mass $M$. There is a free parameter $H_0$ on the order of the Planck mass $M_{Pl}$ that determines the Hubble parameter near the Planck epoch and which depends upon the number and type of particles creating the vacuum fluctuations in the Starobinsky model. In our model, we assume the existence of particles with a mass $m$, on the order of $M$, conformally coupled to gravity, that have a long decay time. Taking  $m\\equiv FM$, we investigate values of $F=0.5$ and $0.3$. These particles, produced $\\sim 60\\,e$-folds before the end of inflation, created the nearly scale invariant scalar density fluctuations which are observed. Gravitational waves (tensor fluctuations) were also produced at this epoch. At $t_{\\rm end}$, the Hubble parameter begins to oscillate rapidly, gravitationally producing the bulk of the $m$ particles, which we identify as the origin of the matter content of the Universe today. The time required for the Universe to dissipate its vacuum energy into $m$ particles is found to be $t_{\\rm dis} \\simeq 6\\,M_{Pl}^2/M^3F$. We assume that the reheating time $t_{RH}$ needed for the $m$ particles to decay into relativistic particles, is very much greater than that necessary to create the $m$ particles, $t_{\\rm dis}$. A particle physics theory of $m$ can, in principle, predict their decay rate $\\Gamma_{mr}\\equiv t_{RH}^{-1}$. From the ratio $f\\equiv t_{\\rm dis}/t_{RH}$, $F$ and $g_{\\ast}$ (the total number of degrees of freedom of the relativistic particles) we can, then, evaluate the maximum temperature of the Universe $T_{\\rm max}$ and the reheat temperature $T_{RH}$ at $t_{RH}$. From the observed scalar fluctuations at large scales, $\\delta\\rho/\\rho\\sim 10^{-5}$, we have the prediction $M\\cong 1.15\\times 10^{-6} M_{\\rm Pl}$ and the ratio of the tensor to scalar fluctuations, $r\\cong 6.8\\times 10^{-4}$. Thus our model predicts $M$, $t_{\\rm dis}$, $t_{\\rm end}$, $T_{\\rm max}$, $T_{RH}$, $t_{\\rm max}$, and $t_{RH}$ as a function of $f$, $F$, and $g_{\\ast}$ (and to a weaker extent the particle content of the vacuum near the Planck epoch). A measured value of $r$ that is appreciably different from $r=6.8\\times 10^{-4}$ would discard our model (as well as the Starobinsky and $\\epsilon R^2$ models). ",
        "introduction": "The standard model of inflation, based on the existence of a scalar inflaton field, makes the following assumptions:  \\begin{enumerate} \\item The beginning of inflation occurs at an energy $<< M_{Pl}$ (Planck energy). Its origin is unknown and the state of the Universe before the beginning of inflation is undefined; \\item A large initial displacement of $\\phi$ from the minimum of $V(\\phi)$ is necessary for the onset of inflation (such as in the chaotic inflation model); \\item The potential energy of the inflaton dominates its kinetic energy during inflation; and \\item The inflaton potential, $V(\\phi)$, and its first derivative are defined by observations $\\sim 60\\,e$-folds before the end of inflation. \\end{enumerate} In complex inflation theories, there can be more than one inflaton. (See \\cite{basset05} for a recent review of inflation theory with inflatons.)  Here we present a model which links the Starobinsky and $\\epsilon R^2$ models of inflation, where $R$ is the Ricci scalar to the reheating era. All three models, the Starobinsky, $\\epsilon R^2$ and ours, do not involve inflatons to create inflation. They also avoid most of the above assumptions.  In the Starobinsky model \\cite{star1}, an $R^2$ term in the effective Lagrangian dominates inflation at late times (see also \\cite{suen}, \\cite{mijic}). There is no sharp boundary between the Starobinsky model and $\\epsilon R^2$ model since the latter is the particular case of the former in the limit $M \\ll H_0$ (using the notation of Eq.(9), with some small non-local terms (due to non-zero rest masses of conformally coupled quantum fields) omitted. However, the Starobinsky and $\\epsilon R^2$ models have the same qualitative behavior at $\\sim 60\\,e$-folds before the end of inflation, when the presently observed scalar and tensor fluctuations were produced.  Both the Starobinsky and $\\epsilon R^2$ cosmologies are characterized by a single mass $M$ ($\\equiv M_{Pl}/\\sqrt{24\\,\\epsilon}$ in the $\\epsilon R^2$ cosmology and $M_{Pl}/\\sqrt{48\\pi k_1}$ in the Starobinsky model, where $M_{Pl}\\equiv G^{-1/2}$ is the Planck mass and $k_1$ is the coefficient of the term which contains the second derivative of $R$ in the quantum corrected vacuum expectation value of the energy-momentum tensor [Eq.(\\ref{anomaly})]). The mass $M$ characterizes the end of the inflation period, during which, the Hubble parameter varies slowly. A period then begins, in which $H$ oscillates rapidly as $H\\propto (1/t)\\cos^2{\\omega t}$ and the cosmological scale factor varies as $a(t)\\propto t^{2/3}[1+\\sin{[2\\omega t]}/(3\\omega t)]$, where $\\omega \\simeq M/2$. When averaged over several oscillations, the Universe expands as a classical matter-dominated Universe.  Although the $\\epsilon R^2$ model can be considered to be the simplest way to produce inflation, i.e., by means of a simple modification of the gravitational Lagrangian, we concentrate here on the Starobinsky model since it is more complete. It links the beginning of the inflation period to the beginning of the Universe and also describes the end of inflation in detail.  The Starobinsky model suggests that for energy densities and curvatures near the Planck scale, quantum corrections to Einstein's equations become important (as discussed in detail by Vilenkin \\cite{vile85a}). In the Starobinsky model, inflation is driven by one-loop corrections due to quantized matter fields \\cite{star1} (see also \\cite{star2,fhh,Mukhanov81,ander,anapel1,anapel2}). The model is consistent with a Universe that was spontaneously created, as discussed by Grishchuk and Zel'dovich \\cite{grish1}.  The beginning of the Starobinsky inflation period can be associated with the beginning of the Universe due to quantum fluctuations of the vacuum. Tryon \\cite{tryon} was the first to suggest that a closed Universe can be created spontaneously as a result of a quantum fluctuation. Vilenkin \\cite{vile85a,vile8234,vile85b}, Zel'dovich and Starobinsky \\cite{zel1}, and Linde \\cite{linde} were the first to attempt to describe the quantum creation of a Universe in the framework of quantum gravity. The picture that emerges is one of a Universe tunneling quantum mechanically to a de Sitter space time. At the moment of nucleation ($t=0$), the Universe has a size $a(0)=H_{\\rm in}^{-1}$. This is the beginning of time and, from that point on, the Universe evolves along the lines of the inflation scenario.  In the Starobinsky model, inflation is produced by the vacuum energy $\\rho_V$, which has negative pressure, $P = -\\rho_V$. Inflation in both the Starobinsky and our models can be described by an effective geometric scalar particle $M$. In our model, there is an additional massive particle $m$ produced at the end of inflation, which is freely moving and which produces positive pressure.   Structure in the Universe primarily comes from almost scale-invariant superhorizon curvature perturbations \\cite{giudice01,giudice04}. In our model, a mass $m$ is much less than the Hubble parameter during inflation. The mechanism of $m$ particle production from inflation is based on the observation that particles that are massive in the present-day vacuum, could have been very light during inflation. This implies that fluctuations of a generic scalar field $\\chi$ with mass $m \\ll H$ during inflation are copiously generated, with an almost scale invariant spectrum \\cite{linde90,liddle,riotto}. The particles become heavy and non-relativistic at the end of inflation.   The end of the Starobinsky inflation period has been suggested to be due to the masses of the particles in the vacuum fluctuations [14]. Thus, since the mass $M$ describes the end of inflation in both the Starobinsky and our models, $M$ is a natural mass scale for the particles that are created at the end of inflation. In our model, we then have the scenario that particles of mass comparable to the mass $M$ in the vacuum fluctuations first create the inflation, after which, particles of mass $m$ comparable to, but slightly less than $M$, are produced from the vacuum due to the rapid change of the Hubble parameter. The particles $m$ are conformally coupled to gravitation (Ricci scalar). These massive particles create the reheating epoch of the Universe.  Our model can be compared with that of the Starobinsky model, in which massless particles, non-conformally coupled to gravitation, directly create the reheating era. Here we assume that non-conformal coupling does not exist.   Previously, gravitational production of massive particles has been investigated in order to explain the observed ultra-high energy cosmic rays, produced as a result of heavy particle decay (masses $>10^{12}\\,{\\rm GeV})$ \\cite{cosmray}. The gravitational particle production of the heavy particles  $m$ can be described assuming a given background metric \\cite{kuzmin,kolb}.  The paper is organized as follows. In section II, we give the main results of the Starobinsky model, as discussed by Vilenkin \\cite{vile85a}. We discuss the gravitational production of the $m$ particles in section III. In section IV, we derive the scalar density fluctuations produced $\\sim 60\\,e$-folds before the end of inflation. We obtain the ratio $r$ of the tensor to scalar fluctuations in section V. The reheating of the Universe is discussed in section VI. Our conclusions and discussion are presented in section VII. ",
        "conclusions": ""
    },
    "0512/astro-ph0512100_arXiv.txt": {
        "abstract": "The software package aXe provides comprehensive spectral extraction facilities for all the slitless modes of the ACS, covering the Wide Field Channel (WFC) grism, the High Resolution Channel (HRC) grism and prism and the Solar Blind Channel (SBC) prisms. The latest developments to the package apply to all ACS slitless modes leading to improved spectral extraction. Many thousands of spectra may be present on a single deep ACS WFC G800L image such that overlap of spectra is a significant nuisance. Two methods of estimating the contamination of any given spectrum by its near neighbours have been developed: one is based on the catalogue of objects on the direct image; another uses the flux information on multi-filter direct images. An improvement to the extracted spectra can also result from weighted extraction and the Horne optimal extraction algorithm has been implemented in aXe. A demonstrated improvement in signal-to-noise can be achieved. These new features are available in aXe-1.5 with the STSDAS 3.4 release. ",
        "introduction": "As part of a collaborative project between STScI and the Advanced Camera for Surveys (ACS) IDT, the ST-ECF provides comprehensive support for the slitless spectroscopy modes of the ACS. As well as support to users of the ACS grism (G800L for the Wide Field Channel, WFC and High Resolution Channel, HRC) and prisms (PR200L for the HRC and PR110L and PR130L for the Solar Blind Channel, SBC) and contributions to the ground and in-orbit calibrations of the slitless modes, a primary pillar of this project has been the provision of an extraction package, called aXe. The package consists of a number of self-contained modules, which perform the basic steps - defining apertures for extraction of a spectrum, assigning wavelengths to pixels, flat fielding, extracting rectified 2D spectra and a 1D spectrum, and applying flux calibration. The modules are scripted in Python to allow easy integration into Pyraf/STSDAS (see K\\\"{u}mmel et al. 2005 for more details). The aXe user manual ({\\tt http://www.stecf.org/software/aXe/documentation.html}) provides full details for installing and running the package.  The fundamental aspect of slitless spectroscopy is that the individual objects define their own `slit' in terms of position on the detector and object height of the dispersed spectrum; the object width in the dispersion direction affects the spectral resolution. aXe uses a catalogue of the observed targets, which is usually taken from a matched direct image as the starting point of the reduction process. In the design of aXe it was decided to make the software as general as possible, so that in the longer term not only slitless spectra from ACS could be extracted. This flexibility is engineered by putting all instrument specific parameters in a configuration file. Thus the specification of the spectral traces, the dispersion solutions, the name of the flat field file, the sensitivity file name, etc are all listed in a single file for each instrument mode. Thus for ACS, there are six configuration files; three for the G800L with the WFC (one for each chip) and HRC, one for PR200L, and one each for the SBC with PR110L and PR130L. The built-in flexibility has paid off since aXe has also been used to extract spectra from multi-object spectra taken with the VLT FORS2 instrument (K\\\"{u}mmel et al. 2006).  Since the first release of aXe in 2002 ready for installation of ACS into HST, the package has evolved. In particular in 2004 a major enhancement was added with the use of `drizzle' (Fruchter \\& Hook, 2002) to combine 2D spectra of individual objects when the data is taken with dithers of the telescope. This turns out to be a very common observational procedure, at least for the grism modes, in order to recover the undersampling of the Point Spread Function (PSF) and to mitigate the effect of pixel sensitivity variations and hot pixels. Since the 2004 release (aXe-1.4), two enhancements have been added which are here described. The sensitivity of the ACS slitless spectroscopy modes implies that, despite the small pixel scale and compact PSF, even for high Galactic latitude fields, the surface density of detected spectra on moderate exposures ($>$ thousands of seconds) displays crowding and overlap. A high priority was to indicate to the user which pixels in an extracted spectrum are affected by overlap with spectra of other objects. This has now been implemented in a quantitative way, whereby the estimated value of the contaminating flux contributing to a spectrum pixel is output. The second enhancement was to apply the well known technique of weighting by the spatial profile when forming a 1-D spectrum from the 2D spectrum on the detector (optimal extraction of Horne 1986). Both these enhancements are described. ",
        "conclusions": ""
    },
    "0512/astro-ph0512192_arXiv.txt": {
        "abstract": "The acceleration of charged particles in the presence of a magnetic field and gravitational waves is under consideration. It is shown that the weak gravitational waves can cause the acceleration of low energy particles under appropriate conditions. Such conditions may be satisfied close to the source of the gravitational waves if the magnetized plasma is in a turbulent state. ",
        "introduction": "The interaction of a charged particle with a gravitational wave (GW) has been studied  by using a Hamiltonian formalism (\\cite{VP}-\\cite{KVP3}). This approach is similar to the one used for the study of the interaction of ions with electrostatic waves in magnetized plasmas (\\cite{Fuku}-\\cite{SmK}). The existence of chaos in phase space has been shown as a consequence of the overlap of resonances. Recently, the acceleration of the charged particles due to the chaotic diffusion has been associated with bursts \\cite{VVP}. However, such a diffusion is obtained for relatively high energy particles and becomes efficiently present only for large amplitude values of the GW ($a>0.1$).  Furthermore, the existence of a parametric resonance in the general problem of interaction between particles and waves  is discussed in the classical book of Landau and Lifshitz  \\cite{LaLi}. For the particular problem Kleidis {\\it et al.} \\cite{KVP3} showed that a parametric resonance arises for a set of frequencies of the GW  of the form $\\nu=2/n,\\: n=1,2,...$ and an estimation of its width is given. Also, the parametric resonance has been associated with the existence of chaotic motion.  In the present paper we study the possible fast acceleration of low energy particles for small amplitudes of the GW by taking into account two special properties of the system. The first is the parametric resonance which, as we shall show, arises only for $\\nu=2$. The second property is the integrability of the system, and subsequently the absence of chaos, when the GW propagates in a parallel direction with that of the magnetic field.  In section 2 we review the Hamiltonian formalism of the model which describes the interaction. In section 3, we study systematically the rise of a parametric resonance and in section 4 we prove the integrability of the system in the case of parallel propagation. It is shown in section 5 that when the above integrable case is in parametric resonance fast acceleration of particles occurs. The astrophysical implications are discussed in section 6. ",
        "conclusions": "We have shown that gravitational waves with frequency $\\omega$ and amplitude $\\alpha$ can accelerate charged particles to very high energies when two conditions are met (a) the gravitational wave propagates along the external magnetic field ($\\theta\\approx 0^\\circ$) and (b) the ratio of the gravitational wave frequency to gyro frequency is about 2 ($\\nu\\approx 2)$. The model used has  a number of important simplifications: (1) We assume that the magnetic field is constant (2) The energy absorption by the particles is small compared to the energy curried by the gravitational wave and (3) we ignore the collective plasma effects associated with the particle acceleration.  Retaining the above simplifications we may proceed to estimate the energy gained by the particles in a realistic astrophysical system. We assume that a volume of length $L \\approx 10^{12} cm$ is close to a source of gravitational waves. The volume is filled with plasma with density $n\\approx 10^{12} \\textrm{particles}/cm^{-3}$ and the magnetic field is in a fully turbulent state. The gravitational wave is crossing  the volume in few seconds and every time it travels along the ambient field ($\\theta \\approx 0^\\circ$) with almost zero strength its energy increases. The zero strength magnetic field, which is associated to null surfaces, is an important requirement since the gravitational wave has low frequency (of few KHz) and acceleration is possible only for $\\nu\\approx 2$).  The energy increase is fast and can be substantial ($\\Delta\\gamma\\approx 10$) if the amplitude of the gravitational wave is $a\\approx 10^{-4}-10^{-5}$. Therefore the particles diffuse in energy space not by small random steps, as it was the case in the stochastic interaction studied earlier (\\cite{VP}- \\cite{KVP3}), but through large random energy jumps.  It is beyond the scope of this article to present a detailed analysis of the above process but we can present a rough estimate of the energy transferred from the gravitational wave to the plasma. Let us assume that a particle travel a distance $L_p\\geq L$ (its trajectory is not a straight line) before exiting the interaction volume. Let  us also assume that only in small parts of its trajectory ($L_{int}\\approx $ a few kilometers) the particle is able to meet the conditions needed for its acceleration. For a small portion of protons, e.g $n_{acc}=n_{total}/100$), which participate in the interaction, these regions can be a fraction of $N_{max}=L_p/L_{int}\\approx 10^7$ and the energy gain will be $E_{kin} \\approx 10^6 (\\Delta \\gamma-1) m_p c^2 \\approx 10^4 ergs$ (or $10^{16} eV$), where $m_p$ is the mass of proton. The number of particles participating in this interaction can be as many as $N_{total} \\approx n_{acc} L^3 \\approx 10^{46}.$ The total energy lost by the gravitational wave is, approximately, $10^{48}-10^{49}$ ergs, which constitute a small fraction of its total energy $10^{58}$ ergs (see \\cite{Putten}).   According to the model, the acceleration conditions ($\\theta\\approx 0^\\circ$, $\\nu\\approx 2$) can be met sporadically along the trajectory of a charged particle in the vicinity of a gravitational wave source when the magnetized plasma is in a turbulent state. A few particles can then reach very high energies in a few seconds absorbing only a very small fraction of the gravitational wave. It remains to be seen, by performing further simulations, if the process described above for particle acceleration can be significant in a realistic astrophysical environment. In this case the electromagnetic emission from the accelerated particles can be a precursor of the gravitational wave for easier detection."
    },
    "0512/astro-ph0512471_arXiv.txt": {
        "abstract": "{} {We address the topic of the Intra-Night Optical Variability of the BL Lac object \\sbz.} {To this purpose a long term observational campaign was performed, from 1996 to 2003, which allowed the collection of a very large data set, containing 10,675 photometric measurements obtained in 102 nights.} {The source brightness varied in a range of about 2~mag, although the majority of  observations were performed when it was in the range $13.0 < R < 13.75$. Variability time scales were estimated from the rates of magnitude variation, which were found to have a distribution function well fitted by an exponential law with a mean value of 0.027~mag/h, corresponding to an e-folding time scale of the flux $\\tau_F =$ 37.6~h. The highest rates of magnitude variation were around 0.10--0.12 mag/h and lasted less than 2~h. These rates were observed only when the source had an $R$ magnitude $<$ 13.4, but this finding cannot be considered significant because of the low statistical occurrence. The distribution of $\\tau_F$ has a well defined modal value at 19~h. Assuming the recent estimate of the beaming factor $\\delta \\sim$ 20, we derived a typical size of the emitting region of about 5$\\times$10$^{16}/(1 + z)$ cm. The possibility to search for a possible correlation between the mean magnitude variation rate and the long term changes of the velocity of superluminal components in the jet is discussed.} {} ",
        "introduction": "\\label{sec:Intro}  The radio source \\sbz~ was identified with a bright and highly variable BL Lac object, characterized by a strong  featureless optical continuum (\\citealt{Bie81}). The failure in detecting a host galaxy both in HST direct imaging (\\citealt{Urr00}) and in high S/N spectra (\\citealt{RecSto01}) suggests that its redshift $z$ should be greater than 0.52 (\\citealt{Sba05}; see also \\citealt{Sch92}  and \\citealt{Wag96}). Variations on short time scales (a fraction of hour) have been detected in several occasions at frequencies ranging from radio to X-rays (\\citealt{Qui91}; \\citealt{Wag96}; \\citealt{Gab00}; \\citealt{HeiWag96}; \\citealt{Nes02}; \\citealt{Gio99}; \\citealt{Vil00}; \\citealt{Wu05}). The recent history of the flux of \\sbz~ in the $R$ (Cousins) band is plotted in Fig.~\\ref{fig:lc0716}, which spans the time interval from 1997 to 2003. The photometric points up to 2001 have been extracted from the data base by \\citet{Rai03}, taking only one measurement per day, whereas points from 2002 to 2003 are new data obtained by our group (\\citealt{Nes05}). The general structure of this curve shows that the mean flux of \\sbz~ varies in the range 5-20~mJy with some flares in which it reaches and overwhelms 30~mJy. These flares are separated by time intervals variable from $\\sim$~1 to $\\sim$~3 years. In Fig.~\\ref{fig:lc0716}  there are four large flares having a typical duration of the order of 1--2~months from which it is possible to estimate a flaring duty cycle of about 5--10\\%.  \\begin{figure*} \\vspace{1.0cm} \\hspace{2.0cm} \\epsfysize=9cm \\epsfbox{fig1.ps} \\caption[]{ The light curve in the $R$(Cousins) band of \\sbz~ from February 1997 to March 2003. Vertical bars on the top mark the epochs of our INOV observations. } \\label{fig:lc0716} \\end{figure*}  In this paper we focus our attention on the so called Intra-Night Optical Variability (INOV) or {\\it microvariability}. Brightness changes of BL Lac objects, having amplitudes of about 10--20\\% and occurring on time scales as short as a fraction of an hour, were studied since the eighties by Miller and coworkers (\\citealt{Mil89}; \\citealt{Car91}). This phenomenon was after detected in many sources of this class and it could be considered one of their characterizing properties. INOV in BL Lac objects and other blazars has been investigated by several authors, using either single or multi band photometry: among BL Lac objects showing such activity we recall AO 0235+164 (\\citealt{Rom00}), S4 0954+658 (\\citealt{Pap04}), BL Lacertae (\\citealt{Mas98}; \\citealt{Nes98}; \\citealt{Pap03}) whereas studies on samples of sources are those on LBL (Low energy peaked BL Lacs) objects (\\citealt{HeiWag96}), EGRET blazars (\\citealt{Rom02}) and other BL Lac objects and radio-core dominated blazars (\\citealt{Sag04}; \\citealt{Sta05}). These works are generally based on data sets, for single sources, obtained in a not high number of nights and/or observations. Moreover, the analysis is mainly focused on the search for recurrence time scales detectable in the individual light curves.  A first detailed study of INOV in \\sbz~ is that of \\citet{Wag96}, who investigated rapid variations in multifrequency (from radio to X-rays) campaigns and observed a quasi-periodic behaviour with a typical recurrence time of about two days and a high correlation between the optical and radio flux changes. The possibility of an harmonic component in the optical flux of \\sbz~ was after confirmed by \\citet{HeiWag96}, who did not detect a similar effect on the other BL Lacs of their sample. \\citet{Sag99} reported INOV multiband observations covering 4 weeks in 1994 and found only three major events of rapid variability in which the highest magnitude variation rate was around 0.03~mag/h. \\citet{Vil00} reported the results of a WEBT (Whole Earth Blazar Telescope) campaign on \\sbz~ from 16 to 22 February 1999 in which a relevant INOV was observed every night with magnitude variation rates up to $\\sim$~0.1~mag/h.  A well established definition of INOV properties for a given source (or for a class of sources) can be obtained only from the analysis of a relatively large number of observations, possibly performed in different brightness states.  The definition of time scale for non-periodic phenomena is not univocal and can depend on the type of variation: in this work we consider a time scale based on the magnitude variation rates.  We worked extensively on \\sbz, bright enough to obtain good photometric data with small aperture telescopes and short exposure times. Our INOV observational campaign of \\sbz~ started in 1996 and since November 1998 we undertook a more intense data acquisition concluded in spring 2003. In this paper we report a large set of observational data, containing 10,675 photometric points, obtained in 102 nights, which is up to now the largest database of INOV for any BL Lac object. Our statistical analysis will give new informations on the distribution of the variability time scales and other properties of this source. ",
        "conclusions": "\\label{sec:Disc} The observational results described in this paper are useful to extend the present knowledge about INOV of BL Lac objects. In an about six year long campaign we obtained a large collection of data on \\sbz, not available before for a single source, useful to develop a statistical study of the main INOV properties. We give in Table~A1 (see Appendix) the whole data set, including all 10,675 photometric measurements. They can be used for further investigations and for a comparison with other data on the same and other sources.  Our statistical analysis was essentially based on the evaluation of magnitude variation rates $|\\Delta m/\\Delta t|$ over several time intervals, selected using rather uniform criteria to minimize possible biases. We found that the resulting distribution is fully compatible with an exponential one having a mean $\\langle |\\Delta m/\\Delta t| \\rangle =$ 0.027~mag/h, corresponding to a flux variation time scale of~37.6~h. This finding implies that the probability to observe a magnitude variation rate higher than 0.2~mag/h is smaller than 10$^{-3}$, and therefore one would require more than 500 nights of observations like ours to detect an episode having such a high rate.  The interpretation of the variability of blazars is not a simple problem because it involves the description of rapidly changing processes characterized by several physical quantities, whose mean values and statistical distributions are poorly known. The fact that we found an exponential distribution for $|\\Delta m/\\Delta t|$ without any evidence of a typical time scale suggests that the INOV is essentially a stochastic process. A possibility already considered in some papers is that of a turbulent jet. A model in which relativistic electrons emit synchrotron radiation in a turbulent magnetic field was described by \\citet{Mar92}: the resulting light curves show trends and oscillations like those described in the previous sections. This agreement is, however, only qualitative and a larger observational effort should be performed to achieve a more detailed description of the turbulence parameters. For instance, the discovery of a relation between the amplitude and the duration of small oscillations on a robust statistical ground can help to model the turbulence.  Another investigation can concern the possible long term variations of INOV parameters. \\citet{Nes05} recently presented the historic light curve of \\sbz~ that shows a well apparent brightening trend since about 25 years. A study of the apparent ejection velocities of superluminal blobs in the jet $\\beta_{ej}$ (\\citealt{Bac05}) showed that it decreased from $\\sim$~15 to $\\sim$~5 in the period from 1986 to 1997. Both effects are consistent with a scenario of a precessing jet having a stable Lorentz factor $\\Gamma \\simeq 12$  but approaching the line of sight from an angle of about 5$^{\\circ}$ to 0$^{\\circ}$.5. The corresponding Doppler factor $\\delta = 1/\\Gamma (1-\\beta\\,cos\\theta)$ ($\\Gamma$ is the bulk Lorentz factor, and $\\theta$ the angle between the jet direction and the line of sight) increased from $\\sim$~13 to $\\sim$~25.  We can use this independent estimate of $\\delta$ to constrain the size of the emitting region responsible for INOV. Considering the mode $\\tau_{Fm}$ of the distribution in Fig.~\\ref{fig:histvartau} as a typical (short) timescale and assuming $\\delta \\simeq 20$ (\\citealt{Nes05}), we can derive an upper limit to the characteristic size of the emitting region in the comoving frame $r' \\simeq c \\,\\delta\\, \\tau_{Fm}/(1 + z) \\simeq 5\\times10^{16}/(1 + z)$ cm, a value which agrees well with those usually adopted in modelling BL Lac jets.  The distribution of the magnitude variation rates is useful to detect a change in the jet orientation. Indeed, if synchrotron emission is originated in a relativistic jet, the observed flux is related to the one emitted in the comoving frame~$F'$ (assumed practically steady) by the relativistic boosting: \\begin{equation}   \\label{eq:ft} F(t) = (\\delta(t))^k F' \\end{equation} where $k = 3 + \\alpha$ ($\\alpha$ is the spectral index). After converting the flux in magnitude and deriving with respect to the time: \\begin{equation}  \\label{eq:absmdotdelta} |\\dot{m}| = 1.086~k\\,|\\dot{\\delta}/\\delta| \\end{equation} If the variation of $\\delta$ is due only to the change of $\\theta$, we obtain \\begin{equation}  \\label{eq:absdeltadotdelta} |\\dot{\\delta}/\\delta| = \\beta_{app} |\\dot{\\theta}| \\end{equation} and \\begin{equation} \\label{eq:absmdotbeta} |\\dot{m}| = 1.086\\,k\\,\\beta_{app} |\\dot{\\theta}| \\end{equation} where $\\beta_{app} =(\\beta~sin \\theta)/(1 - \\beta~cos\\theta) = - d (1 - \\beta~cos\\theta)^{-1}/d\\theta$ is the apparent velocity of superluminal components along the jet. This relation suggests that in the case of a regular variation of $\\theta$  (precessing jet) it is possible to expect a positive correlation between $\\langle|\\dot{m}|\\rangle$ and $\\beta_{app}$. Under this respect, it will be important to continue to study this blazar: we expect that an increase of $\\theta$, after the very low value reached in the past years, would imply an increasing $\\beta_{app}$ and consequently a larger $\\langle|\\dot{m}|\\rangle$ (see also \\citealt{Nes05}).  This consideration supports the importance of a multifrequency approach to distinguish geometrical from physical effects affecting the emission properties of BL Lac objects. \\citet{MasMan04} pointed out how the study of both optical long-term variability and VLBI imaging can be useful for the understanding of geometrical and structural changes of the synchrotron radiation in jets of BL Lacs. Now we suggest the possibility that the mean properties of INOV, say for instance $\\langle |\\dot{m}| \\rangle $ or $ \\tau_{Fm}$, may also change on such long time scales and could be related to the kinematics of the jet derived from VLBI imaging. This kind of work requires the acquisition and storage of a large amount of INOV observations for a sample of BL Lac objects, covering time intervals of several decades. Such a great observational effort can be performed only with the collaboration of several groups, possibly working with automatic/robotic small aperture telescopes, and with the creation of homogeneous and well organized databases."
    },
    "0512/astro-ph0512647_arXiv.txt": {
        "abstract": "We have made a comparative study of morphological evolution in simulated DM halos and X-ray brightness distribution, and in optical clusters. Samples of simulated clusters include star formation with supernovae feedback, radiative cooling, and simulation in the adiabatic limit at three different redshifts, $z = 0.0, 0.10,$ and $0.25$. The optical sample contains 208 ACO clusters within redshift, $z \\leq 0.25$. Cluster morphology, within 0.5 and 1.0 h$^{-1}$ Mpc from cluster center, is quantified by multiplicity and ellipticity.  We find that the distribution of the dark matter halos in the adiabatic simulation appear to be more elongated than the galaxy clusters. Radiative cooling brings halo shapes in excellent agreement with observed clusters, however, cooling along with feedback mechanism make the halos more flattened.  Our results indicate relatively stronger structural evolution and more clumpy distributions in observed clusters than in the structure of simulated clusters, and slower increase in simulated cluster shapes compared to those in the observed one.  Within $z \\leq 0.1$, we notice an interesting agreement in the shapes of clusters obtained from the cooling simulations and observation. We also notice that the different samples of observed clusters differ significantly in morphological evolution with redshift. We highlight a few possibilities responsible for the discrepancy in morphological evolution of simulated and observed clusters. ",
        "introduction": "The hierarchical clustering is the most popular model for the Large Scale Structure (LSS) formation. The model relies on the assumption that larger structures result from the merging of smaller sub-clumps. Theoretical paradigm of the hierarchical evolution is the Cold Dark Matter (CDM) scenario which assumes that baryonic matter (stars, hot X-ray gas) evolves in the dark matter (DM) potential through violent processes. Structural evolution in cosmological objects such as galaxies or clusters of galaxies is the underlying principle in this scenario. A generic prediction of the CDM model is the nonsphericity of the DM halos. The degree of flattening of the halos evolves in cosmological time, from highly irregular at the distant past towards more regular at the present. In principle, the model prediction can be tested comparing the DM halo shapes with that of the (baryonic) matter distributions. A comparative morphological analysis between model and observation could help constraining the nature of the DM and its role on the LSS.  Melott, Chambers \\& Miller (2001; hereafter MCM) have reported evolution in the gross morphology of galaxy-clusters (quantified by ellipticity) for a variety of optical and X-ray samples for $z < 0.1$. They infer that the evidence is consistent with a low matter density universe. Using a similar shape measure as well as intra-cluster medium temperature and X-ray luminosity, Plionis (2002) has presented evidence for recent evolution in optical and X-ray cluster of galaxies for $z \\leq 0.18$. In both studies evolution is quantified by the change of cluster ellipticity with redshift. In a recent study, Jeltema et al. (2005) have reported structural evolution of clusters with redshift where cluster morphology is quantified by the power ratio method (Buote \\& Tsai 1995). Jeltema et al. used a sample of 40 X-ray clusters over the redshift range $\\sim 0.1 - 0.8$ obtained from Chandra Observatory. In spite of methodological differences, the results of these studies indicate evolution in the morphology of the largest gravitationally bound systems over a wide range of look-back time.  The observational evidence prompted concerns about the formation and evolution of structures in the CDM scenario via numerical simulations. If the results of simulations provide faithful representations of the evolutionary history of cosmological objects, then one would expect a similar trend in the structure of simulated objects. So far almost all studies of simulated clusters are focused either on understanding the nature of the background cosmology within which the present universe is evolving (Jing et al. 1995; Crone, Evrard \\& Richstone 1996; Buote \\& Xu 1997; Thomas et al. 1998; Valdarnini, Ghizzardi \\& Bonometto 1999; Suwa et al. 2003) or on understanding the distribution and shape of the DM halos in various types of simulations, e.g., simulations with or without baryons and gas physics (Dubinski \\& Carlberg 1991; Dubinski 1994; Aninos \\& Norman 1996; Tissera \\& Dominguez-Tenreiro 1998; Bullock 2001; Buote et al. 2002; Jung \\& Suto 2002; Gao et al. 2004a,b; Kazantzidis et al. 2004; Springel, White \\& Hernquist 2004; Allgood et al. 2005; Nagai \\& Kravtsov 2005; Flores et al. 2005; Libeskind et al. 2005; Maccio et al. 2005; van der Bosch et al. 2005; Zentner et al. 2005). Until recently a comparative study of morphological evolution in simulated and real clusters was absent. Floor et al. (2003) and Floor, Melott \\& Motl (2004; hereafter FMM) have investigated evolution in cluster morphology simulated with different initial conditions, background cosmology, and different physics (e.g. simulation with or without radiative cooling). They have used eccentricity as a probe to quantify evolution. Their studies, emphasizing shape in the outer regions of clusters, suggest slow evolution in simulated cluster shapes compared to the observed one. However, the studies of Floor and collaborators are indirect in a sense that they did not analyze observed clusters using the same measurement technique applied to their simulated data sets.  In this paper we make a comparative analysis between simulated and observed clusters where both data sets are juxtaposed and analyzed using the same set of structural measures. We analyze cluster morphology and its evolution using shape measures such as multiplicity ($M$) and ellipticity ($\\epsilon$) derived from the Minkowski functionals (Rahman \\& Shandarin 2003, 2004, hereafter RS03 and RS04; Rahman et al. 2004). The MFs provide a non-parametric description of the images with no prior assumptions made on the shapes of the images. The measurements based on the MFs appear to be robust and numerically efficient when applied to various cosmological studies, e.g., galaxies, galaxy-clusters, CMB maps etc. (Mecke, Buchert \\& Wagner 1994; Schmalzing et al. 1999; Beisbart 2000; Beisbart, Buchert \\& Wagner 2001; Beisbart, Valdarnini \\& Buchert 2001; Kerscher et al. 2001a, 2001b; Shandarin, Sheth \\& Sahni 2004). Various measures, constructed from the two-dimensional scalar, vector, and several tensor MFs have been described and tested in RS03 and RS04. To derive the parameters applied in this study we use the extended version of the numerical code developed in RS03 and RS04.  We study evolution in the simulated clusters in a flat CDM universe ($\\Lambda$CDM; $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$) obtained from three different sets of high resolution simulations (Motl et al. 2004). The first set has clusters simulated in the adiabatic limit, the second set contains clusters with radiative cooling (RC), and the last set includes clusters with cooling + star formation and supernovae feedback (SFF). Each sample contain DM as well as X-ray brightness distributions at three different redshifts, $z = 0.0, 0.1,$ and $0.25,$. For comparision we also analyze a sample of ACO (Abell, Corwin \\& Olowin 1989) clusters within $z \\leq 0.25$. The sample contains 208 optical clusters derived from 10-inch photographic plates taken with the 48-inch Palomar Schmidt Telescope (Tr\\`{e}vese et al. 1992; Flin et al. 1995; Tr\\`{e}vese et al. 1997; Flin et al. 2000).  The objective of our study is twofold: first, to check the efficiency of the parameters differentiating various sets of objects, and second, to explore (statistical) correspondence in the morphological properties of the distributions of DM halos, X-ray emitting gas, and optical clusters using measures that are sensitive to shape and sub-structures.  In the CDM model (satellite) galaxies are associated with the DM sub-halos that are accreted by their (current) parent halo, a bigger structure usually associated with a galaxy cluster. If this is the case statistical properties of galaxies regarding mass, sub-structure, shape etc., would show a similar trend to that of the sub-halos. On the other hand X-ray emitting hot gas, evolving in the DM background potential, would not directly follow the DM distribution because of its isotropic pressure support. Therefore, a statistical analysis of various properties of DM halos, galaxy clusters, and X-ray gas distributions will be useful to probe possible bias of luminous galaxies toward sub-halos and their correspondence with the distribution of hot gas. This is the motivation behind the second objective.  The organization of the paper is as follows: simulation technique and the observational data are described briefly in $\\S2$, a short discussion of shape measures is given in $\\S3$. The results are presented in $\\S4$ and the conclusions are summarized in $\\S5$. ",
        "conclusions": "Numerical simulations provide an unique opportunity to follow the hierarchical nature of the LSS formation in both linear and nonlinear regimes (Frenk et al. 1985, 1988; Quinn, Salmon \\& Zurek 1986; Efstathiou et al. 1988). In order to be representative of the reality, results from simulations should agree with observations. Observations provide evidence of morphological evolution in galaxy-clusters (Melott, Chambers \\& Miller 2001; Plionis 2002; Jeltema et al. 2005), simulations should show similar trend. Besides, in the CDM model luminous galaxies are associated with the DM sub-halos which reside in bigger parent halos, closely associated with galaxy clusters. According to this model statistical properties of galaxies, e.g. mass, sub-structure, shape etc., would show a similar trend to that of the sub-halos while X-ray emitting hot gas would have different properties than galaxies and sub-halos. A statistical analysis of various properties of halos, galaxy clusters, and X-ray gas could provide clues to find possible biasing of luminous galaxies toward DM sub-halos and whether or not they have any correspondence with the distribution of hot gas. With this in mind, we have studied redshift evolution of cluster morphology simulated, respectively, in the adiabatic limit, with radiative cooling, and with star formation including SN feedback at three different redshifts, $z = 0.0, 0.10,$ and $0.25$. For comparision we have also studied a sample of observed clusters containing 208 ACO clusters within redshift, $z \\leq 0.25$.  Since observed clusters are projected along the line of sight and lack the full three dimensional information we, therefore, use projected simulated clusters. Each cluster image is a 8 h$^{-1}$ Mpc frame containing 360 $\\times$ 360 pixels. Clusters are analyzed at two different density/brightness threshold levels corresponding to radii 0.5 and 1 h$^{-1}$ Mpc from the cluster center. To quantify morphological evolution we use multiplicity and ellipticity as two different probes that are sensitive to cluster sub-structures and shape.  Our results indicate that optical clusters have, in general, more sub-structures than simulated halos and X-ray brightness distributions. Cluster components, in both observed and simulated clusters, evolve with redshifts and the evolution is different at different regions from cluster centers. In terms of total multiplicity ($M_{max}$), observed clusters have stronger evolution compared to DM halos. The X-ray brightness distributions, however, show steeper evolution (than that of galaxy clusters) in dissipationless simulation.  We find that in terms of overall shape, simulations do model the observed universe in an interesting way. The simulated clusters evolve with redshift, consistent with the hierarchical formation scenario. However, observed clusters appear to be slightly more flattened at higher redshift than the simulated one indicating slower evolution in simulated objects. This may reflect some form of incompleteness in our understanding in simulating the LSS. Our results differ from those of FMM (2004) who reported that the evolution in the simulated cluster shape is significantly slower than the observed one. We not only find stronger structural evolution in simulated clusters, but also find that observed cluster shapes appear to be consistent with dissipative simulations, at least, in the redshift range $z < 0.1$. The discrepancies noted in FMM is due to the different redshift ranged probed as well as intrinsic methodological differences while comparing simulations with observations.  We note that on one hand shapes of optical clusters seems to be compatible with both the halos and X-ray brightness distributions, one the other hand, both of these components appear to be less clumpy than the distribution of galaxies. Therefore, it seems puzzling whether or not there is any correspondence between the DM halos and galaxies. The existence of any such correspondence is still a matter of ongoing debate as there are conflicting results based on systematics of numerical simulations such as nature of simulations (dissipationless or dissipative) and the effect of mass and force resolution (see Maccio et al. 2005). In the context of the CDM model we would expect that the optical clusters would have similar morphology and evolutionary trend to that of the halos and would be different than the properties of the distribution of hot gas traced in the X-ray region of the spectrum.  Within the uncertainties and systematics involved in our optical sample, the results indicate that the properties of optical clusters do not exactly represent either the distribution of the halos or that of the X-ray emitting gas in any of the simulations. We find offsets in the measured parameters, such as multiplicities and ellipticities, between observations and simulations, and are unable to find any clear signature of DM-galaxy biasing based on our morphological analysis. This may be an indication, although in no way conclusive, of the fact that these components of the LSS may represent intrinsically different populations, and galaxies may not trace the DM distributions (see Gao et al. 2004a,b; Nagai \\& Kravtsov 2005). However, this is merely a speculation and we stress that care must be exercised in interpreting our results as one must be careful in selecting proper measures, radius, mass range, and most importantly, well defined samples of clusters to have unbiased and meaningful results in any morphological analysis comparing observations and simulations.  We find that the measurements from different samples do not agree on the evolution rate. Take, for example, optical clusters with $z < 0.1$, and radius, 0.75 h$^{-1}$ Mpc. In this case, the APM sample shows $d\\epsilon/dz \\sim 1.02$. FKB, on the other hand, finds much weaker evolution, $d \\epsilon / d z \\sim 0.2$. As mentioned in FKM, the discrepancy may be due to differences in adopting cluster centers, smoothing, and applied method of shape determination.  A preliminary analysis of a sample of 800 clusters constructed from the Sloan Digital Sky Survey (SDSS) shows that ellipticity evolution of optical clusters, for $z < 0.1$ and within $\\sim 1$ h$^{-1}$ Mpc, is weaker than that of the APM clusters. The result indicates that clusters with different mass limits evolves differently. Large, massive clusters ($M \\sim 10 ^{15} M_{\\odot}$) have stronger evolution compared to the less massive clusters ($M \\sim 10^{13}-10^{14} M_{\\odot}$) (C. Miller 2004, private communication). This is an interesting observation. If it is confirmed then the scaling relation between axes ratios and mass noted in simulations (Bullock 2002; Jing \\& Suto 2002) must be modified to be consistent with observations.  The SDSS sample is uniform with a well documented selection function and high degree of completeness. We may then infer that the cluster samples discussed previously have less uniformity in mass range: the APM catalog and FKB samples are biased toward massive clusters whereas the MCM samples contain more less massive clusters. The discrepancy may also arise from the techniques applied in ellipticity estimates (see also Flores et al. 2005 in this regard). Unfortunately we are unable to check the evolution strength - mass relation for our optical sample because, apart from an approximate range, no well defined criteria has been used to sort clusters into different mass bins.  The discrepancy in the optical samples is an indication of different selection criteria used to construct the catalogs. Larger and more complete catalogs obtained from the SDSS and XMM-Newton survey may be able to shed more light into this issue. It is also likely that numerical simulations may lack crucial physics that needs to be included (see FMM for discussion). In the future we will analyze clusters simulated with various gas physics, e.g. thermal conduction and AGN heating, and compare them with the SDSS clusters. The results of these studies may give us some clues to gain better insight of the current discrepancy. \\\\ \\\\ \\noindent{\\bf Acknowledgments} We thank the anonymous referee for constructive comments and criticisms which help to improve the quality of this paper. We thank M. Plionis for providing the APM cluster ellipticity data and Scott W. Chambers for the MCM data sets. NR thanks Hume Feldman and Bruce Twarog for many useful discussions. \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f01.eps} \\caption{Contour plots of toy clusters at different brightness levels (in arbitrary scales). The multi-modal clusters have clumps with different peak brightness. For all clusters the outer line represents the percolation level where the sub-structures merge and form a single, large system. \\label{toy_img}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f02.eps} \\caption{Multiplicity as a function of contour area ($A_S$) for toy clusters as shown in Fig. \\ref{toy_img}. The circle (star) represents the effective multiplicity $\\bar{M}_{eff}$ (maximum multiplicity $M_{max}$) as defined in the text. The x-coordinates of these legends are chosen only for the convenience of demonstration. Recall that the position of the highest peak along the x-axis corresponds to the $M_{max}$ whereas $\\bar{M}_{eff}$ is obtained after averaging along the x-axis. See text for details. \\label{toy_mul}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f03.eps} \\caption{Ellipticity as a function of contour area ($A_S$) for toy models as shown in Fig. \\ref{toy_img}. Dotted and solid lines represent, respectively, $\\epsilon_{eff}$ and $\\epsilon_{agg}$. The circle (star) represents the $\\bar{e}_{eff}$ ($\\bar{\\epsilon}_{agg}$) for these toy clusters as defined in the text. Once again the x-coordinates of these legends are chosen only for the convenience of demonstration. See text for details. \\label{toy_ell}} \\end{figure}   \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f04.eps} \\caption{Multiplicity (${M}$) as a function of contour area ($A_S$) for a selection of clusters at $z = 0.50, 0.25,$ $0.10,$ and $0.0$. Two clusters from each redshift are shown. Dark, gray, and faint solid lines represent respectively, the adiabatic, radiative cooling (RC), and star formation with feedback (SFF) samples. The dark matter (DM) and X-ray clusters are shown on the left and right panels, respectively. Multiplicity is, in general, greater than 1 in the entire redshift range for clusters simulated with RC (medium line) indicating a slower evolution than in the adiabatic sample (dark line). Redshift $z = 0.5$ is taken for demonstration purpose only. See text for details. \\label{multip_area}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f05.eps} \\caption{Effective (${\\epsilon}_{eff}$) and aggregate (${\\epsilon}_{agg}$) ellipticity as a function of contour area ($A_S$) for the same clusters as in Fig. \\ref{multip_area}. Solid and dotted lines are used to represent ${\\epsilon}_{agg}$ and ${\\epsilon}_{eff}$, respectively. The color style is similar to Fig. \\ref{multip_area}. In most cases the non-spherical central part of these clusters consists of a single peak (i.e. ${\\epsilon}_{eff} = {\\epsilon}_{agg}$) whereas in the outer regions sub-clumps show various shapes. It can be see easily that the central regions of clusters in hydrodynamic simulations appear to be more regular. We notice that cluster centers are slighly more flattened than the outer parts, irrespective of the nature of simulation. Although the trend is weak. See text for details. \\label{ellip_area}} \\end{figure}   \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f06.eps} \\caption{A detailed comparison of the estimate of $\\bar{M}_{eff}$ for the adaibatic DM (left panels) and X-ray (right panels) clusters and the optical sample with ss$ = $50 h$^{-1}$ kpc within 0.5 h$^{-1}$ Mpc (panels 1) and 1.0 h$^{-1}$ Mpc (panels 2) radius. Simulated clusters are shown by (faint) horizontal lines at $z = 0.25, 0.10, 0.0$ and the optical clusters are shown by (dark) crosses. The expressions represent the best fit lines for observation (dark; top line in each panel) and simualtion (faint; second line in each panel). \\label{ad_op_mco}} \\end{figure} \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f07.eps} \\caption{A detailed comparison of the estimate of $\\bar{\\epsilon}_{agg}$ obtained from the adaibatic DM (left panels) and X-ray (right panels) clusters and from the optical sample. Presentation style is similar to Fig. \\ref{ad_op_mco}. \\label{ad_op_eco}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f08.eps} \\caption{Adiabatic sample with 50 h$^{-1}$ kpc smoothing (ss) within 0.5 h$^{-1}$ Mpc radius. Dark and gray lines are used for optical and simulated clusters, respectively. The error bar represents the error in the mean. The expression at each panel relate the evolution of the mean value of the parameter of with redshift. The strength of evolution for optical clusters are: $d\\bar{M}_{eff}/dz \\sim 0.14, \\ dM_{max}/dz \\sim 0.57,$ $d\\bar{\\epsilon}_{eff}/dz \\sim 0.27,$ and $d\\bar{\\epsilon}_{agg}/dz \\sim 0.28$. \\label{ad_sm1_ap1}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f09.eps} \\caption{Radiative cooling (RC) sample with 50 h$^{-1}$ kpc smoothing (ss) within 0.5 h$^{-1}$ Mpc radius. Presentation style is similar to Fig. \\ref{ad_sm1_ap1}. \\label{rc_sm1_ap1}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f10.eps} \\caption{Star formation with feedback (SFF) sample with 50 h$^{-1}$ kpc smoothing (ss) at 0.5 h$^{-1}$ Mpc radius. Presentation style is similar to Fig. \\ref{ad_sm1_ap1}. \\label{sff_sm1_ap1}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f11.eps} \\caption{Adiabatic sample with 50 h$^{-1}$ kpc smoothing (ss) within 1.0 h$^{-1}$ Mpc radius. Presentation style is similar to Fig. \\ref{ad_sm1_ap1}.  The strength of evolution for optical clusters are: $d\\bar{M}_{eff}/dz \\sim 0.13, \\ dM_{max}/dz \\sim 0.79,$ $d\\bar{\\epsilon}_{eff}/dz \\sim 0.22,$ and $d\\bar{\\epsilon}_{agg}/dz \\sim 0.20$. \\label{ad_sm1_ap2}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f12.eps} \\caption{Radiative cooling (RC) sample with 50 h$^{-1}$ kpc smoothing (ss) within 1.0 h$^{-1}$ Mpc radius. Presentation style is similar to Fig. \\ref{ad_sm1_ap1}.  \\label{rc_sm1_ap2}} \\end{figure}  \\begin{figure} \\epsscale{1.10} \\plotone{\\figdir/f13.eps} \\caption{Star formation with feedback (SFF) sample with 50 h$^{-1}$ kpc smoothing (ss) within 1.0 h$^{-1}$ Mpc radius. Presentation style is similar to Fig. \\ref{ad_sm1_ap1}.  \\label{sff_sm1_ap2}} \\end{figure}"
    },
    "0512/astro-ph0512537_arXiv.txt": {
        "abstract": "We present 3-5~$\\mu$m spectroscopy of the interacting system NGC~6240, showing the presence of two active galactic nuclei. The brightest (southern) nucleus shows up with a starburst-like emission, with a prominent 3.3~$\\mu$m emission feature. However, the presence of an AGN is revealed by the detection of a  broad Br$\\alpha$ emission line, with a width of $\\sim1,800$~km~s$^{-1}$. The spectrum of the faintest (northern) nucleus shows typical AGN features, such as a steep continuum and broad absorption features in the M-band.  We discuss the physical properties of the dusty absorbers/emitters, and show that in both nuclei the AGN is dominant in the 3-5~$\\mu$m band, but its contribution to the total luminosity is small (a few percent of the starburst emission). ",
        "introduction": "L-band spectroscopy ($\\sim3-4~\\mu$m) of Ultraluminous Infrared Galaxies (ULIRGs) is a powerful tool to disentangle the starburst and AGN contributions to the huge ($>10^{12}~L_\\odot$) infrared luminosity. Several spectral features can be used as indicators of one of the two components (e.g. Imanishi \\& Dudley~2000, Risaliti et al.~2005, hereafter R05). More specifically:\\\\ - a large equivalent width of the 3.3~$\\mu$m PAH emission feature ($EW_{3.3}\\sim$100~nm) is typical of starburst-dominated sources;\\\\ - a strong absorption feature at 3.4~$\\mu$m ($\\tau_{3.4}>0.2$), due to alyphatic hydrocarbon grains, is an indicator of an obscured AGN;\\\\ - a steep red continuum ($f_\\lambda \\propto \\lambda^\\Gamma, \\Gamma>2$) suggests the presence of an obscured, reddened AGN.  M-band spectra ($4.5-5~\\mu$m) of ULIRGs are available only for a small number of sources. From the analysis of the M-band emission of the nearby obscured AGN NGC~4945 (Spoon et al.~2003) strong absorption features due to CO ices are expected for obscured AGNs.  NGC~6240 is an interacting system consisting of two nuclei with a separation of 1.8~arcsec (Fried \\& Schulz~1984), corresponding to $\\sim$800~pc\\footnote{We adopt $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, e.g Spergel et al.~2003).} and with a total infrared luminosity $L_{IR}=10^{11.8}~L_\\odot$ (Genzel et al.~1998). It is optically classified as a LINER (Rafanelli et al.~1997), and no indications of an AGN, such as broad Pa$\\alpha$ or Br$\\gamma$ lines, are present in the near-IR. L-band spectroscopy, performed with a four meter class telescope, did not resolve the two nuclei, and showed a typical starburst emission, with a flat continuum, a strong 3.3~$\\mu$m emission feature (EW$\\sim$70~nm), and no absorption features (Imanishi \\& Dudley~2000). In the hard X-rays the AGN emission is dominant above 10~keV (Vignati et al.~1999), while at lower energy only the reflected component is visible, due to the high column density ($N_H\\sim2\\times10^{24}$~cm$^{-2}$) obscuring the direct component. In a recent {\\em Chandra} observation the two nuclei are clearly separated, and both show a prominent iron $K\\alpha$ emission line, with $EW>1$~keV, indicating the presence of an AGN in both nuclei (Komossa et al.~2003). This is the first clear detection of a double AGN in an interacting system.  Here we present new VLT L-band and M-band spectra of the two nuclei of NGC~6240, both showing clear AGN features. ",
        "conclusions": "We have presented 3-5~$\\mu$m low resolution spectra of the two nuclei in the Ultraluminous Infrared Galaxy NGC~6240, showing clear evidence of the presence of an AGN in both nuclei. This confirms the early detection of the double AGN obtained in the hard X-rays with {\\em Chandra} (Komossa et al.~2003).  In the southern, brighter nucleus a broad Br$\\alpha$ emission line is detected, the only known BLR evidence in the spectrum of this source.  The northern nucleus shows a steep L and M band emission, typical of a reddened ($\\tau_L\\sim$2) AGN, with possible strong CO absorption features in the M-band.  In both nuclei the AGN emission dominates in the 3-5~$\\mu$m band, but its contribution to the bolometric luminosity is small. The nuclear activity in sources like NGC~6240, i.e. where the AGN is faint (compared to the starburst) and obscured, is non-negligible only in the infrared between 3 and $\\sim5-8$~$\\mu$m and in the hard X-rays. At other wavelength the AGN is either obscured or strongly diluted by the starburst emission."
    },
    "0512/hep-ph0512078_arXiv.txt": {
        "abstract": "\\medskip We discuss effective interactions among brane matter induced by modifications of higher dimensional Einstein gravity via the replacement of Einstein-Hilbert term  with a generic function $f({\\cal{R}})$ of the curvature scalar ${\\cal{R}}$. After deriving the graviton propagator, we analyze impact of virtual graviton exchanges on particle interactions, and conclude that $f({\\cal{R}})$ gravity effects are best probed by high-energy processes involving massive gauge bosons, heavy fermions or the Higgs boson. We perform a comparative analysis of the predictions of $f({\\cal{R}})$ gravity and of Arkani-Hamed--Dvali--Dimopoulos (ADD) scenario, and find that the former competes with the latter when $f^{\\prime\\prime}(0)$ is positive and comparable to the fundamental scale of gravity in higher dimensions. In addition, we briefly discuss graviton emission from the brane as well as its decays into brane-localized matter, and find that they hardly compete with the ADD expectations. Possible existence of higher-curvature gravitational interactions in large extra spatial dimensions opens up various signatures to be confronted with existing and future collider experiments. ",
        "introduction": "The relative feebleness of gravity with respect to the weak force and its stability under quantum fluctuations, the gauge hierarchy problem, has been pivotal for introducing a number of 'new physics' models to complete the standard electroweak theory (SM) above Fermi energies. The idea \\cite{idea3,idea4,idea5} that the scale of quantum gravity can be much lower than the Planck scale, possibly as low as the electroweak scale itself \\cite{idea2,idea1} (see also the recent standard-like models found in intersecting D-brane models \\cite{ibanez}) since this extreme is not excluded by the present experimental bounds \\cite{exp}, has opened up novel lines of thought and a number of phenomena which possess observable signatures in laboratory, astrophysical and cosmological environments.  The basic setup of the Arkani-Hamed--Dimopoulos--Dvali (ADD) scenario \\cite{idea1} is that $(1+3)$--dimensional universe we live in is a field-theoretic brane \\cite{rubakov} which traps all flavors of matter except the SM singlets $e.g.$ the  graviton and right-handed neutrinos. As long as the surface tension of the brane does not exceed the fundamental scale $\\overline{M}_D$ of $D$--dimensional gravity, at distances $\\gg 1/\\overline{M}_D$ the spacetime metric $g_{A B}$ remains essentially flat. In other words, for singlet emissions (from brane) with transverse (to brane) momenta $\\left|\\vec{p}_T\\right| \\ll \\overline{M}_D$ the background spacetime is basically Minkowski. Therefore, it is admissible to expand $D$--dimensional metric about a flat background \\begin{eqnarray} \\label{metric} g_{A B} = \\eta_{A B} + 2 \\overline{M}_D^{1-D/2} h_{A B} \\end{eqnarray} where $\\eta_{A B} =\\mbox{diag.}\\left(1, -1, -1, \\cdots, -1\\right)$ and $h_{A B}$ are perturbations. The gravitational sector is described by Einstein gravity \\begin{eqnarray} \\label{ah} S_{ADD}=\\int d^{D}x\\, \\sqrt{-g} \\left\\{ - \\frac{1}{2} \\overline{M}_D^{D-2} {\\cal{R}} + {\\cal{L}}_{matter}\\left(g_{A B}, \\psi\\right)\\right\\} \\end{eqnarray} where $\\psi$ collectively denotes the matter fields localized on the brane. There are various ways \\cite{idea1} to see that the Planck scale seen on the brane is related to the fundamental scale of gravity in higher dimensions via \\begin{eqnarray} \\label{planck} \\overline{M}_{Pl} = \\sqrt{V_{\\delta}} \\overline{M}_D^{1+\\delta/2} \\end{eqnarray} which equals $(2 \\pi R)^{1/2} \\overline{M}_D^{1+\\delta/2}$ when $\\delta\\equiv D-4$ extra spatial dimensions are compactified over a torus of radius $R$. Obviously, larger the $R$ closer the $\\overline{M}_D$ to the electroweak scale \\cite{idea1}. Experimentally, size of the extra dimensions, $R$, can be as large as a small fraction of millimeter \\cite{exp}, and thus, quantum gravitational effects can already show up at experimentally accessible energy domains provided that the strength of gravitational interactions on the brane drives from higher dimensional gravity as in (\\ref{planck}). Upon compactification, the higher dimensional graviton gives rise to a tower of massive S, P and D states on the brane, and they participate in various scattering processes involving radiative corrections to SM parameters, missing energy signals as well as graviton exchange processes. These processes and their collider signatures have been discussed in detail in seminal papers \\cite{giudice,han}.  The ADD mechanism is based on higher dimensional Einstein gravity with metric (\\ref{metric}). Given the very fact that general covariance does not forbid the action density in (\\ref{ah}) to be generalized to a generic function $f\\Big({\\cal{R}},$ $\\Box {\\cal{R}},$ $\\nabla_{A} {\\cal{R}} \\nabla^{A} {\\cal{R}},$ ${\\cal{R}}_{A B} {\\cal{R}}^{A B},$ ${\\cal{R}}_{A B C D} {\\cal{R}}^{A B C D}, \\dots\\Big)$ of curvature invariants, in this work {\\it we will derive and analyze effective interactions among brane matter induced by such modifications of higher dimensional Einstein gravity, and compare them in strength and structure with those predicted by the ADD mechanism.} The simplest generalization of (\\ref{ah}) would be to consider, as we will do in what follows, a generic function $f({\\cal{R}})$ of the curvature scalar. Such modified gravity theories are known to be equivalent to Einstein gravity (with the same fundamental scale) plus a scalar field theory with the scalar field \\begin{eqnarray} \\label{phig} \\phi= \\overline{M}_{D}^{(D-2)/2} \\sqrt{\\frac{D-1}{D-2}} \\log \\left|\\frac{\\partial f}{\\partial R}\\right| \\end{eqnarray} in a frame accessible by the conformal transformation $g_{A B} \\rightarrow (\\partial f/\\partial R) g_{A B}$ \\cite{ct}. Therefore, generalized action densities of the form $f({\\cal{R}})$ are equivalent to scalar-tensor theories of gravity, and thus, matter species are expected to experience an additional interaction due to the exchange of the scalar field $\\phi$ \\cite{brans-dicke}. This is the fundamental signature of $f({\\cal{R}})$ gravity compared to Einstein gravity for which simply $f({\\cal{R}})= {\\cal{R}}$. (Though remains outside the scope of this work, see the discussions of Lovelock higher-curvature terms in \\cite{rizzo}.)  In this work we study how $f({\\cal{R}})$ gravity influences interactions among brane matter and certain collider processes to observe them. In Sec. 2 below we derive graviton propagator and describe how it interacts with brane matter. Here we put special emphasis on virtual graviton exchange. In Sec. 3 we study a number of higher dimensional operators which are sensitive to $f({\\cal{R}})$ gravity effects. In Sec. 4 we briefly discuss some further signatures of $f({\\cal{R}})$ gravity concerning graviton production and decay as well as certain loop observables on the brane. In Sec. 5 we conclude. ",
        "conclusions": "In this work we have discussed a number of phenomenological implications of $f({\\cal{R}})$ gravity in higher dimensional spacetimes with large extra spatial dimensions. In Sec. 2 we have derived graviton propagator about flat Minkowski background (which requires $f(0)$ $i.e.$ cosmological constant to vanish), and have determined how it influences interactions among the brane matter. In Sec. 3 we have listed down a set of higher dimensional operators which exhibit an enhanced sensitivity to $f({\\cal{R}})$ gravity (compared to those operators involving light fermions or massless gauge fields). Finally, in this section we have performed a comparative study of ADD and $f({\\cal{R}})$ gravity predictions and determined ranges of parameters where the latter dominates over the former. The analysis therein suggests that $f({\\cal{R}})$ gravity theories with finite and positive $f^{\\prime\\prime}(0)$ induce potentially important effects testable at future collider studies. In Sec. 4 we have discussed briefly how $f({\\cal{R}})$ gravity influences loop processes on the brane as well as decays and productions of gravitons.  The analysis in this work can be applied to various laboratory, astrophysical and cosmological observables (see \\cite{idea1} for a detailed discussion of major observables) for examining non-Einsteinian forms of general relativity in higher dimensions. The discussions presented here are far from being complete in their coverage and phenomenological investigations. The rule of thumb to be kept in mind is that higher curvature gravity influences scatterings of massive (sufficiently heavy compared to the fundamental scale of gravity) brane matter."
    },
    "0512/astro-ph0512067_arXiv.txt": {
        "abstract": "We propose a technique to test the idea that non-standard dynamics, rather than dark matter halos, might be responsible for the observed rotation curves of spiral galaxies.  In the absence of non-luminous matter, a galactic disk's rotational velocity and its vertical velocity dispersion can be used jointly to test the self-consistency of the galaxy's dynamics.  A specific illustrative example, using recent measurements of the disk kinematics of M33, shows this to be a promising approach to assess the viability of Modified Newtonian Dynamics (MOND). ",
        "introduction": "Understanding the structure and kinematics of spiral galaxies, in particular explaining their rotation curves at large galactic radii, remains one of the pressing open questions in astrophysics. Optical observations of galactic rotation curves find that rather than falling off as one would expect from galaxy models where the mass traces the observed light, the rotational velocities remain constant at large radii.  These findings are further borne out by radio observations of the 21 cm line, from HI gas in the outer parts of the galactic disk. An overview of rotation curves is provided in \\cite{Rubin01}, \\cite{Persic96} and \\cite{Salucci01}.  There is strong evidence from the CMB data for considerable amounts of non-baryonic dark matter on the cosmological scale \\citep{WMAP03}. While is it enticing to imagine that the dark matter problems on the galactic and cosmological scales have a common resolution, this need not necessarily be the case. Our focus in this paper will be on the galactic dark matter problem as manifested in the ubiquitous observation of flat rotation curves.  A number of ideas have been put forth to account for flat rotation curves.  All require some form of new physics.  These ideas include:  \\begin{enumerate}  \\item{} New particles: Missing mass in the form of galactic dark matter, most likely non-baryonic \\citep{Alcock00}.  \\item{} New interactions: Non-gravitational long-range couplings might exist, or gravitational physics might be subject to revisions over large distances.  \\item{} New dynamics: scenarios such as Modified Newtonian Dynamics (MOND) in which gravity from visible matter is the only force acting but the system's response takes on new aspects.  \\end{enumerate}  The community consensus at present prefers the dark matter hypothesis, but galactic dark matter has thus far evaded all attempts to detect it. We should strive to test, whenever possible, alternatives to the dark matter scenario.  \\subsection{The MOND Approach to the Rotation Curve Puzzle: Novel Dynamics}  Motivated by the observed spiral galaxy rotation curves, \\citet{Milgrom83} proposed a modification of the dynamics of non-relativistic matter.  This modified behavior, termed MOND for Modification of Newtonian Dynamics, is conjectured to arise only in the regime of low accelerations. MOND is a proposed modification to an object's acceleration under an applied force, such that $a=g/\\mu(x)$ where $g$ is the acceleration expected under Newtonian physics, and $x=a/a_0$ depends upon the MOND acceleration scale $a_0 \\sim 1.2 \\times 10^{-10} m/s^2$, with $\\mu(x \\gg 1) \\simeq 1, \\mu(x \\ll 1) \\simeq x$.  A commonly adopted form is $\\mu(x) = {{x}/{\\sqrt{1+x^2}}}$.  In general a gravitating system's behavior under MOND can be described by taking the Newtonian description and replacing $G$, the coupling constant, by $G/\\mu$, with the understanding that the dynamics is being altered rather than the nature of the gravitational interaction.  In this scenario the mass of a galaxy resides in the ordinary astronomical components that we can detect by their emission or absorption of electromagnetic radiation, and the galaxy's light distribution traces out its mass distribution.  A review of MOND as an alternative to dark matter is presented in \\citet{Sanders02}.  In the MOND model the response of a test particle to an applied force depends upon the magnitude of its absolute acceleration relative to a preferred frame, taken to be the local rest frame of the microwave background. ``Overacceleration'' at low values of $x$ then produces the observed rotation curves of spiral galaxies.  MOND thereby eliminates the need for dark matter, at the expense of novel dynamics at low accelerations.  MOND does a remarkably good job of fitting the rotation curves of galaxies across a wide range of surface brightness, with $a_0$ as the single free parameter \\citep{Sanders02}.  The MOND idea was recently placed on a more formal footing \\citep{Beckenstein04}, but our approach will be cast in the original phenomenological framework, in terms of $\\mu(x)$. From this standpoint, the formulation described above suggests an observational test for the self-consistency of MOND. Because $\\mu(x)$ depends on the (scalar) magnitude of a particle's total acceleration, comparing the vertical and rotational dynamics of test particles in the disk of a spiral galaxy provides a means to test for self-consistency. This paper proposes a framework for carrying out such a test and illustrates the technique with recent data from M33 \\citep{Ciardullo04}.  Other recent efforts to investigate the viability of the MOND hypothesis using kinematics include using galaxy clusters \\citep{Silk05} and globular clusters of stars \\citep{Baum05}.  Our approach differs in that, as discussed in the following section, we are checking the self-consistency of MOND rather than comparing the observed kinematics to a prediction. ",
        "conclusions": "Our objective is to propose a general technique for testing the self-consistency of MOND, using the existing M33 data as an illustrative example. The vertical and circular motions of a galaxy can be jointly used for this test. Potential weaknesses in the argument presented above include i) the assertion that the vertical scale height of galaxies is radius-independent, ii) modeling the galaxy with the form shown in equation (1), and iii) the implicit assertion that either objects overaccelerate, or they don't. The first issue can be addressed with better observations and more statistics, and the second by a more comprehensive treatment of the system's kinematics.  The third concern, namely the isotropy of MONDian dynamics, is an interesting issue. If MONDian behavior arises from a modification of inertia \\citep{Milgrom05}, then this scalar quantity will determine an object's response to {\\it any} applied force, and it will exhibit the same modified dynamics in all directions. On the other hand one might imagine that MOND only applies component by component, with a modified response only to those forces that would give rise to accelerations below the $a_0$ threshold. This could produce a difference in the radial and vertical dynamics and could perhaps account for a ratio of $\\mu_{radial}/\\mu_{vertical}$ that differs from unity. In this circumstance however a terrestrial Cavendish experiment conducted at the North or South pole should see differing effective values of $G$ in different regimes of $\\mu$.  Sensible next steps to obtaining observations that are optimally suited to the test we propose include 1) assessing the relative merits of planetary nebulae vs. integrated starlight as probes of vertical velocity dispersion, 2) selecting a favorable list of target galaxies, and 3) carrying out a set of appropriate observations. It is sensible to include, as a control, examples of high surface brightness disk galaxies which should have their inner regions in the Newtonian disk-dominated regime where $\\mu$ =1, to verify that $CP$ is constant and equal to unity for these systems. H$\\alpha$ and 21 cm observations of the velocity field might also contribute to this technique."
    },
    "0512/astro-ph0512298_arXiv.txt": {
        "abstract": "{We present the luminosity function (LF) of star clusters in M51 based on {\\it HST/ACS} observations  taken as part of the Hubble Heritage project. The clusters are selected based on their size and with the resulting 5\\,990 clusters we present one of the largest cluster samples of a single galaxy. We find that the LF can be approximated with a double power-law distribution with a break around $M_V = -8.9$. On the bright side the index of the power-law distribution is steeper ($\\alpha = 2.75$) than on the faint-side ($\\alpha = 1.93$), similar to what was found earlier for the ``Antennae'' galaxies. The location of the bend, however, occurs about 1.6 mag fainter in M51. We confront the observed LF with the model for the evolution of integrated properties of cluster populations of \\citet{gieles05b}, which predicts that a truncated cluster initial mass function would result in a bend in, and a double power-law behaviour of, the integrated LF. The combination of the large field-of view and the high star cluster formation rate of M51 make it possible to detect such a bend in the LF.  Hence, we conclude that there exists a fundamental upper limit to the mass of star clusters in M51.  Assuming a power-law cluster initial mass function with exponentional cut-off of the form $N\\,\\dr M \\propto M^{-\\beta}\\,\\exp(-M/M_C)\\,\\dr M$, we find that $M_C = 10^5\\,\\msun$. A direct comparison with the LF of the ``Antennae'' suggests that there $M_C = 4\\times10^5\\,\\msun$.  } ",
        "introduction": "There is a relation between the luminosity of the brightest star cluster in a galaxy and the total number of clusters (\\citealt{2003dhst.symp..153W}; \\citealt{2002AJ....124.1393L}), suggesting that sampling statistics is determining the luminosity of the most luminous cluster. Since the luminosity of clusters is heavily dependent on the age, a straightforward translation from most luminous to most massive is not possible.  Recently, \\citet{2003AJ....126.1836H} showed that the maximum cluster mass increases with log(age/yr) in the LMC and SMC, which can be interpreted as a size-of-sample effect.  Also \\citet{2004MNRAS.350.1503W} suggest that the maximum cluster mass in a galaxy depends on the star formation rate in the galaxy, hence the total number of clusters. This suggests that it would be {\\it physically} possible to form a super massive cluster such as W3 in NGC~7252 with a mass of $8\\times10^7\\,\\msun$ \\citep{2004A&A...416..467M} in our Milky Way, but the chance is just very small. This issue is still heavily under debate and is subject of this study.  The cluster luminosity function (LF) is a powerful tool for the study of star cluster populations. In a wide variety of environments the LF can often be well approximated by a power-law distribution: $N\\,\\dr L~\\propto~L^{-\\alpha}\\dr L$, where the index $\\alpha$ is between 1.8 and 2.4 (e.g. \\citealt{2002AJ....124.1393L}; \\citealt{2003MNRAS.343.1285D}). The shape of the LF is related to, but not necessarily identical to the cluster initial mass function (CIMF). It is important to note that it is hard to relate the observed LF directly to the underlying CIMF, since the LF contains clusters of different ages. A star cluster fades about 5 magnitudes in 1 Gyr in the $V$-band, which makes it hard to estimate the mass without knowing the age. %  The LF of clusters in the ``Antennae'' galaxies \\citep{1999AJ....118.1551W}, however, is much better approximated by a {\\it double} power-law distribution. The bright side ($M_V\\,\\la\\,-10$) has a steeper slope ($\\sim -2.7$) than the faint side ($\\sim -2$). The latter is close to the value found for other galaxies. This double power-law nature with a {\\it bend}, was interpreted by the authors as a turn-over in the {\\it mass} function.  In \\citet{gieles05b} we compared a cluster population model with various observed luminosity functions from the literature. We investigated various possible ways of detecting a truncated cluster initial mass function and the possible biases caused by extinction, disruption, variations in the cluster formation rate, etc. We concluded that a truncated CIMF will be observed as a bend in the integrated cluster luminosity function. We showed that tentative hints for a truncation are present in NGC~6946 (from \\citealt{2002AJ....124.1393L}) and M51 (from \\citealt{2005A&A...431..905B}) and are clearly not present in the SMC and the LMC (from \\citealt{2003AJ....126.1836H}).  In this work we present a greatly improved LF of clusters in M51, based on recently released deep {\\it HST} observations with the {\\it Advanced Camera of Surveys (ACS)} covering the entire disk of M51. The great resolution of the {\\it ACS} camera is exploited by selecting clusters based on their size. With this we are able to accurately select clusters, even when they are as faint as individual bright high mass stars. The improved resolution and larger field-of-view make it possible to confirm the suggestion of \\citet{gieles05b} that the LF of M51 is of a double power-law nature.   We show that the bend in the LF is not necessarily related to a corresponding turn-over in the MF, but results naturally if the CIMF is a power-law distribution truncated at the high-mass end.   In \\S~\\ref{sec:data} we describe the data, source selection and photometry. In \\S~\\ref{sec:lf} we present the LF in the three available {\\it ACS} filters of all extended objects in M51. A comparison with the model is done in \\S~\\ref{sec:model} and a discussion and the conclusions are presented in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   Does the fact that other galaxies have a single power-law LF imply that there is no upper limit to the cluster mass there? Probably not. The CIMF has to be sampled well enough to reach the critical \\mmax, only then a bend will show up  in the total cluster LF. M51 and the ``Antennae'' are forming enough clusters such that the CIMF is sampled until \\mmax, and a bend in the LF is observable. The difference between \\mmax\\ in M51 and the ``Antennae'' galaxies suggest that \\mmax\\ is environment dependent.  These environmental differences might be caused by variations in the giant molecular cloud (GMC) mass distribution. \\citet{2003ApJ...599.1049W} shows that the cloud mass function of the ``Antennae'' galaxies is truncated at higher masses than that of M51, which in turn is at much higher masses than in the Milky Way \\citep{1997ApJ...476..166W}. Other galaxies also seem to have GMC mass distributions which are truncated at the high mass end \\citep{2005astro.ph..8679R}.  A truncated GMC mass function might impose a physical limit to the maximum star cluster mass, which will be observable in galaxies with a high star/cluster formation rates."
    },
    "0512/astro-ph0512317_arXiv.txt": {
        "abstract": "The study of Be in stars of differing metal content can elucidate the formation mechanisms and the Galactic chemical evolution of the light element, Be.  We have obtained high-resolution, high signal-to-noise spectra of the resonance lines of Be II in eight stars with the high-dispersion spectrograph (HDS) on the Subaru 8.2 m telescope on Mauna Kea.  Abundances of Be have been determined through spectrum synthesis.  The stars with [Fe/H] values $>-$1.1 conform to the published general trend of Be vs.~Fe.  We have confirmed the high Be abundance in HD 94028 and have found a similarly high Be abundance in another star, HD 132475, at the same metallicity: [Fe/H] = $-$1.5.  These two stars are 0.5 - 0.6 dex higher in Be than the Be-Fe trend.  While that general trend contains the evidence for a Galaxy-wide enrichment in Be and Fe, the higher-than-predicted Be abundances in those two stars shows that there are also local Be enrichments.  Possible enrichment mechanisms include hypernovae and multiple supernova explosions contained in a superbubble.  One of our stars, G 64-37, has a very low metallicity of [Fe/] = $-$3.2; we have determined its Be abundance to look for evidence of a Be plateau.  It's Be abundance appears to extend the Be-Fe trend to lower Fe abundances without any evidence for a plateau as had been indicated by a high Be abundance in another very metal-poor star, G 64-12.  Although these two stars have similar Be abundances within the errors, it could be that their different Be values are indicators of may be indicating that a Be dispersion even at the lowest metallicities. ",
        "introduction": "The study of Be in stars has importance in several areas of astronomy.  The history of the production of Galactic Be can be seen in the increase of Be abundance with stellar metallicity, usually measured by Fe or O (e.g. Boesgaard et al.~1999) but other elements such as Ca and Mg have been employed (e.g. King 2002).  The Fe abundance is used as a surrogate for age, but the O abundance is relevant because it it directly connected to the dominant, and perhaps only, Be production mechanism: spallation reactions in the interstellar medium or in the vicinity of supernovae explosions.  The idea of spallation was invented by Reeves, Fowler \\& Hoyle (1970) and the details were first described by Meneguzzi, Audouze \\& Reeves (1971).  The basic process is that high energy ($\\sim$150 MeV) protons and neutrons bombard interstellar nuclei of C, N, and O creating lighter isotopes.  It has been suggested that the ``bullets'' and ``targets'' might be reversed near supernovae where C, N, and O nuclei could be accelerated into the ambient interstellar gas including protons and neutrons (see, for example, Duncan et al.~1997, 1998, Lemoine, Vangioni-Flam \\& Cass\\'e 1998).  Modified and expanded versions of Galactic cosmic ray spallation can be found, for example, in Ramaty \\& Lingenfelter (1999), Ramaty, Lingenfelter \\& Kozlovsky (2000), Parizot \\& Drury (1999), Parizot (2000).  Although standard Big Band Nucleosynthesis (BBN) is not expected to produce much Be (Be/H $\\leq$10$^{-18}$), some models with inhomogeneous regions in the very early universe can make up to Be/H $\\sim$10$^{-14}$ (Malaney \\& Mathews 1992, Orita et al.~1997).  This amount is near the detection threshold and efforts have been made to search for a plateau in Be similar to the Li plateau (e.g. Boesgaard et al.~1999, Primas et al.~2000a).  The plateau in Li abundance for metal-poor stars, from [Fe/H] = $-$1.5 to $-$3.5, has a small dependence on metallicity and temperature found by Novicki (2005) in a sample of 116 halo dwarfs.  A Be plateau might also show a similar dispersion.  If there is a Be plateau, it need not be the result of production of Be in an inhomogeneous Big Bang.  It could be caused by extra production of Be in superbubbles by multiple supernova outbursts (Parizot \\& Drury 1999).  The increase in Be with Fe may also show a dispersion in the Be abundance for a given Fe.  Hints of this are seen in Figure 5 of Boesgaard et al.~(1999a), particularly at [Fe/H] $\\sim$ $-$1.5 where the spread in Be is $\\sim$0.7 dex but the typical errors are $\\pm$0.10 dex.  A spread could result from differing degrees of efficiency in the formation of Be by spallation in different parts of the Galaxy.  We have made Be observations in a set of disk and halo stars to try to understand the formation and Galactic evolution of Be.  Our sample includes the very metal-poor dwarf, G 64-37, at [Fe/H] = $-$3.2 which will allow us to examine whether there is a plateau in the Be abundances.  Other stars that we observed enable us to investigate whether there is a spread in Be at given metallicity. ",
        "conclusions": "We have determined Be abundances in eight stars from high-resolution, high-S/N spectra from the 8.2 m Subaru telescope using HDS.  The stars cover the metallicity range [Fe/H] from +0.02 to $-$3.20.  The most metal-rich star, HR 8888, has the meteoritic abundance of Be, i.e.~it is undepleted.  The four stars with [Fe/H] between $-$0.3 and $-$1.1 fall perfectly along the previously established relationship between A(Be) and [Fe/H].  There are two stars, HD 94028 and HD 132475, with [Fe/H] = $-$1.5 and these both fall significantly above the observed trend of A(Be) with [Fe/H].  With HD 94028 we confirm the previous observation of a high Be abundance from Keck/HIRES observations (Boesgaard et al.~1999a).  HD 132475 was observed for Be because we had found its Li abundance to be 3$\\sigma$ higher than other Li-plateau stars at its metallicity.  The Be abundance of HD 132475 is 4$\\sigma$ above the Fe-Be trend.  (Another star, BD +23 3912, which has similar stellar parameters and a high Li abundance has normal Be, not the same high Be abundance as HD 94028 and HD 132475.)  Now there are two stars with Be abundances higher than typical at [Fe/H] = $-$1.5 by some 0.5 - 0.6 dex; this implies that there is a cosmic dispersion in the Be abundances, probably as a result of the environment around the star at the time of its formation.  The general trend in Be vs.~Fe shows that there are Galaxy-wide processes at work that increase Be and Fe over time, but the two stars above the trend show that there are local enrichments in addition.  Such Be enhancements could be caused by extra Be production by spallation in superbubbles from multiple supernova explosions (Parizot \\& Drury 2000) or in the vicinity of hypernovae (Fields et al.~2002).  There is some evidence in the most metal-poor stars that there is a plateau in the Be abundance.  We have added a third star to the ``plateau investigation'' below [Fe/H] = $-$3.0 with our observation of G 63-37 with [Fe/H] = $-$3.20. This star has A(Be) = $-$1.30 which puts it in line with the general Fe-Be trend, and {\\it not} evidence of a Be plateau.  We compared the two stars G 64-12 and G 64-37, which have similar temperatures and gravities and very low metallicities ([Fe/H] $\\sim$ $-$3.3).  They have different light element abundance patterns.  G 64-12 has higher Li than the plateau value at A(Li) = 2.40, while G 64-37 is lower than the Li plateau at A(Li) = 2.06.  Primas (2000a) found a high Be abundance in G 64-12 of A(Be) = $-$1.15, above the trend, while our A(Be) for G 64-37 of $-$1.30 is consistent with the general trend.  This adds to the evidence for a dispersion in Be at a given metallicity.  We suggest that there are two types of spallation.  One is responsible for the general increase of Be with Fe (or O) due to Galaxy-wide spallation.  The other results from local enhancements by spallation near multiple supernovae or hypernovae."
    },
    "0512/astro-ph0512629_arXiv.txt": {
        "abstract": "Dynamically significant magnetic fields are routinely observed in molecular clouds, with mass-to-flux ratio $\\lambda \\equiv (2\\pi\\sqrt{G}) \\Sigma/B \\sim 1$ (here $\\Sigma$ is the total column density and $B$ is the field strength).  It is widely believed that ``subcritical'' clouds with $\\lambda < 1$ cannot collapse, based on virial arguments by Mestel and Spitzer and a linear stability analysis by Nakano and Nakamura.  Here we confirm, using high resolution numerical models that begin with a strongly supersonic velocity dispersion, that this criterion is a fully nonlinear stability condition.  All the high-resolution models with $\\lambda \\le 0.95$ form ``Spitzer sheets'' but collapse no further. All models with $\\lambda \\ge 1.02$ collapse to the maximum numerically resolvable density.  We also investigate other factors determining the collapse time for supercritical models.  We show that there is a strong stochastic element in the collapse time: models that differ only in details of their initial conditions can have collapse times that vary by as much as a factor of 3. The collapse time cannot be determined from just the velocity dispersion; it depends also on its distribution.  Finally, we discuss the astrophysical implications of our results. ",
        "introduction": "Molecular clouds evolve under the influence of self-gravity so as to condense part of their mass into dense cores and, ultimately, stars. The presence of magnetic fields can prevent or delay condensation.  The possibility was first studied by \\cite{ms56}, who noted that the magnetic energy and the gravitational energy scale in exactly the same way with the radius $R$ of the cloud ($\\propto 1/R$) if flux freezing obtains. They argued that there was therefore a critical mass below which a cloud threaded by a particular field strength would be unable to collapse.  A more precise but less general argument was advanced by \\cite{nn78}, who studied the linear theory of a self-gravitating, isothermal, equilibrium sheet of plasma threaded by a perpendicular magnetic field. They found that magnetic fields stabilize the sheet against gravitational collapse if the mass-to-flux ratio is smaller than $1/2\\pi\\sqrt{G}$.  These results motivate the definition of a dimensionless mass-to-flux ratio, \\begin{equation} \\lambda \\equiv 2\\pi\\sqrt{G} {\\Sigma\\over{B}}, \\end{equation} where $\\Sigma$ is the column density of the sheet, and $B$ is the magnetic field strength.  The exact coefficient used to define $\\lambda$ depends somewhat on the geometry of the collapse.  Here we have chosen the coefficient most relevant to the magnetic field geometry adopted in this paper, tending to produce thin sheets, in agreement with the expectations for magnetically supported clouds.  Clouds with $\\lambda > 1$ are termed {\\it supercritical}, and clouds with $\\lambda < 1$ are termed {\\it subcritical}.  Both the \\citeauthor{ms56} and the \\citeauthor{nn78} models consider exact equilibria.  Molecular clouds are far from equilibrium, however, with near-virial, highly supersonic velocity dispersion.  These internal velocities must arise from strong turbulence.\\footnote{The most plausible alternative to turbulence, some type of weakly dissipative ordered flow, does not emerge naturally in any relevant numerical experiments that we are aware of.  The mode-mode coupling is always strong.  Even circularly polarized Alfv\\'en waves, which are exact solutions to the compressible equations of motion, suffer from a parametric instability with a dynamical decay rate \\citep{sag69,gol78}.} Turbulence might change the stability properties of the cloud, either by compressing a $\\lambda < 1$ flow until it collapses, or by providing turbulent support to a cloud with $\\lambda > 1$.  Many works have suggested that turbulence could provide support to star-forming clouds.  \\citet{cf53a} included turbulent support in their model for interstellar gaseous structures.  \\citet{ms56} pointed out that turbulence tends to decay, and that turbulence of amplitude large enough to support a cloud against self-gravity would decay especially quickly, although allowing the possibility that a strong magnetic field might perhaps allow longer lived turbulence.  The supersonic linewidths observed in molecular clouds were attributed to radial motions inside the cloud instead of turbulence by \\cite{gk74}.  \\citet{zp74} argued that if this interpretation were true for all clouds where such fluctuations are observed, the star formation rate would be too large by at least one order of magnitude.  \\citet{am75} then suggested that the observed velocity fluctuations are due to hydromagnetic waves.  By the late 1980s, this idea was widely accepted \\citep[e.g.,][]{sal87}.  In the late 1990s, however, a succession of numerical experiments \\citep{ml98,sog98,go96} strongly suggested that the damping time of turbulence in magnetized molecular clouds is close to the dynamical time.  If one accepts this, then turbulent pressure can be effective in supporting self-gravitating clouds only if it is constantly replenished, in which case the support is perhaps more readily identified with the stirring mechanism than with the turbulence itself.  Other work has tended to emphasize the role of turbulence in initiating gravitational collapse \\citep[e.g.,][]{mk04}.  Regions with a convergent velocity field will naturally tend to collapse sooner than regions with divergent velocity fields.  It seems highly likely that some parts of molecular clouds have strongly convergent velocity fields; is this ever enough to overcome the stabilizing effects of the magnetic field? Can a subcritical cloud be induced to collapse by squeezing, or can a supercritical cloud be prevented from collapsing by the introduction of turbulence?  The purpose of this paper is to investigate these questions using a simple series of numerical experiments.  The plan of the paper is as follows.  In \\S 2 we describe the experimental design, our numerical methods, and the diffusion characteristics of our code (based on the ZEUS algorithm). In \\S 3 we describe results, including a ``fiducial'' run, and the influence of physical and numerical parameters on the outcome. \\S 4 summarizes and discusses astrophysical implications. ",
        "conclusions": "Our simulations confirm that the single most important element in determining the long term gravitational stability of turbulent magnetized clouds is indeed the mass-to-flux ratio, dividing supercritical from subcritical clouds. The relevant coefficient is that corresponding to a sheet geometry, as derived by \\citet{nn78}.  Turbulent energy has comparatively little influence on the presence or absence of stability, up to Mach numbers $\\sim 10$. Subcritical clouds will develop density concentrations due to this turbulence, but under an ideal MHD regime, the consequent increase in magnetic pressure prevents further collapse. However, total turbulent energy has some influence on the lifetime of supercritical clouds, especially as the Mach number becomes large enough (of the order of $\\sim 7$ in these simulations).  More interesting is the fact that turbulence introduces a stochastic element.  The collapse time cannot be predicted with certainty from physical parameters such as the mass and field in the cloud, and the typical energy of the turbulence motions, because the random distributions of velocity and density can change the lifetime by some factor, seen to be of the order of 3 in one large sample. The resulting distribution of lifetimes has an asymmetric tail of unusually long-lived clouds.  We suggest that the existence of such a tail may introduce a bias in the observed samples of star-forming clouds. Most star formation will take place in the more frequent, shorter lived clouds, while observations of clouds will tend to focus on the fewer longer lived ones.  We have seen that the numerical resolution requirements needed to study cloud collapse are very stringent, and we expect they will be even more stringent in 3D. There is a necessity of resolving the possible equilibrium structures, such as the Spitzer sheets, which we have seen fully formed in the subcritical clouds, and partially formed during the run-up to instability of the mildly supercritical ones. The thickness of these sheets scale with the number $\\nJ$ of Jeans lengths as $\\nJ^{-2}$. Accommodating a large number $\\nJ$ of Jeans lengths inside the computational volume will therefore be numerically challenging. Increasing $\\nJ$ by only a factor of 2 requires increasing the space resolution by a factor of 4. Unless adaptive mesh refinement (AMR) is used, this requires increasing the simulation runtime by factors on the order of $64=4^3$ in 2D, and $256=4^4$ in 3D. We anticipate that AMR will be used in many of the successful simulations of core formation in the future.   Numerical stability, through the Truelove condition, sets a maximum density that can be accommodated at a given spatial resolution. Shocks in strongly turbulent flows have large compression ratios, sometimes requiring increasing resolution in order to distinguish a transient density increase due to a shock from an authentically unstable accumulation of mass able to form a collapsed object.  We have seen that artificially enforcing numerical density floors, even relatively large ones, on the order of $10^{-4}$ times the background density, had almost no influence in the evolution of the collapse. This result is again not surprising, because wide regions of small density have little influence on the dense, self-gravitating regions that undergo collapse.  Density floors can significantly speed up ideal MHD simulations, whose Courant timestep is often limited by large Alfv\\'en speeds $B/\\sqrt{4\\pi\\rho}$ in the least dense regions.  This work is limited due to the periodic boundary conditions. We believe this may have favored the collection of clumps into larger clumps until the instability can take place.  Some simulations occasionally show fast-moving clumps flowing past each other, and later merging once one of them returns through the other side of the periodic computational volume.  The periodic boundary conditions make it plausible that sooner or later, most of the mass in a given fieldline will collect into a single clump, which then can undergo instability if its mass is even slightly supercritical. In real clouds with ordered magnetic fields, clumps inside the same fieldline but moving in opposite directions are not expected to merge; however, it is improbable this will apply to all of the fieldlines and so we expect that the instability will still take place in a similar form, albeit with an additional stochastic factor in the cloud lifetime.  Two-dimensionality is also a limitation of this work. It has strongly limited the topological possibilities for the fieldlines; it is conceivable that the consequent limitations in motion have favored the collection of mass into massive sheets and other structures. Observations \\citep[e.g.,][]{gbmm90, crutcher04}, and 3D simulations and studies \\citep[e.g.,][]{basu00, glso03} indeed indicate that sheets aligned perpendicular to the magnetic field are not always the preferred possibility for the long term development of clouds.  More variety of clump shapes is expected in a 3D study. The larger variety in motions allowed by a 3D magnetic field is expected to enhance the already observed stochastic effects, and perhaps might also delay mass collection into potentially unstable structures. However, even in 3D, the simulations performed by \\citet{osg01} suggest that the stability criterion will still be dominated by the mass-to-flux ratio.  In some of our models, artificial numerical diffusion has turned an initially uniform mass-to-flux ratio $\\lambda$ into a non-uniform distribution, sometimes with striking effects on the numerical stability. While this has a numerical origin, non-uniform distributions of mass-to-flux are also expected on astrophysical grounds. For instance, turbulence provides structures and shocks with small lengthscales and strong magnetic gradients, conditions favorable to a localized, efficient ambipolar diffusion, which can redistribute mass and magnetic flux independently. Cloud collisions can also merge together portions of gas having different masses and magnetic fields. We plan to study directly the physical effect of a non-uniform mass-to-flux ratio in our future work."
    },
    "0512/astro-ph0512303_arXiv.txt": {
        "abstract": "We calculate the spectra of ultra-high-energy cosmic rays (UHECRs) in an explicit top-down model based on the decays of metastable neutral `crypton' states in a flipped SU(5) string model. For each of the eight specific $10^{th}$-order superpotential operators that might dominate crypton decays, we calculate the spectra of both protons and photons, using a code incorporating supersymmetric evolution of the injected spectra. For all the decay operators, the total UHECR spectra are compatible with the available data. Also, the fractions of photons are compatible with all the published upper limits, but may be detectable in future experiments. ",
        "introduction": "The existence of cosmic rays with energies above the Greisen-Zatsepin-Kuzmin (GZK) cutoff~\\cite{Greisen:1966jv,Zatsepin:1966jv} is one of the most important open problems in high-energy astrophysics~\\cite{Nagano:2000ve,Westerhoff}. These ultra-high energy cosmic rays (UHECRs) may be a tantalizing hint of novel and very powerful astrophysical accelerators, or they may be harbingers of new microphysics via the decays of metastable supermassive particles. {\\it It is remarkable that we still do not know whether the UHECRs originate from macrophysics or microphysics.} If there is no GZK cutoff, as suggested by the AGASA data~\\cite{AGASA}, the sources of the UHECRs would need to be local. In this case, since local magnetic fields are unlikely to have deflected significantly their directions of propagation, the UHECRs would `remember' the directions of their sources. Thus, one would expect some anisotropy in the arrival directions of the UHECRs, associated either with discrete energetic astrophysical sources nearby, such as BL Lac objects~\\cite{Westerhoff}, or the distribution of (mainly galactic) superheavy dark matter. No significant anisotropy of the UHECRs has yet been seen, but the existing experiments have insufficient statistics to exclude one at the expected level~\\cite{Evans}. On the other hand, the GZK cutoff may be present in the HiRes data~\\cite{HiRes}, in which case no exotic microphysics may be needed, and any astrophysical sources would be less restricted and more difficult to trace. The first batch of Auger~\\cite{Sommers:2005vs} data are inconclusive on the possible existence of the GZK cutoff.  Superheavy particles of the type required could have been produced gravitationally around the end of inflation~\\cite{inflation}. Particularly interesting candidates for such superheavy particles are \\lq cryptons\\rq, bound states of the fractionally-charged constituents that arise generically~\\cite{Schellekens:1989qb} in the hidden sectors of models of particle physics derived from the heterotic string~\\footnote{Such states may also be a generic feature of models constructed from intersecting D-branes.}. Cryptons arising in the hidden sector of a heterotic string-derived flipped SU(5) model may have exactly the right properties to play this role~\\cite{cryptons,EMN1}.  These \\lq flipped cryptons\\rq \\ are bound by SU(4) hidden-sector interactions, and include four-constituent meta-stable bound states called {\\it tetrons} that are analogues of the three-constituent baryons of QCD, as well as two-constituent meson-like states. Indeed, it was within this flipped SU(5) model that the confinement solution to avoiding the stringent experimental limits placed on fractional charges was first pointed out, and this model remains the only example to have been worked through in any detail~\\cite{cryptons,EMN1}.  In general, tetrons may decay through $N^{th}$-order non-renormalizable operators in the superpotential, which would yield lifetimes that are expected to be of the order of \\begin{equation} \\tau \\approx \\frac{\\alpha_{string}^{2 - N}}{m_X} \\left( \\frac{M_s}{m_X} \\right)^{2(N-3)}, \\label{lifetime} \\end{equation} where $m_X$ is the tetron mass and $M_S \\sim 10^{18}$~GeV is the string scale. The $\\alpha$-dependent factor reflects the expected dependence of high-order superpotential terms on the effective gauge coupling $g$. The mass scale associated with these states is estimated using the renormalization group for the SU(4) interactions to be $\\Lambda \\sim10^{12}-10^{13}$~GeV, just in the right range for their decays to produce the UHECRs.  The lifetimes of neutral tetrons without electric charge has been estimated to lie in the range $\\tau_0 \\sim 10^{11} - 10^{17}$~years, so that they may still be present in the Universe today, and might produce the necessary flux of UHECRs if they are sufficiently abundant.  We have shown in our previous work~\\cite{EMN1} that the mesons and charged tetrons - whose present-day abundances are subject to very stringent experimental limits - would have decayed with short lifetimes early in the Universe. In the course of studying the possible tetron lifetimes, we identified various 10$^{th}$-order superpotential operators that might govern neutral tetron decays. Thus, in this specific model of flipped cryptons we are able to go beyond generic statements regarding the injected UHECR particle spectra that may result from their decays, and make a number of specific predictions.  This enables us to address an important experimental constraint on such crypton models of the UHECRs. Although \\lq top-down\\rq \\ models such as crypton decays may appear to be natural explanations for the UHECRs (if they exist), they generically share a potential drawback. The spectra of UHECRs that they produce might be expected to have large photon fractions, in possible conflict with the observation that most of the UHECR primaries appear to be protons or nuclei. The Auger collaboration has recently set an upper limit of $26\\%$ on the photon fraction above $10^{19}$~eV~\\cite{Risse:2005hi}, and the limit may even be as low as 7-14\\%~\\cite{Sarkar}), while an upper limit of $50\\%$ at energies above $10^{20}$~GeV has been set by the AGASA collaboration~\\cite{Risse:2005jr}.  In this paper, we first give a review of relevant aspects of the flipped SU(5)  heterotic string model. We then analyze the specific primary multi-body decay modes governed by the various different 10$^{th}$-order superpotential operators found in our previous paper. We then calculate the UHECR particle spectra that would be injected by these decays using one of the most detailed and complete codes currently available~\\cite{Barbot:2003wv}, paying particular attention to the photon fraction. The total UHECR spectra obtained from the various superpotential terms do not differ greatly. On the other hand, we find that different decay operators may give rather different photon fractions, particularly at the highest energies. However, in every case, we find that the calculated spectra after supersymmetric evolution and fragmentation are compatible with the published upper limits on the photon fractions in various energy ranges, when we include the UHECR background resulting from a homogenous extragalactic distribution of sources and incorporate the pile-up expected from the GZK effect. There is no need to appeal to the cosmic radio background to absorb a significant fraction of the photons in order to bring them below the AGASA and other limits. ",
        "conclusions": "We have carried as far as is possible at present the modelling of flipped crypton decay contributions to UHECRs, including all the possible $10^{th}$-order superpotential operators. The experimental data presently available are consistent with all the decay modes possible in this crypton framework. The total UHECR spectra are consistent with a contribution from cryptons weighing between $ 2 \\times 10^{13}~{\\rm GeV}$ and $10^{12}$~GeV, although only a crypton mass $M_X \\geq 5\\cdot 10^{12}$~GeV would provide an unambiguous signal over conventional explanations. The available upper limits on the possible photon fraction do not exclude any of the crypton models we have studied.  In the future, the larger data set expected from Auger may be able to discriminate between crypton decays and other models of UHECRs, and also among different crypton models themselves. Greater statistics will enable the UHECR anisotropy to be measured with sufficient accuracy to discriminate crypton decay from a uniform distribution of astrophysical sources, and more accurate measurements of the photon fraction at higher energies might offer some discrimination between models with lepton and quark primaries, as seen by comparing Figs.~1 to 6 with Figs.~7 and 8 above.  Thus there is hope that, in the near future, we may finally learn whether UHECRs have a macrophysical origin or a microphysical origin and, in the latter case, may start to discriminate between different microphysical models."
    },
    "0512/astro-ph0512135_arXiv.txt": {
        "abstract": "We investigate the consequences of an imperfect dark energy fluid on the large scale structure. A phenomenological three parameter fluid description is used to study the effect of dark energy on the cosmic microwave background radiation (CMBR) and matter power spectrum. In addition to the equation of state and the sound speed, we allow a nonzero viscosity parameter for the fluid. Then anisotropic stress perturbations are generated in dark energy. In general, we find that this possibility is not excluded by the present day cosmological observations. In the simplest case when all of the three parameters are constant, we find that the observable effects of the anisotropic stress can be closely mimicked by varying the sound speed of perfect dark energy. However, now also negative values for the sound speed, as expected for adiabatic fluid model, are tolerable and in fact could explain the observed low quadrupole in the CMBR spectrum. We investigate also structure formation of imperfect fluid dark energy characterized by an evolving equation of state. In particular, we study models unifying dark energy with dark matter, such as the Chaplygin gas or the Cardassian expansion, with a shear perturbation included. This can stabilize the growth of inhomogeneities in these models, thus somewhat improving their compatibility with large scale structure observations. ",
        "introduction": "Dark energy is a fundamental component of the nowadays standard cosmological model. It would be  very difficult to explain the set of present days cosmological observations without it. Specifically, we refer to the luminosity-redshift relationship from observations of supernovae of type Ia (SNIa)   \\cite{Riess:1998cb,Perlmutter:1998np,Riess:2004nr}, the matter power spectrum of large scale structure as inferred from galaxy redshift surveys like the Sloan Digital Sky Survey (SDSS)   \\cite{Tegmark:2003ud} and the 2dF Galaxy Redshift Survey (2dFGRS)   \\cite{Colless:1998yu}, and the anisotropies  in the Cosmic Microwave Background Radiation (CMBR)   \\cite{Spergel:2003cb}.   Despite of its major importance in explaining the astrophysical data, the nature of dark energy is one of the greatest mysteries of modern cosmology. The simplest and most popular candidates for it are the cosmological constant (see e.g.   \\cite{Carroll:2000fy}), and minimally coupled scalar fields (see e.g. \\cite{Wetterich:1994bg,Ratra:1987rm,Caldwell:1997ii,Zlatev:1998tr}). However many other candidates were proposed based on high energy physics phenomenology  ( see e.g. \\cite{Amendola:1999er,Farrar:2003uw,Brookfield:2005td,Bagla:2002yn,Padmanabhan:2002cp,Armendariz-Picon:2000ah,Bertolami:1998dn,Boisseau:2000pr, Caldwell:1999ew,Gibbons:2003gb,Chiba:1999ka}), and many investigations on  their possible astrophysical and cosmological signature were undertaken ( see e.g. \\cite{Seljak:2004xh,Abramo:2004ji,Mota:2003tc,Evans:2004iq,Manera:2005ct,Amarzguioui:2004kc, Alam:2003fg,Mota:2003tm,Melchiorri:2002ux,Mota:2004pa,Hannestad:2002ur, Koivisto:2004ne,Nunes:2004wn,Koivisto:2005nr} ).  With so many possible candidates it is imperative to understand what are the main properties of the dark energy component that could have specific signatures in the astronomical data, and so could help us to discriminate among all these models.  In a phenomenological approach, dark energy might be mainly characterized by its equation of state $w$, its sound speed $c_s$, and its anisotropic stress $\\sigma$   \\cite{Hu:1998tj}. Much effort has been put into determining the equation of state of dark energy, in an attempt to constrain theories. The equation of state determines the decay rate of energy and thus affects both the background expansion and the evolution of matter perturbations (see e.g   \\cite{Peebles:2002gy}). An equally insightful characteristic of dark energy is its speed of sound. This does not affect the background evolution but is fundamental in characterizing the behavior of its perturbations. Hence many authors have explored its effect on the evolution of fluctuations in the matter distribution ( see e.g.   \\cite{Bean:2003fb,Sandvik:2002jz,Avelino:2002fj,Balakin:2003tk}). However, the investigation of the effects of the anisotropic stress has been widely neglected.  The main reason for disregarding the anisotropic stress in the dark energy fluid might be that conventional dark energy candidates, such as the cosmological constant or scalar fields, are perfect fluids with $\\sigma=0$.  However, since there is no fundamental theoretical model to describe dark energy, there are no strong reasons to stick to such assumption. In fact, dark energy vector field candidates have been proposed \\cite{Armendariz-Picon:2004pm,Kiselev:2004py,Zimdahl:2000zm, Novello:2003kh,Wei:2006tn}, and these have $\\sigma\\neq 0$. Of course, if dark energy  is such a vector, one might  break   the   isotropy   of   a Friedmann-Robertson-Walker universe.  However, as long as it remains subdominant,  this violation  is likely to  be observationally irrelevant    \\cite{Barrow:1997as}.   Once  dark  energy  comes  to  dominate though, one would expect an  anisotropic expansion of the universe, in conflict  with  the  significant  isotropy  of  the  CMBR   \\cite{Bunn:1996ut}. But on the other hand there appears to be hints of statistical anisotropy in the CMBR fluctuations  \\cite{Jaffe:2005pw,Bielewicz:2004en,Larson:2004vm,Schwarz:2004gk,Copi:2003kt,deOliveira-Costa:2003pu}.  Recently the possibility of viscous dark energy has gained attention  \\cite{Brevik:2004sd,Brevik:2005ue,Nojiri:2005sr, Brevik:2005bj}. These models are usually restricted to the context of bulk viscosity, although one could expect the shear viscosity to be dominant  \\cite{Brevik:2005bj}. One can allow bulk viscosity in a Friedmann-Robertson-Walker (FRW) universe, but when the shear is not neglected one has to face the difficulties of an anisotropic universe. However, shear viscosity at the perturbative level is compatible with the assumption of an isotropic FRW background. In fact the anisotropic stress perturbation is crucial to the understanding of evolution of inhomogeneities in the early, radiation dominated universe. Therefore an obviously interesting question is whether present observational data could allow for an anisotropic stress perturbation in the late universe which is dominated by the mysterious dark energy fluid.  Motivated by all these possibilities, we investigate if the possible existence of an anisotropic stress in the dark energy component would result in a specific cosmological signature which could be probed using large scale structure data, and if it would still be compatible with the latest CMBR temperature anisotropies and the matter power spectrum.  The article is organized as follows: In section II we discuss the parameters describing a general dark energy fluid with anisotropic stress. In section III we consider dark energy imperfect fluid models parameterized with a constant equation of state, sound speed and anisotropic stress. We investigate the effects on the late time perturbation evolution, in the integrated Sachs-Wolfe (ISW) effect of the CMBR anisotropies and on the matter power spectrum. In section IV we extend the analysis to models unifying dark energy with dark matter. We end the article with a summary of our findings and conclusions. ",
        "conclusions": "In this article we have investigated the effects of an anisotropic stress in the dark energy component on large scale structures.  We have parameterized the dark energy component with three variables. The equation of state determines the decay rate of dark energy, and the sound speed characterizes the evolution of its fluctuations. These two were treated as independent parameters, thus accounting for possible entropy in the fluid. In addition we allowed for shear viscosity in the linear order. We discussed the possibility to apply a Navier-Stokes type viscosity to determine the additional degree of freedom for dark energy fluctuations, the amount of shear viscosity, but we adopted the parameterization utilizing a viscosity parameter $c_{vis}^2$, motivated by the fact that it seems to generalize the familiar and well understood cosmological fluids in a natural way \\cite{Hu:1998kj}.  Using this phenomenological three parameter fluid description we investigated the effect of an imperfect dark energy fluid and of unified dark matter and dark energy models on the matter power spectrum and on the CMBR temperature anisotropies. For most models we find that free streaming effects tend to smooth density fluctuations. However, there are some exceptions, described below.  In dark energy models where $-1\\le w<0$, we found that increasing the anisotropic stress results in a swifter decay of dark energy overdensities, which is seen in the CMBR spectrum as an amplification of the ISW effect. The opposite occurs in the case of phantom dark energy ($w<-1$), for which the anisotropic stress supports the growth of overdensities and thus reduces the ISW effect. However, the impact of anisotropic stress on the CMBR spectrum can be closely mimicked by varying the sound speed of dark energy. This makes it difficult to distinguish between these two fluid properties.  In addition, we found that negative sound speeds are also consistent with observations, if shear viscosity is included. The situation that the pressure perturbation (evaluated in the comoving frame) is of the opposite sign than the density perturbation, is formally unproblematical to define, but when $c_{vis}^2=0$ it will exhibit unlimited growth of density fluctuations. However, when $c_{vis}^2>0$ this does not occur. For a suitable choice of parameters a low amplitude for the CMBR quadrupole is produced, in accordance with observations.  In models unifying dark matter and dark energy extended with shear, it is found that the anisotropic stress can stabilize the effect of the adiabatic pressure perturbation, thus slightly improving the compatibility of these models with large scale structure observations. It remains to be seen how one can loosen the constraints by allowing for an anisotropic stress. Our main objective here was to use these models as examples of dark energy with evolving $w$, $c_s$ and $c_{vis}$. The conclusion taken is that, in contrast to the simplest fluid models with constant $w$, $c_s$ and $c_{vis}$, in specific scenarios the shear stress can have consequences distinguishable with present observational data.  In general, we found that anisotropic perturbations in dark energy is an interesting possibility which is not excluded by the present day observational data. Furthermore, we found that the CMBR large scale temperature fluctuations, due to the the ISW effect, are a promising tool to constrain the possible imperfectness of the dark energy component. Even when the anisotropic stress cannot be directly measured, it can still bias measurements of other parameters, for instance the dark energy speed of sound or its equation of state."
    },
    "0512/astro-ph0512245_arXiv.txt": {
        "abstract": "We present the current performances of the AMBER / VLTI instrument in terms of differential observables (differential phase and differential visibility) and show that we are already able to reach a sufficient precision for very low mass companions spectroscopy and mass characterization. We perform some extrapolations with the knowledge of the current limitations of the instrument facility.  We show that with the current setup of the AMBER instrument, we can already reach $3\\sigma = 10^{-3}$ radians and have the potential to some low mass companions characterization (Brown dwarves or hypothetical very hot Extra Solar Giant Planets). With some upgrades of the VLTI infrastructure, improvements of the instrument calibration and improvements of the observing strategy, we will be able to reach $3\\sigma = 10^{-4}$ radians and will have the potential to perform Extra Solar Giant Planets spectroscopy and mass characterization. ",
        "introduction": "In this paper we discuss the current highest performances of Colour-Differential Interferometry (CDI) on the AMBER instrument and compare these performances to signal amplitude we computed from low-mass companion simulations.  CDI is based on {\\itshape simultaneous} interferometric observations in different spectral channels. As a high-angular resolution and high-dynamic technique, it presents two major advantages. First, the chromatic differences in visibility and phases are much less sensitive to instrumental and atmospheric instabilities, and therefore are easier to calibrate than the absolute complex visibility. Since the beginning of long-baseline optical interferometry with separated apertures, many early astrophysical results have been obtained using this self-calibration feature \\citep{1986A&A...165L..13T,1989Natur.342..520M}. Second, the colour-differential phase can be measured with an accuracy much better than the angular interferometric resolution $\\lambda/B$. For objects much smaller than the diffraction limit, it is proportional to the variation of the object photocentre with wavelength.  This paper is placed in the context of Extra Solar Planets characterization with interferometry \\citep{2005-ESP-MNRAS} and is intending to show that this technique has already the potential to get some scientific results on high contrast binaries. ",
        "conclusions": "The preliminary data reduction of bright sources observed in low spectral resolution with AMBER shows that the measured differential phases are accurate and stable enough to achieve the spectroscopy and angular separation of the most favorable Pegasi planets in a few 15 hours observations. This value should be reduced to 2 hours with the foreseen simple improvements of the VLTI. The resulting spectra would be affected by an instrumental term and/or an atmospheric chromatic differential OPD term producing a smooth $10^{-2}$ radians pattern over the K band.  When the instrumental term will be eliminated by beam commutation, the remaining differential OPD might be possible to fit in the data reduction procedure. However, only a successful use of closure phase guarantees the elimination of the differential OPD. The current quality of the VLTI does not allow accurate closure phase measurements, but this should be improved soon, when the three fringe pattern are better stabilized. We remain very optimistic about the possibility to do spectroscopy of Pegasi planets with AMBER quite sson."
    },
    "0512/astro-ph0512073_arXiv.txt": {
        "abstract": "Galactic bulges are known to harbour central black holes whose mass is tightly correlated with the stellar mass and velocity dispersion of the bulge. In a hierarchical universe, galaxies are built up through successive mergers of subgalactic units, a process that is accompanied by the amalgamation of bulges and the likely coalescence of galactocentric black holes. In these mergers, the beaming of gravitational radiation during the plunge phase of the black hole collision can impart a linear momentum kick or ``gravitational recoil'' to the remnant. If large enough, this kick will eject the remnant from the galaxy entirely and populate intergalactic space with wandering black holes.  Using a semi-analytic model of galaxy formation, we investigate the effect of black hole ejections on the scatter of the relation between black hole and bulge mass. We find that while not the dominant source of the measured scatter, they do provide a significant contribution and may be used to set a constraint, $v_{\\rm kick} \\lsim 500 \\kms$, on the typical kick velocity, in agreement with values found from general relativistic calculations. Even for the more modest kick velocities implied by these calculations, we find that a substantial number of central black holes are ejected from the progenitors of present day galaxies, giving rise to a population of wandering intrahalo and intergalactic black holes whose distribution we investigate in high-resolution N-body simulations of Milk-Way mass halos. We find that intergalactic black holes make up only $\\sim 2-3\\%$ of the total galactic black hole mass but, within a halo, wandering black holes can contribute up to about half of the total black hole mass orbiting the central galaxy. Intrahalo black holes offer a natural explanation for the compact X-ray sources often seen near the centres of galaxies and for the hyperluminous non-central X-ray source in M82. ",
        "introduction": "\\label{introduction} The discovery that galactic bulges harbour supermassive black holes whose masses are correlated with the properties of the bulge suggests a close connection between the formation of galaxies and the formation of black holes. The mass of the galactocentric black hole varies approximately linearly with the stellar mass of the bulge (\\citealt{kormendy95}; \\citealt{Magorrian98}; \\citealt{Mclure02}) and roughly as the fourth power of the bulge velocity dispersion (\\citealt{Ferrarese02}, \\citealt{Gebhardt00}, \\citealt{Tremaine02}).  Correlations with the near infrared bulge luminosity (\\citealt{Marconi03}) and with the bulge light concentration (\\citealt{Graham01}) have also been found.  These correlations are remarkable because they link phenomena on widely different scales - from the parsec scale of the black hole's sphere of influence to the kiloparsec scale of bulges - and thus point to a connection between the physics of bulge formation and the physics of black hole accretion and growth (\\citealt{Milosavljevic}). The simplest interpretation is that both, black hole and bulge growth, are driven by the same process whose nature, however, remains unclear.  Various models for the growth of black holes in galaxies have been studied (e.g. \\citealt{HaehneltRees}; \\citealt{SilkRees}; \\citealt{Cattaneo}; \\citealt{Kauffmann}; \\citealt{Ostriker}; \\citealt{Volonteri}; \\citealt{DiMatteo}, and others). In the context of a hierarchical cold dark matter universe, a plausible explanation for the tight correlation between bulge and black hole properties is that the galaxy mergers or disc instabilities that induce bulge growth via bursts of star formation also feed the central black hole (\\citealt{Kauffmann}; \\citealt{Croton05}; \\citealt{Bower05}.) A simple implementation of this model follows from assuming that, as cold gas condenses into stars, a certain percentage of the gas is forced into the centre of the galaxy and accreted by the black hole. Models based on this and related prescriptions successfully reproduce the \\magr relation (hereafter we use this term to refer to the relation between black hole mass and stellar bulge mass) (\\citealt{Croton05}; \\citealt{Bower05}; \\citealt{Malbon}).  An interesting aspect of the correlations between black hole mass and bulge properties is that they seem to apply over a range of five to six orders of magnitude in black hole mass (\\citealt{Tremaine02}; \\citealt{GebRichHo}). If a simple model of the kind just mentioned for the simultaneous growth of black holes and bulges is correct, then there are two direct consequences which we explore in this paper. The first is that black holes should exist in bulges of all luminosities including dwarf ellipticals and satellites of brighter galaxies like the Milky Way. The second is that black holes will likely merge as their hosts bulges collide.  Binary black holes orbiting each other emit gravitational waves (\\citealt{Peters}). Using quasi-Newtonian methods to study the orbital decay due to gravitational wave emission, \\cite{Fitchett} found that the  system will eventually enter a plunge phase, causing the black holes to coalesce emitting a burst of gravitational waves. \\cite{Peres} found that, in addition to transferring energy out of the emitting system, gravitational radiation can also take with it linear momentum. As a result, the centre of mass of the system recoils in a direction specified by the boundary conditions of the last stable orbit.  The astrophysical implication of this linear momentum kick (the ``rocket effect'' or ``gravitational recoil'') is that black holes may be ejected from galactic bulges if the potential is shallow and the kick is large enough (\\citealt{Madau}; \\citealt{Merritt04}; \\citealt{Enoki}). In theory, this could lead to a sizable population of extragalactic black holes which could, in principle, dominate the black hole mass function. Our aim in this work is to examine the importance of such kicks for the galactic black hole population. In particular, we consider the scatter on the \\magr relation in an attempt to constrain the relatively uncertain kick velocity, as well as the nature and spatial distribution of a possible extragalactic population of ejected black holes. We model the growth of black holes using the semi-analytic galaxy formation model of \\cite{Cole00}, using two methods for obtaining merger trees: Monte Carlo techniques and high resolution N-body simulations (in which the trajectories of recoiling black holes can be tracked.)  This paper is organised as follows. In section \\ref{gravitaionalrecoil}, we review the physics of the gravitational recoil. In section \\ref{gfintro}, we describe how we model the formation and ejection of black holes in our semi-analytic model. In section \\ref{standalone}, we determine the effect of black holes ejected from the progenitors of present day galaxies on the \\magr relation. In section \\ref{nbodyintro}, we use a set of high resolution N-body simulations of galactic halos to track the location of ejected black holes. We conclude in Section \\ref{conclusion} where we discuss the possible consequences of an extragalactic black hole population. ",
        "conclusions": "\\label{conclusion}  The gravitational recoil of merging black holes is an important physical effect in a universe built up hierachically through the repeated merging of galactic units. Whenever galaxies that host black holes merge, the black holes themselves will coalesce and there exists the potential for the remnant black hole to be ejected from the galaxy, provided the recoil velocity is high enough and the galactic potential shallow enough.  We have examined the role that gravitational recoil plays on the demographics of black holes. We find that the process of ejecting black holes from galaxies is efficient if bulges and black holes grow both in galaxy mergers and as a result of discs becoming unstable, because in this case nearly all bulges contain a black hole. In models where disc instability is ignored, this process is not as efficient because fewer bulges and associated black holes exist and because black hole - black hole mergers tend to occur late when the galactic potential wells are deeper. In the former case, black hole ejections produce a significant contribution to the scatter in the \\magr relation.  In fact, conservative estimates of the scatter in the observed relation constrain the recoil prefactor velocity to $v_{\\rm pf} < 500\\, \\kms$, which is consistent with the general relativitistic calculations of \\cite{Favata} and \\cite{Blanchet05}.  Is there any empirical evidence for the kind of black hole processes present in our model? \\cite{Coccato05} have recently reported the discovery of a black hole in NGC~4435 whose mass is smaller than about 20\\% of the value expected from the \\magr relation. Black holes with a smaller than expected mass arise naturally in our model, although masses as extreme as this could be rare. For example, in our model with $v_{\\rm pf} = 300\\, \\kms$, only $2-3\\%$ of black holes have a mass that deviates as much or more from the mean relation as that of the black hole in NGC~4435. However, if $v_{\\rm pf} = 500\\, \\kms$ this fraction raises to $\\sim 20\\%$.  Indirect evidence for black holes with masses below those expected from the \\magr relation has been presented by \\cite{ColbertMushotzky}. They argue that the compact X-ray sources often seen near the centre of elliptical and spiral galaxies could be black holes of mass $\\sim 10^2-10^4 M_{\\odot}$. Interestingly, these sources are often displaced from the centre by a few hundred parsecs. Similarly, \\cite{Neff} have discovered a group of off-centre compact X-ray sources in the merger remnant NGC 3256 which, they argue, could be intermediate-mass black holes. In our model, these objects might be identified with the black holes of infalling satellites or with recently merged black holes that have been kicked out of the galactic centre, their growth stunted as a result.  To investigate the spatial distribution of black holes in Milky-Way like galactic halos we used a set of N-body simulations to track both the satellite galaxies that host black holes and also the black holes that are ejected from their host galaxies. We find that the black hole mass function is bimodal, being composed of two overlapping populations.  The lower mass population consists of black holes at the centre of orbiting low mass satellites that have not undergone recent mergers, while the slightly higher mass population is composed of wandering black holes that have been ejected from mergers that formed the central galaxies and the more massive satellites. Among the latter population, we find a few supermassive ($<10^{6} h^{-1}M_{\\odot}$) black holes that were ejected from the main progenitor of the central galaxy sufficiently early such that the bulge of the central galaxy has had enough time to regrow a sizable black hole.  There is also an intergalactic population consisting of black holes whose recoil velocity was large enough not only to unbind them from their host galaxy but also from its halo.  In the future it may be possible to detect the formation of wandering black holes directly using instruments such as \\textit{LISA}\\footnote{\\texttt{http://lisa.jpl.nasa.gov/}} that can directly measure the gravitational radiation emitted in the black hole -- black hole merger that ejects the remnant. In the meantime, detecting these black holes presents an interesting challenge. In practice, an ejected intergalactic or intrahalo black hole is likely to bring along a small cusp of stars tightly bound to it as it escapes the halo's potential. These stars would provide the only measurable way of detecting such black holes. Most likely, there would not be any gas that could be accreted and hence the black holes would not be observable in the conventional way. However, as the orbits of the stars which the black hole brought with it from the bulge decay, the stars may plunge into the black hole. In addition to emitting a burst of gravitational radiation, this type of infall would tidally disrupt the star and create a small accretion disc which would radiate according to the standard physics of accretion discs. The gravitational radiation signature and accretion disc emission would provide the best way to identify these black holes.  \\cite{Magain} have recently discovered a quasar, HE0450-2958, which has no visible host galaxy and, at first sight, is a good candidate for an escaped black hole. \\cite{Haehnelt05} have indeed suggested that the quasar may have been ejected from a nearby ultraluminous infrared galaxy (ULIRG) either through the gravitational recoil process we are considering here or through a gravitational slingshot associated with three or more black holes involved in the merger responsible for the ULIRG. However, \\cite{Hoffman} have argued that the quasar is much too far from the companion galaxy to have been ejected with a velocity of $300\\, \\kms$ and so favour a slingshot mechanism.  Ultraluminous X-ray sources (ULXS) or micro-quasars have long been regarded as candidates for intermediate-mass black holes, $M_{\\rm BH} \\sim 100-1000 M_{\\odot}$, radiating near the Eddington limit (see e.g. \\citealt{Fabbiano,Mushotzky} and references therein). This is exactly the kind of object that would be naturally identified with the orbiting intrahalo black holes predicted by our model. This interpretation, however, has two difficulties. Firstly, ULXs tend to be associated with star-forming regions and their frequency seems to be correlated with the galactic star formation rate (see e.g. \\citealt{Ward}; \\citealt{Kaaret} and references therein). These facts lend support to the view that ULXs are stellar mass black holes emitting beamed radiation at highly super-Eddington rates. The second arguement against identifying ULXs with the intrahalo black holes in our model is the difficulty of finding a suitable source of material for the black hole to accrete. There are, however, some examples of ULX's, most notably the ULX in M82 (\\citealt{Kaaret01}), that are probably much too bright to be explained even as exotic super-Eddington luminosity stellar mass black holes (\\citealt{King01}). \\cite{King05} call objects like this ``Hyperluminous X-ray sources'' (HLX) and argue that these objects are precisely the intrahalo black holes associated with satellites in our model which switch on when they come close to the galactic centre. Since the exact mechanism by which these black holes would be activated is uncertain, we cannot predict how common this phenomenon should be. However, our model contains, in principle, a plentiful supply of intermediate mass black holes orbiting in the halo of every galaxy which is sufficient to account for the presence of a few ULXs or HLXs in most galaxies."
    },
    "0512/astro-ph0512590_arXiv.txt": {
        "abstract": "Of the 228 sources in the Molonglo Southern 4\\,Jy Sample (MS4), the 133 with angular sizes $<35''$ have been imaged at 5\\,GHz at $2-4''$ resolution with the Australia Telescope Compact Array.  More than 90\\% of the sample has been reliably optically identified, either on the plates of the UK Schmidt Southern Sky Survey or on $R$-band CCD images made with the Anglo-Australian Telescope.  A subsample of 137 sources, the SMS4, defined to be a close southern equivalent of the northern 3CRR sample, was found to have global properties mostly consistent with the northern sample.  Linear sizes of MS4 galaxies and quasars were found to be consistent with galaxy-quasar unification models of orientation and evolution. ",
        "introduction": "\\label{sec.intro}  In the preceding paper (\\citealt{paper1}, hereafter Paper~I), the Molonglo Southern 4\\,Jy sample (MS4) was defined according to the following selection criteria:  declination $-85^{\\circ} < \\delta < -30^{\\circ}$; flux density S$_{\\rm 408\\,MHz} > 4.0$\\,Jy; Galactic latitude $|b| > 10^{\\circ}$; not in the Magellanic Cloud regions; and not a known Galactic source. The sample thus defined contained 228 radio sources.  These sources were all imaged at 843\\,MHz with the Molonglo Observatory Synthesis Telescope (MOST) at a resolution of  $43'' \\times 43'' \\cosec |\\delta|$, giving radio positions accurate to around 1$-$2\\arcsec, and estimates of angular size.  A subsample of 137 sources, the SMS4, was defined to have S$_{\\rm 178\\,MHz} > 10.9$\\,Jy, making it a southern equivalent of the northern 3CRR sample \\citep{lrl}.  This sample was found to have properties mostly consistent with those of 3CRR, though with slight, possibly significant differences in angular-size distributions and source densities, in the sense that SMS4 has slightly higher source density and larger angular sizes than 3CRR.  In this paper we present higher-resolution radio images of those MS4 sources of smaller angular size (Section~\\ref{section.atca}), as well as optical identifications based on the Schmidt plates and on $R$-band CCD images (Section~\\ref{optid.ms4}).  Section~\\ref{sec.indcomm} contains notes on individual sources.  In Section~\\ref{sec.sampcont} we use the radio and optical data to compare derived properties of MS4, SMS4, and 3CRR, such as their radio luminosity distributions, quasar fractions, and radio linear sizes. ",
        "conclusions": "Studies of powerful, steep-spectrum radio sources have for the last forty years drawn largely on the 3C sample and its revisions the 3CR and 3CRR.  It is therefore important to know how representative this sample is, especially considering the comparative scarcity of radio sources of high flux density.  An independent southern sample, the SMS4, has been constructed from the MS4 sample, and compared with 3CRR. The most important result is that, within the errors, most comparisons of similarly measured quantities show the properties of the 3CRR and its southern equivalent, the SMS4, to be very similar.  The MS4 and SMS4 samples have also been used to test some models of relativistic beaming.  The results appear to be consistent with the model of \\citet{linsiz.gkkw} which predicts that quasar-galaxy differences are affected by evolution as well as orientation."
    },
    "0512/gr-qc0512034_arXiv.txt": {
        "abstract": "We investigate the formation via tunneling of inflating (false-vacuum) bubbles in a true-vacuum background, and the reverse process. Using effective potentials from the junction condition formalism, all true- and false-vacuum bubble solutions with positive interior and exterior cosmological constant, and arbitrary mass are catalogued. We find that tunneling through the same effective potential appears to describe two distinct processes: one in which the initial and final states are separated by a wormhole (the Farhi-Guth-Guven mechanism), and one in which they are either in the same hubble volume or separated by a cosmological horizon. In the zero-mass limit, the first process corresponds to the creation of an inhomogenous universe from nothing, while the second mechanism is equivalent to the nucleation of true- or false-vacuum Coleman-De Luccia bubbles. We compute the probabilities of both mechanisms in the WKB approximation using semi-classical Hamiltonian methods, and find that -- assuming both process are allowed -- neither mechanism dominates in all regimes. ",
        "introduction": "It has long been appreciated that in a field theory with multiple vacua -- including some models of cosmological inflation -- the nucleation of true-vacuum bubbles in a false-vacuum background can and does occur.  The study of such transitions, with and without gravity, was pioneered by Coleman and collaborators \\cite{Coleman:1977py, Callan:1977pt,Coleman:1980aw} and has become a large enterprise.  Real understanding of the {\\em reverse} process, nucleation of false-vacuum (inflating) regions in a background of (non-inflating) true-vacuum, has, however, been somewhat more elusive.  It has been proposed that his may occur by the same Coleman-DeLuccia (CDL) instanton responsible for true-vacuum nucleation~\\cite{Lee:1987qc,Banks:2002nm,Garriga:1993fh}, by the tunneling of a small false-vacuum bubble through a wormhole to become an inflating region (the Farhi-Guth-Guven, or `FGG' mechanism)~\\cite{Farhi:1989yr,Fischler:1990pk,Fischler:1989se}, or by thermal activation~\\cite{Garriga:2004nm,Gomberoff:2003zh}.  This paper comprises the second in a series studying the general process of the nucleation of inflating regions from non-inflating ones. In the first~\\cite{Aguirre:2005sv} we cataloged and interpreted all single-bubble thin-wall solutions with an interior false-vacuum de Sitter (`dS') space, and discovered and investigated an instability in such bubbles to non-spherical perturbations. In this paper we attempt to unify the treatment of both false- and true-vacuum bubble nucleations, via the CDL, FGG, and thermal activation mechanisms, in the thin-wall limit.  We find that these can all be studied within a single framework based on the junction condition potentials developed by Guth and collaborators \\cite{Farhi:1989yr,Blau:1986cw} and further generalized by Aurilia et. al. \\cite{Aurilia:1989sb}~\\footnote{The study of thin wall junctions is a vast subject~\\cite{Berezin:1982ur,Maeda:1981gw,Ipser:1983db,Aurilia:1984cm,Sato:1986uz,Berezin:1987bc}. In this paper, we will use the notation of Aurilia et. al.~\\cite{Aurilia:1989sb}.}. This allows us to both catalog all true- or false-vacuum bubble spacetimes, and to calculate tunneling exponents using the semi-classical Hamiltonian formalism of Fischler et al.~\\cite{Fischler:1990pk,Fischler:1989se}.  Understanding the quantum mechanical~\\footnote{It can be shown that this process cannot occur classically unless the weak energy condition is violated~\\cite{Farhi:1986ty,Dutta:2005gt,Aguirre:2005sv}.} genesis of inflating regions is very important in assembling a picture of spacetimes containing fields with multiple false vacua, and in understanding how inflation might have begun in our past.  These are related because if inflation can begin from a non-inflating region like our own, then {\\em our} inflationary past may have nucleated from non-inflation, and this raises troubling questions~\\cite{Dyson:2002pf,Albrecht:2004ke} if spawning inflation is less probable than spawning a large homogeneous big-bang region. This is indeed suggested by singularity theorems showing that inflating false vacuum regions must be larger than the {\\em true} vacuum horizon size \\cite{Penrose:1964wq,Vachaspati:1998dy} according to some observers \\cite{Aguirre:2005sv}. The FGG mechanism provides a potential loophole~\\cite{Albrecht:2004ke} because according to an observer in the background true vacuum spacetime, only a region the size of the black hole event horizon is removed.  There have, however, been lingering questions about whether the Farhi-Guth-Guven~\\cite{Farhi:1989yr} ``tunneling\" process can actually occur. The oldest objection is the fact that the euclidean tunneling spacetime is not a regular manifold~\\cite{Farhi:1989yr}. A more modern objection comes from holography: in the FGG mechanism, an observer in the background spacetime only sees a small back hole, whereas the inflating region ``inside\" should be described by a huge number of states~\\cite{Bousso:2004tv, Banks:2002nm}. This entropy puzzle was recently considered by Freivogel et. al.~\\cite{Freivogel:2005qh}, who have used the AdS/CFT correspondence to study thin-walled dS bubbles embedded in a background Schwarzschild-Anti-de Sitter space (Alberghi et. al. have also used the ADS/CFT correspondence to study charged vacuum bubbles~\\cite{Alberghi:1999kd}). They find that bubbles containing inflating regions which reside behind a wormhole are represented by mixed states in the boundary field theory. This resolves the entropy puzzle, and also implies that inflating regions hidden behind a wormhole cannot arise from a background spacetime by any unitary process, including tunneling.  It does not, however, suggest why semi-classical methods break down, nor how we should interpret the seemingly-allowed tunneling.  The formalism that we outline in this paper indicates that there are two ways to interpret tunneling through the effective potential of the junction conditions. The existing interpretation (the FGG mechanism) requires that the wall of a false-vacuum bubble (and in some cases of true-vacuum bubbles) must tunnel through a wormhole to produce an inflating region. In this paper, we use the global properties of the Schwarzschild-de Sitter spacetime to show that there is another interpretation corresponding to a mechanism that does not require the existence of a wormhole.  In this mechanism, a small bubble of true- or false-vacuum, which would classically collapse, instead tunnels to a large bubble that exists outside of the {\\em cosmological} horizon of the background spacetime. Consequently, this mechanism exists only in spacetimes with a positive cosmological constant. The zero-mass limit of this mechanism correctly reproduces the tunneling exponent for both true- and false-vacuum CDL bubbles~\\cite{Coleman:1980aw,Lee:1987qc}. In light of the objections to the FGG mechanism, this new process may be an alternative, in which case the formation of inflating false-vacuum regions by tunneling is forbidden in flat spacetime. On the other hand, these may just be two competing processes, and we will directly compare the tunneling exponents under this assumption.  In section~\\ref{classy}, we classify the possible thin-wall true and false one-bubble spacetimes using the effective potential formalism. We then introduce the possible tunneling mechanisms and outline the calculation of the tunneling exponents for the various possibilities in section~\\ref{tunneling}. We compare the tunneling rates for the allowed processes in section~\\ref{comparison}, interpret our results in section~\\ref{bl},  and conclude in section~\\ref{conclusions}. ",
        "conclusions": "We have catalogued all possible spherically symmetric, thin-wall, one-bubble (true- and false-vacuum) spacetimes with positive cosmological constant and have provided an exhaustive list of the possible quantum transitions between these solutions. Although there are undoubtedly many more possibilities as one relaxes the assumptions of spherical symmetry and a thin wall, this analysis should provide guidance in searching for more realistic processes.  The effective potentials of the junction condition formalism which were used to construct this catalog clearly indicate the existence of a region of classically forbidden radii separating bound solutions from unbound solutions. There are seemingly two processes which correspond to quantum tunneling through this same region, which we refer to as the L and R tunneling geometries. Both processes begin with a bound solution, which might be fluctuated by the background dS spacetime as we have assumed in Section~\\ref{comparison}. This bound solution then evolves to its classical turning point, where it has a chance to tunnel to an unbound solution, which is typically either through a wormhole in the case of the L tunneling geometries (the Farhi-Guth-Guven, or FGG, mechanism) or through a cosmological horizon in the case of the R tunneling geometries.  The R tunneling geometries without a wormhole have a relatively clear interpretation in terms of the transition of a background spacetime to a spacetime of a different cosmological constant. Indeed, the zero-mass limit corresponds exactly to the nucleation of true- and false-vacuum CDL (Coleman-De Luccia) bubbles, correctly reproducing the radius of curvature of the bubble at the time of nucleation, as well as the tunneling exponent.  The L tunneling geometries (FGG mechanism) have a rather perplexing interpretation, which is most clearly seen by studying the zero mass limit. This corresponds to absolutely nothing happening in the background spacetime, while a completely topologically disconnected universe containing a CDL bubble of the new phase is created from nothing. The massive L tunneling geometries also have an element of this creation from nothing. Before the tunneling event, there is no wormhole, but after the tunneling event, there is a wormhole behind which is a large (eventually infinite) region of the old phase surrounded by a bubble of the new phase. It is unclear how we are to interpret this as the transition of a background spacetime to a spacetime of a different cosmological constant, since the background spacetime remains completely unaffected save for the presence of a black hole.  We have found that the sign of the Euclidean action is opposite for the L and R tunneling geometries, and while the second order constraints on the momenta introduce a sign ambiguity, it is unclear how to correctly fix the signs in light of the existence of two seemingly different processes for tunneling in the same direction through the same potential. A complete explanation of these processes may well require the resolution of some very deep problems in quantum cosmology.  If we take the stance that the L and R tunneling geometries are in competition as two real descriptions of a transition between spacetimes with different cosmological constants,  then we must directly compare their relative probabilities. We have shown in Section~\\ref{comparison} that the zero-mass solution is always the most probable for either the L or R tunneling geometries, and that the L tunneling geometry will be dominant when $2 B_{CDL} > 3\\pi/\\Lambda_{+}$. Therefore, if one is considering drastic transitions of the cosmological constant, the zero-mass FGG mechanism will be the dominant mechanism for upward fluctuations and the nucleation of true-vacuum CDL bubbles will be the dominant mechanism for downward fluctuations. This situation upsets the picture of fluctuations in the cosmological constant satisfying some kind of detailed balance \\cite{Lee:1987qc, Banks:2002nm}.  It does, however, help to explain how spawning an inflationary universe from a non-inflating region might be a feasible cosmology \\cite{Albrecht:2004ke}. In the picture that we have presented, both the L and R tunneling geometries are constructed by carving some volume out of the background spacetime and filling it with the new phase. The size of this region is in some sense a measure of how special the initial conditions for inflation are. In the case of the R tunneling geometries, a huge number of the states of the background spacetime must be put into the false vacuum at high cost in terms of the probability of such a fluctuation occurring~\\cite{Banks:2002nm}. The L tunneling geometries avoid this cost by fluctuating new states already in the false vacuum (seemingly a non-unitary process as discussed by Frievogel et. al.~\\cite{Freivogel:2005qh}), with the result that beginning inflation is no longer prohibitively difficult. The question of how much of the background spacetime must make the transition to the false vacuum is therefore crucial to determining exactly how special the initial conditions for inflation are. Unfortunately, detailed balance and this resolution of the paradoxes associated with the initial conditions for inflation are seemingly incompatible, but hopefully future work will yield further insight into the old but still interesting theory of vacuum transitions."
    },
    "0512/astro-ph0512523_arXiv.txt": {
        "abstract": "The origin of Galactic CRs up the knee energy remains unanswered and provides strong motivation for the study of $\\gamma$-ray sources at energies above 10~TeV. We discuss recent results from ground-based $\\gamma$-ray Cherenkov imaging systems at these energies as well as future observational efforts in this direction. The exciting results of H.E.S.S. give clues as to the nature of Galactic CR accelerators, and suggest that there is a population of Galactic $\\gamma$-ray sources with emission extending beyond 10~TeV. A dedicated system of Cherenkov imaging telescopes optimised for higher energies appears to be a promising way to study the multi-TeV $\\gamma$-ray sky. ",
        "introduction": "The origin of hadronic Cosmic-Rays (CRs) remains an unsolved problem since their discovery nearly 100 years ago. A major area of focus concerns the Galactic component of CRs, which is thought to be dominant at least up to the {\\em knee} ($\\sim10^{15}$~eV; 1 PeV), the energy at which a spectral break is seen in the CR spectrum. It is generally believed that shell-type supernova remnants (SNRs) are responsible for the acceleration of the Galactic component of CRs based on energy budget considerations (of SNRs and of the Galactic CRs), and the good agreement between the observed properties of the CR spectrum and diffusive shock acceleration (DSA) theory. While DSA theory is able to explain many aspects of the observed Galactic CRs, a controversial aspect concerns the maximum particle energies attained \\cite{Kirk:1,Drury:1}. A maximum energy limit of $\\sim10^{14}$~eV is generally predicted \\cite{Lagage:1} which falls somewhat short of the knee energy. Considerably higher energies approaching the {\\em ankle} ($\\sim10^{18}$~eV) are thought possible if one accounts for significant fluctuations in the magnetic field \\cite{Bell:1,Lucek:1}. Alternative source scenarios have also been considered such as local Gamma-Ray-Bursts (GRB) \\cite{Waxman:1,Dermer:1}. In the conventional context of SNR acceleration of CRs, so-called Superbubbles, which benefit from the combined effects of many SNR shocks and possibly also from the stellar winds of OB associations, may also accelerate CRs to well beyond the knee \\cite{Drury:2,Bykov:1,Parizot:1}.  We are therefore left with critical questions concerning the origin of CRs at and above the knee. Unfortunately, identifying localised sources of Galactic CRs at such energies is impossible by direct detection since the CR trajectories are are randomised by Galactic magnetic fields. However, the interaction of CRs with local matter can produce a flux of $\\gamma$-rays\\footnote{and neutrinos} in the multi-GeV to multi-TeV energy band which act as an accessible {\\em tracer} of CR accelerators. Moreover, the energy spectrum of the $\\gamma$-ray flux is expected to follow that of the parent CR spectrum at the site of $\\gamma$-ray production, thus providing important information concerning the CR acceleration properties. We note that the interaction of CRs with ambient matter proceeds with a typical inelasticity of order $\\sim$0.15. A source of $\\gamma$-rays, when interpreted in the above framework, is therefore a source of hadronic CRs of energies roughly a factor 10 higher. Establishing beyond doubt a hadronic origin for a $\\gamma$-ray flux is non-trivial since we must often deal with additional leptonic sources of $\\gamma$-ray production such as inverse-Compton scattering and non-thermal Bremsstrahlung. These latter components are often diminished at higher energies due to the strong cooling suffered by electrons in magnetised post-shock environments. Morphology differences between the hadronic and leptonic components are also expected, due to a number of physics issues such as energy-dependent diffusion. Spatially-resolved spectra and energy-resolved morphology studies of $\\gamma$-ray sources are necessary to disentangle the hadronic and leptonic components. In particular, establishing the spectra and morphology of sources of $\\gamma$-rays at energies 10~TeV and above is vital to understanding where and how hadronic CR acceleration to the knee is taking place.  The detection of TeV $\\gamma$-rays is the exclusive domain of ground-based methods, and a major goal of the field is to establish the types of sources responsible for Galactic CRs. Given the maximum energy problem discussed above, it is not sufficient however just to establish numerous $\\gamma$-ray source types and their spectral slope, but to clearly establish the maximum energies of their $\\gamma$-ray fluxes well into the multi-TeV domain. The performance of ground-based $\\gamma$-ray detectors at energies $E>10$~TeV will therefore attract increasing interest over the coming years. ",
        "conclusions": "Studies of $\\gamma$-ray sources at energies above 10~TeV are vital if we are to understand how and where Galactic cosmic-rays are accelerated to  and above the knee. Since the present ground-based telescopes such as H.E.S.S. are optimised for energies $E\\leq$10~TeV, it is suggested that a new dedicated system of telescopes be employed to open up the $\\gamma$-ray sky at higher energies. Monte-Carlo simulations of a {\\em cell} of four modest-size telescopes employing the same stereoscopic imaging technique as in H.E.S.S. and the HEGRA IACT-System suggest that sufficient sensitivity is achieved in the $E>10$~TeV range. These encouraging results, from rather conservative simulations, suggest that the stereoscopic imaging approach is a viable method to explore the multi-TeV sky. More detailed simulations are now underway."
    },
    "0512/astro-ph0512009_arXiv.txt": {
        "abstract": "{}{We observed the center of the supernova remnant Vela Jr in radio continuum in order to search for a counterpart to the compact central X-ray source \\cxou , possibly a neutron star candidate which could be the remnant of the supernova explosion.} {Observations were made with the Australia Telescope Compact Array at 13 and 20 cm. Spectral indices were obtained using flux density correlations of the data which were spatially filtered to have the same u-v coverage. A multiwavelength search for counterparts to the compact central X-ray source was made.} {We compiled a new catalogue of 31 small diameter radio sources, including the previously known source PMN J0853-4620, listing the integrated flux densities at 20 cm and, for half of the sources, the flux densities at 13 cm with the corresponding spectral indices. All sources are unresolved at the present angular resolution except for Source 18, which is clearly elongated and lies strikingly close to \\cxou . Our observations show no evidence for the existence of a pulsar wind driven nebula associated with the point X-ray source. Furthermore, Source 18 has a thermal spectrum with index $\\alpha = +0.8 \\pm 0.4$ ($S \\propto \\nu^{\\alpha}$), and appears to be the counterpart of the optical source Wray 16-30. In spite of the absence of \\OIII \\ emission lines as reported in the literature, we find that this object could be explained as a low emission planetary nebula belonging to the ``butterfly'' morphological class.} {We conclude that if the radio source 18 is actually a planetary nebula, then \\cxou \\ is more likely to be related to it rather than to Vela Jr.}   \\keywords {ISM: supernova remnants -- supernova remnants: individual: Vela Jr -- planetary nebula: individual: Ve 2-27 -- stars: individual: Wray 16-30 -- X-rays: individual: \\cxou~ -- catalogues}    \\titlerunning{Interior of Vela Jr at radio wavelengths} \\authorrunning{\\textsc{E. Reynoso et al.}} ",
        "introduction": "Rotation powered pulsars can lose their energy in the form of a wind of electrons and positrons, creating a synchrotron emitting nebula or {\\it plerion}, generally called a ``pulsar wind nebula'' (PWN). In the last few years, new data with unprecedented spatial resolution have revealed that PWNe are highly structured objects which can fall into a variety of morphological classes depending on the properties and evolutionary state of the pulsar and its surroundings \\citep{bg04}. The morphology of a PWN depends on how the wind particles are confined. Therefore, the study of PWNe are of particular interest since their properties can provide information on both, the pulsar's wind and the surrounding interstellar medium.  In some cases, neutron stars (NS) are not detected as ordinary radio pulsars. Several X-ray sources with no radio counterpart and very high X-ray to optical flux ratios have been detected projected on the interior of supernova remnants (SNR). These sources, generally called ``compact central objects'' \\citep[CCO; see][]{pskg}, are believed to be NS with peculiar characteristics \\citep[e.g.][] {frail,gpz}, or simply normal pulsars with unfavorable beaming \\citep{bj}. In recent papers, we have reported the results of an {\\hi} \\ and radio continuum survey towards candidate NSs to search for their imprints in the interstellar medium and confirm or reject their physical association with the host SNR \\citep{iafeusyd1, iafeusyd2, geminga}. In this paper we analyze another CCO-SNR possible association: the X-ray source \\cxou \\  and the SNR Vela Jr.  Vela Jr (G266.1--1.2, RX J0852.0--4622) is a shell-like SNR, about 2\\deg in diameter, first discovered with ROSAT \\citep{asch}. The name ``Vela Jr'' is due to its position at the south-east corner of the Vela SNR, and to a possible age of $\\sim 700$ yr based on the detection of $\\gamma$-ray emission from the radioactive isotope $^{44}$Ti \\citep{iyetal}. However, subsequent observations show that this source is likely to be older and ten times more distant (2 kpc) than was originally believed \\citep{shepmta}. \\citet{crb} reported a radio counterpart based on the 2.42 GHz survey of \\citet{dshj}. More recently, \\citet{dg} find that several features believed to be extensions to the radio shell are actually unrelated to Vela Jr. A detection with H.E.S.S. of TeV $\\gamma$-ray extended emission, whose spatial distribution correlates with the X-ray observations, has recently been reported \\citep{hess}.  Several authors noted the presence of an X-ray point-like source close to the SNR center: \\citet{asch99} based on the ROSAT All Sky Survey, \\citet{pskg} using the Chandra Spectrometer, \\citet{shepmta} based on ASCA data, and \\citet{mere} with BeppoSAX observations. This point-like source is named slightly differently according to the different authors and has independent entries listed in SIMBAD. Most frequently, it is referred to by its Chandra identification \\cxou. Due to its location, this X-ray source is believed to be the NS remaining after the supernova explosion. This identification of \\cxou \\ as a NS appears to be supported by the very high X-ray-to-optical flux ratio \\citep{pskg}. Although the distance to \\cxou \\ is poorly known, if it is located at 2 kpc the emission region and luminosity are similar to the CCO detected in Cas A. \\citet{pelliz} found a faint \\ha \\ nebula at a position compatible with \\cxou, which could be explained as a bow shock driven by a NS. However, \\citet{iyetal2} claim that the very high ejecta velocity derived for Ti implies that the SN was of type Ia, in which case no compact remnant is expected. The nature of \\cxou \\ thus remains uncertain.  \\citet{iafeusyd3} observed a $\\sim 30^\\prime$ diameter field towards \\cxou \\ with the Australia Telescope Compact Array\\footnote{The ATCA is part of the Australia Telescope, which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO} (ATCA; Frater, Brooks \\& Whiteoak 1992) in the $\\lambda$21 cm line and at 20 cm in the radio continuum. The observations revealed a compact continuum source and an elongated region ($30^\\prime \\times 14^\\prime$) of enhanced emission at the position of the X-ray source. If confirmed, this could be a PWN and would be good evidence that the compact X-ray source is a neutron star. However, this enhancement is only marginally above the noise level of the image (less than 3$\\sigma$ rms), making the detection of the nebula doubtful. In this paper, we analyze new observations of the central region of Vela Jr performed at 13 and 20 cm, which confirm the detection of the compact source. A new interpretation of the source origin is presented.  \\medskip ",
        "conclusions": "We have made a high resolution, high sensitivity image of the center of the SNR Vela Jr at 13 and 20 cm. The aim of the project was to detect a radio counterpart to the X-ray source \\cxou, a neutron star candidate for the compact remnant of the SN explosion. There is no evidence for the existence of a PWN associated with \\cxou. The closest radio feature detected is a compact source of thermal origin, named Source 18 in this paper, which appears to be the counterpart of the optical source Wray 16-30 detected also at infrared wavelengths as IRAS 08502-4606. In spite of the lack of \\OIII \\ emission lines, we believe there is good evidence that Wray 16-30 is a PN, probably belonging to the morphological class known as ``butterfly'' PNe. If so, it is also possible that \\cxou \\ is physically associated with it. In that case, \\cxou \\ would then not be the CCO of SNR Vela Jr. This interpretation is strengthened by the radio frequency results presented here, which do not indicate an association between \\cxou \\ and Vela Jr.  The angular resolution of the present observations is insufficient to show whether Source 18 has an internal structure similar to the optical images. Higher resolution data could be obtained with the ATCA in an extended configuration and at a higher frequency.   A by-product from this study is the discovery of 30 new, uncatalogued unresolved sources, which may be useful for comparison with high energy searches. For almost half of these sources we were able to measure the flux densities at both frequencies, 1348 and 2368 MHz, and have computed spectral indices using the flux-flux method. The remaining sources were too weak at 13 cm to derive reliable results. Among those sources for which spectral indices could be measured, six were found to be extragalactic and five are clearly thermal, including the previously known source PMN J0853-4620. The only one with a spectral index close to 0.6, typical of stars with mass outflows, is Source 18.  \\bigskip"
    },
    "0512/astro-ph0512379_arXiv.txt": {
        "abstract": "We present the discovery of the first X-ray counterpart to a Rotating RAdio Transient (RRAT) source.  \\src\\ is a relatively highly magnetized (B $\\sim 5\\times10^{13}$\\,G) member of a new class of unusual pulsar-like objects discovered by their bursting activity at radio wavelengths. The position of \\src\\ was serendipitously observed by the {\\sl Chandra} ACIS-I camera in 2005 May.  At that position we have discovered a pointlike source, \\srcx, with a soft spectrum well fit by an absorbed blackbody with $N_H = 7^{+7}_{-4} \\times 10^{21}$ cm$^{-2}$ and temperature $kT=0.12 \\pm 0.04$\\,keV, having an unabsorbed flux of $\\sim2 \\times 10^{-12}$\\,ergs\\,cm$^{-2}$\\,s$^{-1}$ between 0.5 and 8 keV.  No optical or infrared (IR) counterparts are visible within $1''$ of our X-ray position.  The positional coincidence, spectral properties, and lack of an optical/IR counterpart make it highly likely that \\srcx\\ is a neutron star and is the same object as \\src.  The source showed no variability on any timescale from the pulse period of 4.26~s up to the five-day window covered by the observations, although our limits (especially for pulsations) are not particularly constraining.  The X-ray properties of \\srcx, while not yet measured to high precision, are similar to those of comparably-aged radio pulsars and are consistent with thermal emission from a cooling neutron star. ",
        "introduction": "\\label{intro}  The discovery of a new class of ``Rotating RAdio Transients'' (RRATs) has recently been reported by McLaughlin et al.~(2005).  These objects, 11 so far identified, are characterized by repeated radio bursts with durations between 2 and 30 ms and average intervals between bursts ranging from 4 minutes to 3 hours.  Their dispersion measures (DMs) place them within the Galactic plane at distances from 2 to 7 kpc.  If bursts are periodic, periods can be found from the greatest common divisor of the differences between burst arrival times.  For ten of the sources, this calculation results in periods between 0.4 and 7 seconds, suggesting that the objects are rotating neutron stars. The periods measured for the RRATs are longer than those of most normal radio pulsars and more similar to those of the populations of X-ray dim isolated neutron stars (XDINSs; Haberl~2004) and magnetars (Woods \\& Thompson~2006).  For the three sources with the highest bursting rates, period derivatives, $\\dot P$, have been measured.  No binary motion is detected. If the $\\dot P$ values are interpreted as due to magnetic dipole spin-down, they imply characteristic ages and magnetic field strengths in the general range of pulsars.  In this paper we report the X-ray detection of \\src, the first detection at other wavelengths of any of the RRATs. This source has a 4.26-s period, a relatively high inferred characteristic surface dipole magnetic field strength of $5\\times10^{13}$~G, a characteristic age $P/2\\dot{P} = 117$~kyr, and a spin-down luminosity of $3\\times10^{32}$~ergs~s$^{-1}$. The distance of this source inferred from its DM using the electron-density model of Cordes \\& Lazio~(2002) is 3.6~kpc, with considerable uncertainty.  \\src\\ is characterized by radio bursts of average duration 3~ms, with one burst detected every $\\sim$~3~minutes. This object was fortuitously in the ACIS-I field of a 30 ks {\\sl Chandra} observation toward the Galactic supernova remnant G15.9+0.2 (Reynolds et al., in preparation).  The brightest source on any of the 6 CCD chips of the {\\sl Chandra} field, besides G15.9+0.2 itself, is coincident to within $2''$ with the radio position of \\src\\ (whose error ellipse has semimajor axes $5'' \\times 32'' $).  The positional coincidence, as well as properties we describe below, make us confident that this new source, which we designate \\srcx, is the X-ray counterpart to the radio-bursting source \\src. ",
        "conclusions": "We have discovered the first X-ray counterpart to a Rotating RAdio Transient source, \\srcx.  The X-ray source is well described by a thermal spectrum consistent with emission from a cooling neutron star of age $10^4-10^5$~years, broadly consistent with the characteristic age of \\src.  The X-ray properties also suggest possible connections to the population of X-ray dim isolated neutron stars and to the transient magnetar XTE J1810--197. A search for an X-ray modulation at the aliased radio pulse frequency was unsuccessful. No variations were seen in the X-ray flux on longer time scales either, and no optical and infrared counterparts to the source have been found.  Deeper X-ray observations are required to search for pulsations, bursts, and in general to clarify the nature of the RRATs."
    },
    "0512/astro-ph0512465_arXiv.txt": {
        "abstract": "We characterize the mass-dependent evolution in a large sample of more than 8,000 galaxies using spectroscopic redshifts drawn from the DEEP2 Galaxy Redshift Survey in the range $0.4 < z < 1.4$ and stellar masses calculated from $K$-band photometry obtained at Palomar Observatory. This sample spans more than 1.5 square degrees in four independent fields.  Using restframe $(U-B)$ color and [OII] equivalent widths, we distinguish star-forming from passive populations in order to explore the nature of ``downsizing''---a pattern in which the sites of active star formation shift from high mass galaxies at early times to lower mass systems at later epochs.  Over the redshift range probed, we identify a mass limit, $M_Q$, above which star formation appears to be quenched. The physical mechanisms responsible for downsizing can thus be empirically quantified by charting the evolution in this threshold mass. We find that $M_Q$ decreases with time by a factor of $\\approx$3 across the redshift range sampled according to $M_Q \\propto (1+z)^{3.5}$.  We demonstrate that this behavior is quite robust to the use of various indicators of star formation activity, including morphological type. To further constrain possible quenching mechanisms, we investigate how this downsizing signal depends on local galaxy environment using the projected 3$^{rd}$-nearest-neighbor statistic $D_{p,3}$ which is particularly well-suited for large spectroscopic samples. For the majority of galaxies in regions near the median density, there is no significant correlation between downsizing and environment.  However, a trend is observed in the comparison between more extreme environments that are more than 3 times overdense or underdense relative to the median.  Here, we find that downsizing is accelerated in overdense regions which host higher numbers of massive, early-type galaxies and fewer late-types as compared to the underdense regions.  Our results significantly constrain recent suggestions for the origin of downsizing and indicate that the process for quenching star formation must, primarily, be internally driven. By quantifying both the time and density dependence of downsizing, our survey provides a valuable benchmark for galaxy models incorporating baryon physics. ",
        "introduction": "The redshift interval from $z \\approx 1$ to $z=0$ accounts for roughly half of the age of the universe and provides a valuable baseline over which to study the final stages of galaxy assembly. From many surveys spanning this redshift range, it is now well-established that the global star formation rate (SFR) declines by an order of magnitude \\citep[e.g.,][]{broadhurst92, lilly96, cowie99, flores99, wilson02}.  An interesting characteristic of this evolution in SFR is the fact that sites of active star formation shift from including high mass galaxies at $z \\gtrsim 1$ to only lower mass galaxies subsequently. This pattern of star formation, referred to by \\citet{cowie96} as ``downsizing,'' seems contrary to the precepts of hierarchical structure formation and so understanding the physical processes that drive it is an important problem in galaxy formation.  The observational evidence for the downsizing of star formation activity is now quite extensive.  The primary evidence comes from field surveys encompassing all classes of galaxies to $z \\approx 1$ and beyond \\citep{BE00, bell05:SFR, bauer05, juneau05, faber05, borch06}. However, the trends are also seen in specific populations such as field spheroidals, both in their stellar mass functions \\citep{fontana04,bundy05} and in more precise fundamental plane studies \\citep{treu05:downsize, vdwel05} which track the evolving mass-to-light ratio as a function of dynamical mass.  The latter studies find massive spheroidals completed the bulk of their star formation to within a few percent prior to $z\\simeq$1, whereas lower mass ellipticals continue to grow by as much as 50\\% in terms of stellar mass after $z \\sim 1$. Finally, detailed analyses of the spectra of nearby galaxies suggest similar trends \\citep{heavens04, jimenez05}.  Downsizing is important to understand as it signifies the role that feedback plays in the mass-dependent evolution of galaxies.  As a consequence, its physical origin has received much attention theoretically.  Recent analytic work by \\citet{dekel04}, for instance, suggests that the distinction between star-forming and passive systems can be understood via several characteristic mass thresholds governed by the physics of clustering, shock heating and various feedback processes. Some of these processes have been implemented in numerical and semi-analytic models, including mass-dependent star formation rates \\citep{menci05}, regulation through feedback by supernovae \\citep[e.g.,][]{cole00, benson03, nagashima04} and active galactic nuclei (AGN) \\citep[e.g.,][]{silk98, granato04, dekel04, hopkins05:red, croton05, delucia05, scannapieco05, bower05, cattaneo06}.  However, most models have, until now, primarily addressed the mass distinction between star-forming and quiescent galaxies as defined at the present epoch \\citep[e.g.,][]{kauffmann03:bimodal}. Quantitative observational measures of the {\\em evolving mass dependence} via higher redshift data have not been available.  This paper is concerned with undertaking a systematic study of how the mass-dependent evolution of galaxies progresses over a wide range of epochs.  The goal is to quantify the patterns by which evolution proceeds as a basis for further modeling.  Does downsizing result largely from the assembly history of massive early-type galaxies or is there a decline in the fraction of massive star-forming systems?  In the quenching of star formation, what are the primary processes responsible and how are they related to the hierarchical framework of structure assembly as envisioned in the CDM paradigm?  Does downsizing ultimately result from internal physical processes localized within galaxies such as star formation and AGN feedback, or is it caused by external effects related to the immediate environment, such as ram pressure stripping and encounters with nearby galaxies in groups and clusters?  In this paper, we combine the large spectroscopic sample contained in the DEEP2 Galaxy Redshift Survey \\citep{davis03} with stellar masses based on extensive near-infrared imaging conducted at Palomar Observatory to characterize the assembly history and evolution of galaxies since $z \\approx 1.2$.  Our primary goal is to quantify the downsizing signal in physical terms and test its environmental dependence so that it is possible to constrain the mechanisms responsible.  A plan of the paper is as follows.  Section $\\S$\\ref{data} presents the observations and characterizes the sample while $\\S$\\ref{methods} describes our methods for measuring stellar masses, star formation activity and environmental density.  We discuss how we estimate errors in the derived mass functions in $\\S$\\ref{mfn_description} and present our results in $\\S$\\ref{results}. We discuss our interpretations of the results in $\\S$\\ref{discussion} and conclude in $\\S$\\ref{conclusions}.  Where necessary, we assume a Chabrier IMF \\citep{chabrier03} and a standard cosmological model with $\\Omega_{\\rm M}=0.3$, $\\Omega_\\Lambda=0.7$, $H_0=100 h$ km~s$^{-1}$~Mpc$^{-1}$ and $h = 0.7$. ",
        "conclusions": "Using a large sample that combines spectroscopy from the DEEP2 Galaxy Redshift Survey with panoramic near-IR imaging from Palomar Observatory, we have investigated the mass-dependent evolution of field galaxies over $0.4 < z < 1.4$.  We have constructed stellar mass functions for active  and quiescent populations, defined in several ways, and divided into different samples according to accurate measures of the environmental density determined from the extensive spectroscopic data.  We summarize our conclusions below:  \\begin{itemize} \\item The mass functions of active and quiescent galaxies integrated over all environments conclusively demonstrate a downsizing signal.  We quantify this by charting the evolution in a ``quenching mass,'' $M_{Q}$, which describes the mass scale above which feedback processes suppress star formation in massive galaxies.  We find that $M_{Q} \\propto (1+z)^{3.5}$ with a factor of $\\approx$3 decrease across the redshift range probed.  \\item The growth in the abundance of quiescent or ``early-type'' galaxies occurs first at the highest masses and then proceeds to lower mass systems.  The relative abundance of early-types with $M_* \\sim 10^{10}$\\msun~has increased by a factor of $\\approx$3 from $z \\sim 1.2$ to $z \\sim 0.55$, whereas the total mass function exhibits little evolution in shape and normalization (less than 0.2--0.3 dex).  This implies that early-type systems result largely via the transformation of active star-forming galaxies, indicating that ``dry mergers'' are not a major feature of their assembly history.  \\item Alternative ways of dividing active and quiescent galaxies, including the use of [OII] equivalent widths and HST-derived morphologies, show qualitatively similar mass-dependent evolution and quenching. Interestingly, we observe that morphological evolution appears to take place on longer timescales than changes in the apparent star formation rate which operate at lower mass scales at each redshift.  \\item For the majority of galaxies in our sample, downsizing shows little dependence on environment.  An environmental signal is apparent when the extremes of the density distribution are compared.  In this case, downsizing in high-density regimes appears moderately accelerated compared to low-density ones, with values of $M_{tr}$ lower by a factor of $\\sim$2.  \\item We discuss several possibilities for the origin of downsizing based on our results.  We clearly rule out a scenario in which the incidence of star formation decreases uniformly for galaxies at all masses.  The weak density dependence also argues against explanations that rely on the accelerated assembly of structure in dense environments, favoring internal mechanisms instead.  \\item Through comparisons to the expected behavior of dark matter halos, we argue that the dynamical timescale resulting from the growth of structure is at least 5 times longer in galaxies hosted by halos with $\\log M/M_{\\odot}\\approx 12.5$ ($\\log M_*/M_{\\odot} \\approx 10.8$) compared to $\\log M/M_{\\odot}=14.5$ ($\\log M_*/M_{\\odot} \\approx 11.6$).  Because global dynamical scales are also independent of halo mass at a given redshift, this suggests that the quenching mechanism is strongly mass-dependent with the potential for different physical processes acting in different mass ranges.  \\end{itemize}  There are two obvious avenues for further studies of downsizing.  In a forthcoming paper (Bundy et al., in preparation) we discuss the constraints on merging and the growth of galaxies determined by our observations of the total mass function.  This will help dissect the role of merging in driving downsizing.  In the near future, it will also be possible to chart the incidence of AGN among the galaxy population and compare it to the incidence of star formation to probe the link between quenching and AGN.  The significant Chandra follow up observations currently underway in the EGS will make that field particularly exciting for such work.  Other efforts from the DEEP2 collaboration have provided new insight on environmental trends \\citep{cooper06} and will focus on precise measures of the evolving star formation rate (Noeske et al., in preparation)."
    },
    "0512/astro-ph0512186_arXiv.txt": {
        "abstract": "{} {We study in this paper the ice composition in the envelope around intermediate-mass class I Young Stellar Objects (YSOs).} {We performed a spectroscopic survey toward five intermediate-mass class I YSOs located in the Southern Vela molecular cloud in the $L$ (2.85--4.0 $\\mu$m) and $M$ (4.55--4.8 $\\mu$m) bands at resolving powers $\\lambda/\\Delta \\lambda=$~600-800 up to 10,000, using the Infrared Spectrometer and Array Camera mounted on the {\\em Very Large Telescope-ANTU}. Lower mass companion objects were observed simultaneously in both bands.} {Solid H$_{\\mathrm{2}}$O at 3~$\\mu$m is detected in all sources, including the companion objects. CO ice at 4.67 $\\mu$m is detected in a few main targets and one companion object. One object (\\object{LLN 19}) shows little CO ice but strong gas-phase CO ro-vibrational lines in absorption.  The CO ice profiles are different from source to source. The amount of water ice and CO ice trapped in a water-rich mantle may correlate with the flux ratio at 12 and 25 $\\mu$m. The abundance of H$_2$O-rich CO likely correlates with that of water ice. A weak feature at 3.54 $\\mu$m attributed to solid CH$_\\mathrm{3}$OH and a broad feature near 4.62 $\\mu$m are observed toward \\object{LLN 17}, but not toward the other sources. The derived abundances of solid CH$_{\\mathrm{3}}$OH and OCN$^{-}$ are $\\sim$10 $\\pm$ 2\\% and $\\sim$1 $\\pm$ 0.2\\% of the H$_{\\mathrm{2}}$O ice abundance respectively. The H$_{\\mathrm{2}}$O optical depths do not show an increase with envelope mass, nor do they show lower values for the companion objects compared with the main protostar. The line-of-sight CO ice abundance does not correlate with the source bolometric luminosity.} {Comparison of the solid CO profile toward \\object{LLN 17}, which shows an extremely broad CO ice feature, and that of its lower mass companion at a few thousand AU, which exhibits a narrow profile, together with the detection of OCN$^{-}$ toward \\object{LLN 17} provide direct evidences for local thermal processing of the ice.} ",
        "introduction": "Dust grains play an important role in the evolution of clouds from protostellar cores to circumstellar disks. Since dust grains are the main source of opacity, they control the thermal balance of clouds. The surfaces of cold grains act as heat sink for highly exothermic reactions to occur (e.g., formation of H$_2$) or provide sites for atoms and molecules to freeze on.  The freeze-out of molecules like CO is found to be important for regulating the gas phase chemistry of other species (e.g., \\citealt{Bergin1997ApJ...486..316B}; \\citealt{Bergin2001ApJ...557..209B}; \\citealt{Jorgensen2004A&A...416..603J}). The frozen atoms and molecules accumulate on top of a refractory core (silicates, carbonaceous compounds) and form an icy mantle or react with other species to synthesize more complex molecules.  The relative chemical composition of this ice mantle is well determined after three decades of studies using both ground-based and space-borne telescopes.  The core-mantle grain model is supported by spectropolarimetry studies (e.g., \\citealt{Holloway2002MNRAS.336..425H}). Solid H$_{\\mathrm{2}}$O, CO, and CO$_{\\mathrm{2}}$ abound in most lines of sight where ices are detected \\citep{deGraauw1996A&A...315L.345D}. Sometimes, minor species such as CH$_{\\mathrm{4}}$ ($\\sim$2\\%), HCOOH ($\\sim$2\\%), OCN$^{-}$ ($\\sim$0.2--1\\%) (e.g., \\citealt{vanBroekhuizen2005A&A}), and H$_{\\mathrm 2}$CO (3--6\\%) are found (e.g., \\citealt{Boogert1998A&A...336..352B} ; \\citealt{Keane2001A&A...376..254K}; \\citealt{Boogert2004ApJS..154..359B}).  By contrast, the presence of other minor constituents such as NH$_{\\mathrm{3}}$ is controversial; its abundance relative to water ice is likely less than $\\sim$10\\% (e.g., \\citealt{Ehrenfreund2000ARA&A..38..427E}; \\citealt{Dartois2001A&A...365..144D}; \\citealt{Dartois2002A&A...394.1057D}; \\citealt{Lacy1998ApJ...501L.105L}; \\citealt{Taban2003A&A...399..169T}). It should be emphasized that most of these studies refer to high mass young stellar objects.  Solid methanol (CH$_{\\mathrm{3}}$OH) is a particular case. It epitomizes the importance of molecular solids in the understanding of the gas phase chemistry. The radiative association of CH$_3^+$ and H$_2$O is an inefficient gas-phase process that yields methanol abundances of $\\sim~10^{-9}$ relative to H$_2$ while abundances of 10$^{-7}$--10$^{-6}$ have been found in hot cores (e.g., \\citealt{Blake1987ApJ...315..621B}; \\citealt{Sutton1995ApJS...97..455S}). The prevalent view is that the high abundance of gas phase methanol come from the release of large amounts of frozen methanol formed on grain surfaces (\\citealt{Charnley1995ApJ...448..232C}; \\citealt{vanderTak2000A&A...361..327V}; \\citealt{Horn2004ApJ...611..605H}).  Methanol ice abundance relative to water ice is found to vary from less than 3\\% w.r.t. water ice in quiescent regions up to 30\\% around massive protostars \\citep{Dartois1999A&A...342L..32D}. If all the methanol ice is released in the gas phase, the abundance of methanol in the gas phase with respect to H$_2$ will amount to 3 $\\times$ 10$^{-7}$ -- 3 $\\times$ 10$^{-6}$ assuming that the abundance of water ice with respect to H$_2$ is $\\sim$10$^{-5}$ (e.g., \\citealt{Whittet2003dge..conf.....W}). A similar situation is found for low-mass objects.  \\cite{Chiar1996ApJ...472..665C} set a stringent limit of 5\\% of methanol with respect to water ice for sources located in the Taurus molecular cloud, while \\cite{Pontoppidan2003A&A...404L..17P} found abundant methanol ice (14--25\\% of water ice) in 4 out of $\\simeq$ 40 envelopes around protostars observed with the VLT. Possible formation routes of solid methanol are also disputed. The formation rate of solid methanol by hydrogenation of CO ice in the absence of energy input (i.e. hot atoms, UV or particle irradiation) measured in laboratory experiments remains controversial (e.g., \\citealt{Hiraoka2002ApJ...577..265H}; \\citealt{Hiraoka2005ApJ...620..542H}; \\citealt{Watanabe2004ApJ...616..638W}).  The advent of 8m class telescopes equipped with large format arrays opens up the opportunity to study large samples of low and intermediate-mass sources. We present here the first observations of molecular ice features in the $L$ (2.8--4.1 $\\mu$m) and $M$ (4.5--5.1 $\\mu$m) bands toward class I intermediate-mass young stellar objects (YSOs) located in the Vela molecular cloud complex.  The spectra were obtained in the context of a large programme using the {\\it Infrared Spectrometer And Array Camera} (ISAAC) mounted at the {\\it Very Large Telescope} ANTU (VLT-ANTU) of the {\\it European Southern Observatory} (ESO). Two major absorption features are observable with ground-based telescopes, along with some weak features.  The first strong feature is centered around 3.01 $\\mu$m ($\\sim$ 3300 cm$^{-1}$) and is usually attributed to the stretching mode of solid H$_{\\rm 2}$O.  The study of the solid-water profile has been used to better explore the ice structure in the water matrix (e.g., \\citealt{Smith1989ApJ...344..413S}; \\citealt{Smith1993MNRAS.263..749S}; \\citealt{Maldoni1998MNRAS.298..251M}). The feature shows a broad excess absorption beyond 3.2 $\\mu$m whose origin is still unclear, although scattering by the larger grains (0.1--1 $\\mu$m in radius) in the size distribution is the best candidate (\\citealt{Smith1989ApJ...344..413S}; \\citealt{Dartois2001A&A...365..144D}).  The other important feature is the solid-CO band at 4.67 $\\mu$m (2140 cm$^{-1}$), whose profile is sensitive to the shape and size of the grains as well as the ice composition and temperature and is therefore a diagnostic of the evolutionary state and thermal history of ices \\citep{Sandford1988ApJ...329..498S}. As soon as the grain is warmed to $\\simeq$~12--15~K by the luminosity of the object, CO molecules can diffuse into the ice and form new bonds, changing the morphology of the mixture and thus the profile at 4.67 $\\mu$m, or they can sublime back to the gas phase (\\citealt{Al-Halabi2004A&A...422..777A}; \\citealt{Collings2003ApJ...583.1058C}; \\citealt{Givan1997JPhysChem}). The mobility of CO and its high abundance in cold icy mantles also explain why it is a key species for surface reactions leading to polyatomic molecules such as CO$_{\\rm 2}$ and CH$_{\\rm 3}$OH (\\citealt{Rodgers2003ApJ...585..355R}; \\citealt{Chiar1998ApJ...498..716C}; \\citealt{Teixeira1998A&A...330..711T}).  The VLT-ISAAC observations are used to constrain the physical and chemical conditions in the envelopes of a sample of intermediate-mass stars.  The paper is organized as follows. We present the objects of our sample in Sect.~\\ref{vela:obj} and the observations and data reduction procedures in Sect.~\\ref{vela:obs}.  The results on the water ice band are described in Sect.~\\ref{vela:L_band}. Evidence for solid CH$_{\\rm 3}$OH is presented in Sect.~\\ref{vela:methanol}.  The CO-ice band is presented in Sect.~\\ref{vela:co}. A possible correlation between the abundance of CO embedded into a water ice matrix and the IRAS 12~$\\mu$m/ 25~$\\mu$m color ratio is discussed in Sect.~\\ref{vela:cloud}. Conclusion are provided in Sect.~\\ref{vela:conclusion}. These data complement the survey of CO and other species for a sample of $\\sim$40 low mass YSO's by \\cite{Pontoppidan2003A&A...404L..17P,Pontoppidan2003A&A...408..981P} obtained in the same programme.   \\begin{figure}[ht] \\centering \\resizebox{\\hsize}{!}{\\includegraphics[]{vela_fig1.eps}} \\caption[Position of the observed sources]{CO (1--0) contour map of the Vela Molecular Cloud complex \\citep{Murphy1991A&A...247..202M}. The complex is composed of 4 clouds named A, B, C and D in the map. The positions of the observed sources are indicated and the filled circles are the location of the 50 other embedded young stellar objects.  Indicated as well are the locations of the Vela Pulsar and the H{\\sc II} region 259,-4. The map is an adaptation of Fig.1 of \\cite{Liseau1992A&A...265..577L}.} \\label{vela:fig_vela_cloud} \\end{figure} ",
        "conclusions": "\\label{vela:conclusion}  We have obtained $L$- and $M$- band spectra of a sample of intermediate-mass protostars in the Vela molecular cloud with the VLT-ISAAC. This is the first significant sample of intermediate mass protostars for which ice data are published.  A broad absorption feature at $\\sim$3.01 $\\mu$m is detected in all sources (main and companion objects). The features show an extended wing beyond 3.25 $\\mu$m, which can be reproduced in part by scattering by grains at radius 0.4--0.5 $\\mu$m. The water ice feature is dominated by absorption from cold amorphous ice although the spectroscopic signature of warm water ice can be masked if the ice is porous.  Methanol ice is only detected around the protostar \\object{LLN 17} (\\object{IRAS 08448--4343}). The derived abundance is 10 $\\pm$ 2~\\% relative to water ice. The upper limit on the methanol abundance toward the other sources is between 5 and 10\\% with respect to water ice.  Solid CO is detected in four main objects and one companion object. The profiles show a large variety of shapes. A strong variation of the total CO ice column density is found. We decompose the CO ice feature into three components. The column of CO becomes significant (i.e. larger than 10$^{17}$ cm$^{-2}$) only at $R$(12/25) greater than $\\sim$0.3. The color $R$(12/25) may trace the abundance of water ice with respect to silicate. There is no clear trend between the column density of pure CO ice, traced by the middle component, and other characteristics of the YSOs ($L_{\\mathrm{bol}}$, A$_{\\mathrm{V}}$, IRAS bands flux ratios). On the other hand, we find a possible correlation between the ratio of the flux at 12~$\\mu$m and 25~$\\mu$m, $R$(12/25), which is a measure of the warm dust temperature (100$<$ $T$ $<$ 250~K), and the amount of CO ice trapped in a water rich ice mantle, traced by the red component.  Likewise, the amount of CO ice trapped in a water and that of water may correlate.  This possible correlation is consistent with the idea that the water ice and the CO embedded in it sublime simultaneously.  These possible correlations should be tested in other star-forming regions. However, CO ice trapped in water-rich ice cannot solely account for the large amount of CO in the red component seen toward YSOs in Vela. Other factors such as scattering by the larger grains in the size distribution can probably contribute to the red component.  A strong absorption feature centered at 4.62 $\\mu$m is detected toward \\object{LLN 17} (\\object{IRAS 08448--4343}). The feature is likely caused by OCN$^-$.  The derived abundance relative to water is 1~$\\pm$~0.2~\\%. This feature is not detected in any other object in our sample. Together with the detection of methanol and the broad CO feature, the detection of OCN$^{-}$ suggests that the ice has been thermally and/or UV processed in \\object{LLN 17}. The processing generated by UV from the central object is not essential and perhaps shocks induced processing is at play. Further theoretical investigations on the possibility to form large amounts of methanol ice together with water ice in post-shocked regions are needed.  Observations of a larger sample of high signal-to-noise ratio spectra obtained with 8-10 meter class telescopes (VLT, Gemini, Keck) and with Spitzer at longer wavelengths of protostars of varying luminosities combined with sophisticated laboratory experiments will improve our understanding of the nature of the ices and their role in the synthesis of complex molecules in the interstellar medium."
    },
    "0512/astro-ph0512269_arXiv.txt": {
        "abstract": "Evidence of small-scale condensations in the magnetic molecular clouds has been accumulating over the past decades through radio and optical/ultraviolet observations. The origin and shape of these small-scale condensations is a disputable issue. Nejad-Asghar \\& Ghanbari (2004 hereafter NG) have recently studied the effect of the linear thermal instability on the formation of fluctuations in molecular clouds. The authors inferred that under certain conditions (e.g., depending on expansion or contraction of the background) thermal instability and ambipolar diffusion can produce spherical, oblate, or prolate condensations. ",
        "introduction": "Firstly, about presence of these fluctuations. We have two method to find them: (1) direct imaging of nearby clouds reveals substructures on scales to lengths of $0.01pc$ and masses of $0.01M_\\odot$ \\cite{peng98,saka03}, (2) studies of the time variability of absorption lines indicates the presence of fluctuations on scales of $10^{-4} pc$ (5-50 AU) and masses of $10^{-9} M_\\odot$ \\cite{moor95,bois05}.  How these fluctuations may be formed and is the thermal instability important? To answer the question, first, we must consider time-scales. According to the prior results of Gilden \\shortcite{gild84}, thermal instability time-scale in molecular clouds is in the order of $10^3-10^4$ year. On the other hand, Larson \\shortcite{lars81} showed that dynamical time-scale for transient structure of molecular clouds is in order of $10^7$ year. Thermal instability time-scale is very smaller than dynamical time-scale. Thus, it is reasonable to consider thermal instability as an important mechanism in formation of small fluctuations in molecular clouds.  It is now accepted that the magnetic field is very important in dynamical evolution of all ISM. It will affect on ions directly and on neutral particles indirectly, via collision with ions. In molecular clouds, ionization degree is very small, thus, drift of neutral particles between tied ions is important. It can heat the medium. In this way, for complete investigation of the thermal instability in molecular clouds, we must consider plasma drift (ambipolar diffusion). ",
        "conclusions": ""
    },
    "0512/astro-ph0512096_arXiv.txt": {
        "abstract": "Velocity oscillations in sunspot umbrae have been measured simultaneously in two spectral lines: the photospheric Silicon {\\sc i}  \\hbox{10827 \\AA} line and the chromospheric Helium {\\sc i} \\hbox{10830 \\AA} multiplet. From the full Stokes inversion of temporal series of spectropolarimetric observations we retrieved, among other parameters, the line of sight velocity temporal variations at photospheric and chromospheric heights. Chromospheric velocity oscillations show a three minute period with a clear sawtooth shape typical of propagating shock wave fronts. Photospheric velocity oscillations have basically a five minute period, although the power spectrum also shows a secondary peak in the three minute band which has proven to be predecessor for its chromospheric counterpart. The derived phase spectra yield a value of the atmospheric cut-off frequency around $4$ mHz and give evidence for the upward propagation of higher frequency oscillation modes. The phase spectrum has been reproduced with a simple model of linear vertical propagation of slow magneto-acoustic waves in a stratified magnetized atmosphere that accounts for radiative losses through Newton's cooling law. The model explains the main features in the phase spectrum, and allows us to compute the theoretical time delay between the photospheric and chromospheric signals, which happens to have a strong dependence on frequency. We find a very good agreement between this and the time delay obtained directly from the cross-correlation of photospheric and chromospheric velocity maps filtered around the 6 mHz band. This allows us to infer that the 3-minute power observed at chromospheric heights comes directly from the photosphere by means of linear wave propagation, rather than from non-linear interaction of 5-minute (and/or higher frequency) modes. ",
        "introduction": "\\noindent The study of the generation and propagation of waves in the solar atmosphere is a hot topic of research in astrophysics, since it provides information about the atmospheric structure and dynamics (e.g., Lites 1992; Bogdan 2000; Socas-Navarro, Trujillo Bueno and Ruiz Cobo, 2000; Bogdan and Judge 2006), while at the same time it helps us identify the key mechanisms of chromospheric and coronal heating.  In fact, acoustic and magnetic waves and magnetic field reconnection have been mentioned in the literature as the most promising heating mechanism candidates (Alfv\\'en 1947; Biermann 1948; Schwarzschild 1948; Parker 1979; Ulmschneider and Musielak 2003).   Historically, sunspot oscillations have been classified into three different groups (e.g. Lites 1992): (1) photospheric umbral oscillations, which have basically a 5-minute period with an average rms amplitude of 75 m$\\rm{s^{-1}}$. These oscillations are the apparent response of the umbral photosphere to the 5-minute p-mode oscillations. (2) Chromospheric umbral oscillations with periods around 180 s and amplitudes of a few kilometers per second, and (3) running penumbral waves, seen in $\\rm{H_{\\alpha}}$ as disturbances propagating radially outwards from the umbra. They all seem to be different manifestations of the same dynamical global phenomenon, though (e.g., Rouppe van der Voort et al. 2003).  \\noindent Simultaneous time-series observations of various spectral lines that sample different regions of the solar atmosphere is one of the most useful techniques for studying wave propagation (e.g., the review by Lites 1992 and references therein). For instance, Lites (1986) could provide hints of shock wave formation via the Doppler shifts observed in the Stokes $I$ profiles of the He {\\sc i} 10830 \\AA\\ multiplet. By measuring the phase difference of the oscillations in different spectral lines, this author could also infer the upward propagation of waves in the frequency band around 6.5 mHz (Lites, 1984). Other pioneering investigations on this topic are those by Kneer et al. (1981).  \\noindent In the last 35 years, since the first report on chromospheric umbral oscillations was made (Beckers and Tallant, 1969), many works have been published on this subject, accompanied by nearly an equal number of differing findings, conclusions and contradictions yielded by the literature in this time. We refer the reader to recent reviews (such as those by Bogdan 2000 and Bogdan \\& Judge 2006) for a comprehensive overview of present knowledge of oscillatory phenomena in sunspots, both from the theoretical and the observational points of view.   Nowadays, theoretical investigations on this topic are mainly carried out by means of detailed numerical simulations. For instance, the hydrodynamical simulations of Carlsson \\& Stein (1995) suggest that acoustic shock waves in the internetwork regions of the solar atmosphere intermittantly heat the plasma there, but with an acoustic shock heating that is insufficient to explain quantitatively the emission line cores observed in far-UV lines. Similar numerical simulations have recently begun to be extended to strongly magnetized regions of the solar atmosphere, taking into account the coupling among various MHD wave modes (e.g., Stein et al. 2004), but much work remains to be done prior to reaching a level of realism for which it becomes reasonable to start contrasting computed Stokes profiles with spectropolarimetric observations. In this respect, one of the aims of this paper is to provide high-quality observational information on the phenomenon of oscillations in sunspot umbrae, based on full Stokes-vector IR spectropolarimetry in photospheric and chromospheric lines.   This paper is organized as follows: Observations, data redution and inversion techniques are presented in sections 2 and 3. For the analysis, we follow a similar approach to that of Lites (1984, 1986) but measuring instead the full Stokes-vector of the photospheric Silicon {\\sc i}  \\hbox{10827 \\AA} line and of the chromospheric Helium {\\sc i} \\hbox{10830 \\AA} multiplet. The analysis of the photospheric and chromospheric LOS velocities, and the relation between them, are shown in section 4. As we shall see below, we are able to provide very clear observational evidence for the upward propagation of waves from the photosphere to the chromosphere within the umbra of a sunspot, including an unprecedent measurement of the time delay between the signals and the detection and characterization of the photospheric driving piston. A brief discussion can be found in section 5, followed by some final remarks in section 6. ",
        "conclusions": "The spectropolarimetric investigation we have presented here provides observational evidence for the upward propagation of slow magneto-acoustic waves from the photosphere to the high chromosphere inside the umbra of a sunspot. The time delay between the signals corresponding to both regions varies strongly with the frequency of the oscillation, going from a few tens of seconds to several minutes. As the photospheric perturbations propagate upwards, their amplitude increases due to the rapid decrease in density, and they eventually develop into shock waves at chromospheric heights.  Interestingly, the observed temporal variability of the Stokes profiles in the Si {\\sc i} line at 10827.09 \\AA\\ may help to establish the required initial condition for performing realistic MHD simulations, which are needed for a full physical understanding of the phenomenon of wave propagation in sunspot atmospheres. Our future work on this topic will focus on similar spectropolarimetric investigations, but for atmospheric plasma structures with lower manetic fluxes, such as pores, active region plages and the chromospheric network of the `quiet' Sun.   It is a pleasure to thank H\\'ector Socas-Navarro and Jos\\'e Antonio Bonet for fruitful discussions and for sharing with us their knowledge on Stokes inversion techniques and Fourier analysis. We are also grateful to Bruce Lites and Tom Bogdan for suggesting improvements to an earlier version of this paper. This research has been funded by the Spanish Ministerio de Educaci\\'on y Ciencia through the proyect AYA2004-05792, and is part of the European Solar Magnetism Network."
    },
    "0512/astro-ph0512575_arXiv.txt": {
        "abstract": "We present the results of a systematic analysis of gamma-ray burst afterglow spectral energy distributions (SEDs) in the optical/near-infrared bands.  Our input list includes the entire world sample of afterglows observed in the pre-\\emph{Swift} era by the end of 2004 that have sufficient publicly available data. We apply various dust extinction models to fit the observed SEDs (Milky Way, Large Magellanic Cloud and Small Magellanic Cloud) and derive the corresponding intrinsic extinction in the GRB host galaxies and the intrinsic spectral slopes of the afterglows. We then use these results to explore the parameter space of the power-law index of the electron distribution function and to derive the absolute magnitudes of the unextinguished afterglows. ",
        "introduction": "According to the most popular progenitor model of long-duration Gamma-Ray Bursts (GRBs), the collapsar model \\citep{Woosley1993}, a GRB is the result of ultra-relativistic jets ejected by an accreting black hole formed by the core-collapse of a massive star (most probably a Wolf-Rayet star). This predicts a physical link between GRBs and supernovae (SNe) that has been spectroscopically confirmed in four cases so far: XRF 020903 \\citep{Soderberg2005}, GRB 021211/SN2002lt \\citep{DellaValle2003}, GRB 030329/SN2003dh \\citep[e.g.,][]{Hjorth030329, Stanek030329} and GRB 031203/SN2003lw \\citep[e.g.,][Mazzali et al. 2005, in preparation] {Malesani2004}. Further evidence comes from the observation of late-time bumps in GRB afterglows that can be modelled very well by an underlying SN component \\citep{Bloom1999, ZKH}, and which  have led to the conclusion that in fact all long-duration GRBs are  physically related to SN explosions \\citep[][Paper I]{ZKH}.  Furthermore, the collapsar model implies that the progenitors of long-duration GRBs are associated with regions of high-mass star formation \\citep{Paczynski1998}, which might reveal themselves by a detectable extinction in the GRB host galaxies along the lines of sight towards the burster. This idea is further supported by the so-called dark bursts \\citep{Groot1998}, for which no optical afterglow has been discovered despite rapid and deep searches in small error box regions. While most none-discoveries are the result of too shallow searches and too large error boxes \\citep[e.g.,][]{Jakobsson2004b, Rol2005}, a small percentile remains that require intrinsic extinction to dim the afterglow, e.g., GRB 970828 \\citep{Groot1998, Djorgovski2001}, 990506 \\citep{Taylor2000} and 020819 \\citep{JakobssonIPR}, while others might have been intrinsically underluminous \\citep[for a discussion, see e.g.,][]{Fynbo2001b, Lazzati2002, Klose2003}.  Of particular interest within the dust extinction model is the  statistical distribution of the amount of visual extinction in the GRB host galaxies, as well as the nature of the corresponding dust. Given the fact that GRBs and their afterglows can be observed up to high redshift, they offer the possibility to get insight into the nature of cosmic dust when the universe was much younger and galaxies were much less evolved. Furthermore, since the optical properties of the dust grains trace the  environmental conditions in the interstellar medium  \\citep{FitzpatrickMassa1986, Draine2000, Bradley2005}, they are to some degree an indicator for the physical conditions to make a GRB progenitor.  The present paper is the third in a series of papers where we employ a large database of photometry gathered from the literature to reanalyze all optical afterglow light curves of GRBs in the pre-\\emph{Swift} era in a consistent manner to derive a homogenous sample, which is then used to study afterglow properties in a statistical way. In Paper I and an update \\citep{ZKH05}, the properties of the supernovae underlying nearby afterglow light curves were explored. In Paper II \\citep{ZKK} the optical light curve parameters were derived for a large (and basically complete) sample of afterglows observed by the end of 2004, up to the launch of the \\emph{Swift} mission. In this paper, we extend this systematic analysis to the spectral energy distribution (SED) of GRB afterglows  in the optical/near-infrared bands in order to search for signals from extinction by dust in the GRB host galaxies.  In \\kref{data}, we present the methods with which we analyzed the afterglow SEDs. Among the 59 GRBs studied in Paper II, 30 had data quality sufficiently high to be included in a sample for an investigation of the SEDs. We then further reduce this sample to a Golden Sample of 19 GRB afterglows with parameters derived from specific dust model fits. In \\kref{RaD} we present the results derived from our analysis and discuss our findings in the  context of the standard fireball model. The fitting process allows us to derive the host galaxy extinction $\\AV$ along the line of sight and the intrinsic, extinction-corrected spectral slope $\\beta$ ($F_\\nu\\propto\\nu^{-\\beta}$) of the afterglow light in the optical/near-infrared bands. The $\\AV$-$\\beta$ sample is then further analyzed statistically to derive conclusions about the environment of GRBs and the dust properties of high-redshift galaxies. ",
        "conclusions": "\\label{conclusions}  We have presented a sample of 30 GRB afterglow spectral energy distributions in the optical/NIR bands which have been modeled with various dust extinction curves (Milky Way, Large Magellanic Cloud and Small Magellanic Cloud) to derive the source frame extinction, $\\AV$, intrinsic to the host galaxies and the spectral slope, $\\beta$, of the afterglows unaffected by any dust extinction. As all afterglows have been analyzed in a systematic way, the results are fully comparable, making this sample unique in terms of both size and consistency. For the further statistical study, we selected 19 afterglows, our Golden Sample, which have physically reasonable results and small error bars.  The preferred dust models we find (\\kref{SMCchapter}) as well as the deduced source frame dust-to-gas ratios (\\kref{NH}) based on the inclusion of data taken from the literature, both indicate that the majority of GRBs we have investigated, covering the redshift range from 0.1 to 4.5, occur in low-metallicity environments. The $\\AV$ distribution that we have derived from these data (\\kref{AVchapter}) highlights a sparsity of strongly extinguished afterglows, creating a Dark Burst Desert, even though it is unclear if the preference of low extinctions is more than an observational and sample selection bias. Our finding that most afterglows suffer from only low extinction in their hosts could indicate that afterglows are usually not obscured by dust close to the burster. One would then expect that the most extinguished afterglows are in fact located in globally dusty hosts. Indeed, we find weak evidence for  a correlation between the submm flux of GRB host galaxies and  the source frame extinction $\\AV$. Although the statistical significance is low due to the small sample size and the large errors, this finding calls for a more thorough investigation.  Knowledge of $\\beta$ and $\\AV$ allowed us then to correct the afterglow light curves for intrinsic extinction and to derive the true luminosity distribution of our afterglow sample at chosen times in the host galaxy frame (\\kref{redshifting}). We find that, on average, low-$z$ afterglows are less luminous than high-$z$ afterglows. The most likely explanation we have at hand for this finding is an observational bias against intrinsically faint afterglows at high redshifts. A bimodal distribution found by \\cite{GB2005} in similarly corrected X-ray afterglows is not clearly seen in the optical although, on average, GRBs with fainter X-ray afterglows also have fainter optical afterglows. Unfortunately, the available sample size is still too small to reach definite conclusions. A search for correlations between prompt emission parameters and the luminosity of the optical afterglows at jet break time has come up empty.  Since our sample is exclusively composed of GRBs from the pre-\\emph{Swift} era, a similar study in a few years time on \\emph{Swift}-discovered GRB afterglows will shed light on the Dark Burst Desert and the true afterglow luminosity distribution by removing observational bias factors via rapid and highly precise GRB localizations. Already, \\emph{Swift} has lead to the discovery of very faint afterglows \\citep[e.g. GRB 050126, GRB 050607,] []{BergerSwift2005, Rhoads2005} including what may be the \"darkest\" burst ever, GRB 050412 \\citep{KosugiGCN, JakobssonIPR}. The recent discovery of the first afterglows of short GRBs \\citep{HjorthShort, FoxShort, CovinoShort, BergerShort, SB051221} opens the possibility of finally making a comparison of the environment of the two different classes of GRBs."
    },
    "0512/physics0512062_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\lvm  This talk is about the possibility that the Solar System belongs to the territory of a hypercivilization spanning our galaxy or a large region of it. I will start introducing the Fermi Paradox (why we do not see aliens around?) and some of its solutions. Then I will present my own solution which includes two proposals called the Subanthropic Principle and the Undetectability Conjecture. This solution states that, at present, all typical galaxies like ours are already colonized by very advanced technological civilizations spread through large regions or the whole galaxies, many of them containing primitive subcivilizations like ours. After discussing some consequences of this solution for our planet and our civilization I will make some comments on recent, very popular theories in the scientific community of Particle Physics and Cosmology. These theories, known as `brane worlds', assume that our visible Universe with three space dimensions is embedded in a much larger Cosmos with more space dimensions. Therefore it would be most natural if other universes would also exist located along the extra space dimensions. As a result, these theories open up enormous possibilities regarding the visitation or colonization of the Solar System by alien civilizations, strengthening the Fermi Paradox. Finally, in the appendix I have included some questions and answers that came up during this Forum. ",
        "conclusions": ""
    },
    "0512/astro-ph0512433_arXiv.txt": {
        "abstract": "We generalize the predictions for the CMB anisotropy patterns arising in Bianchi type VII$_h$ universes to include a dark energy component. We consider these models in light of the result of \\cite{jaffe:2005a,jaffe:2005b} in which a correlation was found on large angular scales between the \\emph{WMAP} data and the anisotropy structure in a low density Bianchi universe. We find that by including a term $\\Omega_\\Lambda > 0$, the same best-fit anisotropy pattern is reproduced by several combinations of cosmological parameters. This sub-set of models can then be further constrained by current observations that limit the values of various cosmological parameters. In particular, we consider the so-called geometric degeneracy in these parameters imposed by the peak structure of the \\emph{WMAP} data itself. Apparently, despite the additional freedom allowed by the dark energy component, the modified Bianchi models are ruled out at high significance. ",
        "introduction": "While cosmology appears to be converging on a ``concordance model'' that describes the universe as inflationary and isotropic, there remain unexplained anomalies in the CMB data, and other models are not yet ruled out.  The \\emph{WMAP} data provide some of the most accurate measurements yet of the cosmic microwave background and contribute to high accuracy determinations of cosmological parameters \\citep{bennett:2003,spergel:2003}. However, there are several studies that show that at large scales, the CMB is not statistically isotropic and Gaussian, as predicted by inflation theory \\citep{de Oliveira-Costa:2004, eriksen:2004, hansen:2004b, vielva:2004}.  In \\cite{jaffe:2005a,jaffe:2005b}, we examined a particular set of anisotropic cosmological models: the Bianchi type VII$_h$ class of spatially homogeneous generalizations of Friedmann universes that include small vorticity (universal rotation) and shear (differential expansion) components.  Surprisingly, we found evidence that one of these models correlates with the CMB sky and that subtracting this component resolves several of these observed anomalies.  The models used in that study were derived by \\cite{BJS} and include no cosmological constant component in the total energy density.  The best-fit model found by that study required $\\Omega_{\\textrm{m}0}=0.5$, implying, $\\Omega_k=0.5$, \\ie a significantly negatively curved universe.  \\citet{land:2005} subsequently considered flat Bianchi models and sought explicitly to resolve the problems of the low quadrupole and the low-$\\ell$ alignments using a statistic constructed to achieve that purpose. Their analysis does indeed find a model that fixes these anomalies, % but it remains statistically insignificant as a detection. Our original result has the benefit of being a detection that is based entirely on a simple least-squares fit of the Bianchi template to the data, yet we find that it serendipitously resolves several anomalies % without that requirement having been built-in to the search algorithm. The importance of our result, therefore, lies in the fact that it resolves the problems of the low quadrupole, low-$\\ell$ alignments, power asymmetry, and non-Gaussian cold spot.  However, both analyses neglect the fact that the existing Bianchi solutions include no dark energy component. Furthermore, the best-fit results require an energy content that is consistent neither with other astronomical observations nor with the CMB itself. In this work, we present a modification of the Barrow \\etal Bianchi type VII$_h$ solution so that it includes a cosmological constant term in the evolution.  We discuss the impact of that term on the structure of the resulting CMB anisotropy pattern, particularly the degeneracy that is introduced in the model space by the addition of $\\Omega_\\Lambda$.  We then discuss the viability of $\\Lambda$ models that are morphologically identical to our best-fit $\\Omega_\\Lambda=0$ models by considering the constraints imposed by different measurements of the cosmological parameters. ",
        "conclusions": "We have presented solutions for Bianchi type VII$_h$ type universes that include a dark energy term and examined how their morphological properties change over the parameter space.  The addition of dark energy adds a degeneracy such that different combinations of the three parameters $(\\Omega_{\\textrm{m}},\\Omega_\\Lambda, x)$ can lead to the same observed structure as in the best-fit model of \\cite{jaffe:2005a,jaffe:2005b}.  A template can be constructed that has the identical structure of that best-fit model and also falls on the geometric degeneracy curve for the parameters as measured by \\emph{WMAP} data.  This model, however, lies well outside the 95\\% confidence region in $\\Omega_{\\textrm{m}}-\\Omega_\\Lambda$ space for \\emph{WMAP} data, ruled out by the over-prediction of large scale power.  It also lies outside the 95\\% likelihood contours from the Supernova Cosmology Project, and is further inconsistent with the measurement of the Hubble constant from the Hubble Key Project.  Bianchi models that are more consistent with these other measurements are no longer good fits to the \\emph{WMAP} large scale structure.  One of the most difficult problems for these models is to account for the acoustic peak structure.  The anisotropy at early times might influence the nature of the fluctuations at last scattering, and the geometry could affect the power spectrum on small angular scales due to the geodesic focusing between last scattering and the observer. Detailed predictions for these effects would be needed, but it is difficult to envision such an anisotropic scenario that happened to reproduce the observed acoustic peak structure, mimicking the concordance cosmology so well.  There is currently no prediction for the CMB polarization anisotropy in a Bianchi universe, but such a geometry-induced signal could provide an additional test of these models.  If the preferred direction indicated in the temperature data is also reflected in the full sky polarization data (expected from further \\emph{WMAP} data releases and, eventually, from Planck), there will be even more motivation to consider non-standard models.   We have shown that our best-fit Bianchi type VII$_h$ model is not compatible with measured cosmological parameters, despite the additional freedom from adding dark energy.  It is worth reiterating, however, that the serendipitous discovery of a theoretically derived template that correlates well with the data also happens to resolve several anomalies that cannot be explained in the standard picture. These particular models may not be viable, but lacking any plausible scenario for systematics or foregrounds to be the source of the anomalies, non-standard models that reproduce a similar morphology merit continued interest."
    },
    "0512/astro-ph0512119_arXiv.txt": {
        "abstract": "We analyze a portion of the SDSS photometric catalog, consisting of approximately 10,000 objects that have been spectroscopically classified into stars, galaxies, QSOs, late-type stars and unknown objects (spectroscopically unclassified objects, SUOs), in order to investigate the existence and nature of subclasses of the unclassified objects.  We use a modified mixture modeling approach that makes use of both labeled and unlabeled data and performs class discovery on the data set.  The modeling was done using four colors derived from the SDSS photometry: $(u-g)$, $(g-r)$, $(r-i)$, and $(i-z)$.  This technique discovers putative novel classes by identifying compact clusters that largely contain objects from the spectroscopically unclassified class of objects.  These clusters are of possible scientific interest because they represent structured groups of outliers, relative to the known object classes.  We identify two such well defined subclasses of the SUOs.  One subclass contains 58\\% SUOs, 40\\% stars, and 2\\% galaxies, QSOs, and late-type stars.  The other contains 91\\% SUOs, 6\\% late-type stars, and 3\\% stars, galaxies, and QSOs.  We discuss possible interpretations of these subclasses while also noting some caution must be applied to purely color-based object classifications. As a side benefit of this limited study we also find two distinct classes, consisting largely of galaxies, that coincide with the recently discussed bimodal galaxy color distribution. ",
        "introduction": "One of the main goals of modern photometric and spectroscopic surveys is to understand and isolate the different types of measured objects. This is best done automatically, both due to the overwhelming number of objects in the surveys and because of the relative objectivity gained by using automated methods.  There are a number of approaches to classification of objects, traditionally falling into two major groups: supervised and unsupervised.  Supervised classifiers make use of labeled training data in order to train the software to recognize certain patterns in the data features that are characteristic of the different object classes known to be present.  Unsupervised classifiers, or clustering algorithms, try to find natural groupings of objects in feature space.  \\citet{Storrie92} were the first to use neural networks, a form of supervised classification, to automatically classify galaxies according to their morphology.  \\citet{Odewahn92} have developed successful neural network based methods for star/galaxy separation.  Since the mid-1990s there has been a lot of work using neural networks for stellar object classification.  \\citet{Naim95a, Naim95b} used backpropagation neural networks to classify galaxies based on their morphology and showed the results compared favorably to classification by humans.  \\citet{Vonhipple94} and \\citet{Bailer-Jones97, Bailer-Jones98} used neural networks to successfully predict effective temperatures, MK classifications, and luminosity classes of stars based on stellar spectra.  \\citet{Ball04} used a neural network and a variety of morphological and photometric parameters to predict the eClass of SDSS galaxy spectra.  The eClass is a continuous parameterization of the galaxy type derived from principal component analysis applied to a collection of SDSS galaxy spectra \\citep{Connolly95, Connolly99}.  \\citet{Willemsen05} determined the metallicities of stars from two different globular clusters using neural networks.  Another supervised classification method that has been used with astronomical data is the decision tree inducer.  \\citet{Weir95a} used a decision tree to classify stars and galaxies from the Palomar Sky Survey (DPOSS) based on a set of eight features.  \\citet{Owens96} used an oblique decision tree classifier (OC1) to study the same data as \\citet{Storrie92} used with a neural network, achieving comparable results.  \\citet{Jarrett00} describe the use of OC1 for classifying 2MASS extended sources.  \\citet{Bazell01} compared ensembles of decision trees with ensembles of other types of classifiers.  They found the largest reduction in classification error was for ensembles of decision trees.  More recently, \\citet{Suchkov05} used an ensemble of decision tree classifiers to determine types and redshifts of objects in the SDSS photometric catalog.  The Kohonen Self-Organizing Map (SOM) \\citep{Kohonen01} has been used as an unsupervised method of exploring associations within astronomical data sets \\citep{Mahonen95, Miller96, Naim97a, Rajaniemi02}.  This approach, an unsupervised neural network, clusters similar objects together in a way that conserves the topology of the input data.  In other words, input vectors that are similar to each other are mapped to neighboring regions of what is usually a two-dimensional output lattice.  In this way SOMs have been found useful as a way to visualize high dimensional data.  A common approach to unsupervised clustering is the use of mixture models \\citep{Duda, Mclachlan, Raftery}. Mixture models produce probabilistic, or soft, assignments of data points to each of the mixture components, or clusters.  The number of clusters to be learned must be specified as part of the algorithm.  There is no single standard approach for choosing this parameter.  Furthermore, without supervising examples, there is no guarantee that the learned clusters will acceptably characterize the known classes within the data set.  To overcome some of these shortcomings {\\it semisupervised learning} algorithms were proposed, including the work of \\citet{Shashahani}, \\citet{Miller97}, and \\citet{Nigam}.  In these approaches, the input data consists of both labeled and unlabeled samples.  This affects the clustering process in two ways.  First the labeled data helps guide the cluster definitions to properly reflect the known ground-truth classes \\citep{Miller97, Basu1}.  Second, the unlabeled data may help to more accurately learn parameters of the model, so as to better define the shapes of the learned classes in feature space \\citep{Shashahani, Miller97}.  Until recently, most semisupervised methods focused on building models and classifiers for the set of known classes, i.e. those for which labeled training examples are available.  However, given a data set with many unlabeled examples, it is quite possible that, latent in the unlabeled data, some {\\it unknown} classes are present.  These unknown classes are compact clusters of unlabeled objects that have not, as yet, been recognized by domain experts as distinct classes or categories of interest.  Such clusters, once identified, could be the focus of further scientific inquiry, to either validate them as new classes of interest or to reject them as uninteresting outlier groups. Recent work in \\citet{Miller03a, Miller03b} developed semisupervised mixture modeling methods with a built-in capability to discover these unknown classes in mixed labeled and unlabeled data.  Standard supervised learning algorithms do not have good means to identify unknown classes because they can only be trained using examples from the known classes.  Standard unsupervised clustering algorithms make no distinction between clusters of purely unlabeled data and clusters containing some labeled samples.  Purely unlabeled clusters are of particular scientific interest because they represent groups of objects that are well separated from objects belonging to known classes.  Thus, they may represent new object types or groups of unusual objects, i.e, outliers from existing known classes.  In a previous study, \\cite{Bazell05} applied the semisupervised class discovery approach to investigate astronomical data.  In this paper, we analyze a portion of the SDSS photometric catalog consisting of approximately 10,000 objects.  The SDSS data in our sample have been classified via the spectral pipeline into several different types of objects.  Here we particularly focus on the spectroscopically unclassified objects, SUOs, which consist of objects that were not readily classified as stars, galaxies, QSOs, or late type stars (red stars of type M or later).  The papers describing the SDSS data releases, in particular, \\citet{Stoughton}, as well as the SDSS website \\footnote{http://www.sdss.org} describe the procedure that leads to objects being labeled as spectroscopically unclassified (the SDSS class called ``unknown'').  SUOs are objects that passed through the various filters in the spectroscopic pipeline and were then examined by a person who could not reliably classify them because they were too noisy, too featureless, or for some other reason.  We are interested in applying our semisupervised class discovery technique to the SDSS data for two main reasons: 1) to see whether this approach determines significant substructure within the class of SUOs, i.e., subclasses of SUOs; 2) to determine whether some SUOs are clustered into groups that mainly consist of known class objects.  An alternative approach would be to apply a simple clustering algorithm, such as K-means, to identify substructure solely using as data the set of SUO samples.  However, both the validity of the learned SUO clusters and the nature of these clusters are best assessed relative to the known class clusters, by learning a clustering solution using the data from all the classes, both known and spectroscopically unclassified.  Our method clusters on the basis of {\\it all} available information -- feature vectors, available class labels, and the {\\it fact} of label presence/absence for each sample \\citep{Miller03a}.  Absence of labels in a compact cluster is suggestive that a meaningful subclass of the unknown class may have been found. \\citet{Miller03a} demonstrated that use of label presence/absence information may help to achieve more accurate clusters and to better discern unknown subclasses than methods which do not use this information.  Our semisupervised mixture modeling approach has several benefits not seen by using standard unsupervised clustering.  In particular, some objects labeled SUOs may have measured features that are similar to those of objects from known classes.  In performing clustering using all the data, both from the known classes as well as the SUO class, we can learn clustering solutions that reveal these similarities.  That is, if an SUO is assigned to a cluster of predominantly ``known'' objects of a given type, the clustering is indicating the possibility that the SUO may be related to (e.g. may be a variant of) the known object class.  This may also suggest the SUO was mislabeled. Likewise, consider a cluster that primarily contains SUOs, but which also owns some known class objects.  The known class composition of the cluster may hint at the underlying physical nature of the SUO subclass represented by this cluster.  On the other hand, the ownership of some known class objects by the cluster may indicate that these known class objects were mislabeled.  Finally, consider the problem of choosing the {\\it number} of clusters to represent each class.  Model order selection methods ``match'' the model complexity (number of clusters in each class) to the amount of data these clusters own.  If a significant number of SUO class objects can be explained by, i.e. belong to, clusters of known class objects, then fewer SUO clusters will be needed in the model.  This means that learning and model selection using only the SUO class objects could lead to an overestimate of the number of SUO clusters (subclasses).  There has been limited work using mixture models to classify astronomical data.  \\citet{Strateva01} used a mixture model to verify the bimodal galaxy distribution they found by other means. \\citet{Nichol01} discuss the use of mixture models for finding clusters of galaxies.  \\citet{Yip04} used a simple mixture model to identify classes of objects following principal component analysis (PCA) of spectra.  They applied a mixture model to the histogram of their parameter $\\phi_{KL}$, the angle between the first two eigencoefficients of the PCA decomposition of the galaxy spectra. They performed model selection using the Akaike Information Criterion \\citep{Akaike} and found three components best fit their data. \\citet{Kelly04, Kelly05} used a similar approach following PCA of the shapelet decomposition of galaxy images and used the Bayesian Information Criterion (BIC, \\citet{Schwarz}).  They also found that a three class representation of their data was optimal, although their data set was completely different from that of \\citet{Yip04}. ",
        "conclusions": "The plot of the BIC cost vs. number of components, Figure 1, shows a distinct global minimum, but there is also a lot of scatter evident. We performed additional runs using models with between 10 and 30 components in an attempt to quantify the reproducibility of the number of nonpredefined components in each model. Table \\ref{tab:multirun} displays some results that illustrate this issue.  \\begin{deluxetable}{rrrrr} \\tabletypesize{\\scriptsize} \\tablecaption{Number of Nonpredefined Components Resulting from Repeated Algorithm Runs \\label{tab:multirun}} \\tablewidth{0pt} \\tablehead{ \\colhead{$N_{comp}$} & \\colhead{$N_{np}$} & \\colhead{Avg} & \\colhead{$\\sigma$}} \\startdata 10 & 1 2 2 3 2 & 2.0 & 0.63 \\\\ 11 & 2 2 1 2 2 & 1.8 & 0.40 \\\\ 12 & 2 2 2 2 3 & 2.0 & 0.40 \\\\ 13 & 2 1 2 2 3 & 2.0 & 0.63 \\\\ 14 & 2 3 1 2 2 & 2.0 & 0.63 \\\\ 15 & 2 3 2 1 5 & 2.8 & 1.47 \\\\ 16 & 3 3 3 4 2 & 3.0 & 0.63 \\\\ 17 & 3 4 3 2 3 & 3.0 & 1.26 \\\\ 18 & 4 1 2 4 4 & 3.0 & 1.26 \\\\ 19 & 5 2 4 5 6 & 4.4 & 1.46 \\\\ 20 & 3 2 6 6 3 & 4.0 & 1.67 \\\\ 21 & 5 3 4 3 3 & 3.6 & 0.80 \\\\ 22 & 3 3 5 4 3 & 3.6 & 0.80 \\\\ 23 & 3 6 5 2 4 & 4.0 & 1.41 \\\\ 24 & 8 7 4 6 4 & 3.4 & 1.60 \\\\ 25 & 3 4 5 2 3 & 3.4 & 1.02 \\\\ 26 & 5 3 4 4 2 & 3.6 & 1.02 \\\\ 27 & 7 4 3 3 6 & 4.6 & 1.62 \\\\ 28 & 5 5 5 5 4 & 4.8 & 0.40 \\\\ 29 & 6 6 5 2 4 & 4.6 & 1.50 \\\\ 30 & 6 4 5 5 7 & 5.4 & 1.02 \\\\ \\enddata \\tablecomments{Column $N_{np}$ denotes the number of nonpredefined components for each of 5 runs with a total of $N_{comp}$ components.} \\end{deluxetable}  Column 1 gives the number of components in the model.  Column 2 gives the number of nonpredefined components for each of the five additional runs performed.  Column 3 gives the average number of nonpredefined components and Column 4 gives the standard deviation.  We see that there is a trend toward increasing number of nonpredefined components as the model complexity (total number of components) increases.  This is expected since {\\it all} classes, including the SUO class, should be represented by more components as the total number of components is increased.  Note that for $N_{comp} = 16$ the additional model runs produced an average of 3 nonpredefined components.  However, these models all had higher BIC cost than the one we chose for analysis, which was based on the model with the lowest BIC cost; this model had only two nonpredefined components.  For $N_{comp} = 14$, which also had a low BIC cost, the average number of nonpredefined components is 2, with little scatter.  For $N_{comp} = 24$, the other model with especially low BIC cost, the average number of nonpredefined components is 3.4, with three models producing six or more nonpredefined components. This is significantly different from the model we analyzed in detail and deserves further study.  We hope to examine this in more detail in a future paper.  The comparison of Table \\ref{tab:comp-membership} with Table \\ref{tab:comp-membership2} provides some insight into the different results that are achieved with the semisupervised vs. unsupervised algorithms.  The best unsupervised model had 15 components compared with the best semisupervised model, which contained 16 components. Note that the numbers identifying each component are arbitrary so we cannot perform a direct component to component comparison between the two models.  However, we do note the following.  Component 0 of the semisupervised model (SS0 for short) is very similar to component 1 of the unsupervised model (US1 for short); they captured only late-type stars, with 494 and 487 stars respectively.  Also, SS12 corresponds closely to US2 since both captured the bulk of the QSOs in the sample (80\\% and 74\\% respectively).  Still, the semisupervised algorithm did a better job at isolating the QSOs.  SS10 and SS11, the nonpredefined components, appear to have unsupervised analogs in US12 and US4.  SS10 and SS11 together capture 37\\% of the SUOs while US12 and US4 capture 36\\%.  According to the SDSS spectroscopic classification procedure, the objects that comprise component 10 are a mixture of mainly SUOs and stars.  It is notable in Figure 3 that many of the SUO class points are strongly clustered together between $-0.1<(u-g)<0.3$ and $-0.3<(g-r)<0$. This component overlaps the white dwarf exclusion region of \\citet{Richards02} in all three color-color projections they defined.  It also overlaps a region of low redshift, ugri selected quasars in $(u-g)$ vs. $(g-r)$ space (see figure 13 of \\citet{Richards02}; also \\citet{Hall02}). It also overlaps the high density region of confirmed quasars from \\citet{Richards04} in all three projections.  All the works just cited base their quasar identifications on photometric properties, while our identification of the objects as SUOs is based on their spectral properties being unusual in some way.  Indeed, a visual inspection of the spectra labeled as SUOs in component 10 shows that they are almost all blue, relatively featureless spectra.  Most of these are likely to be stars, sdO or DA white dwarfs \\citep{Kleinman04}, with a small number of BL LAC objects.  Figure 5 shows spectra from several representative objects in component 10. \\clearpage \\begin{figure} \\epsscale{0.8} \\plotone{f4.eps} \\label{fig:Spec10} \\caption{\\small Typical spectra from objects in component 10.  These are all blue, featureless spectra which are most likely sd0 or DA white dwarfs.  The lower curve in each plot is the associated error. The spectra have been smoothed to about $5\\AA$ resolution.} \\end{figure} \\clearpage Component 11 in our mixture model has very different characteristics from component 10.  It consists of almost purely SUO type objects (80\\%), and it has a much broader distribution of all four colors. Furthermore, all the points in component 11 are much redder than those in component 10.  Component 11 has captured a region of color space that is largely in the HiZ QSO region of the $(u-g)$ vs. $(g-r)$ color-color diagram of \\citet{Richards02}.  However, it also corresponds to a region of relatively high density of objects initially classified as quasars but then rejected following a cut on stellar density (see figure 2 of \\citet{Richards04}).  A visual inspection of component 11 objects labeled as SUOs indicates that a majority of these objects are low signal to noise G-K stars.  See Figure 6 for sample spectra from objects in component 11.  SDSS target selection is fainter for QSOs than stars and galaxies since their prominent broad features make QSOs easier to identify at low signal to noise.  Stars incorrectly targeted as QSOs are thus more likely to be classified as SUOs. \\clearpage \\begin{figure} \\epsscale{0.8} \\plotone{f5.eps} \\label{fig:Spec11} \\caption{\\small Five typical spectra of objects from component 11. These are probably low SNR G-K stars.  The relatively smooth line is the error level.  The spectra have been smoothed to about $5\\AA$ resolution.} \\end{figure} \\clearpage The small rms values for several of the components suggest, on the surface, a homogeneity of object types within the components.  For example, component 10 has a relatively small rms in all colors (see Table \\ref{tab:phot-stats}).  This is consistent with all the objects being owned by a single mixture component but it does not necessarily mean all the objects are the same.  That determination would have to be made with more detailed examination and comparison of the spectra of the objects.  Moreover, there are several possible reasons why the nonpredefined components contain significant numbers of points from other classes. First, for some objects, the fact that the object is assigned to the SUO class is indicative that there is significant uncertainty about its class of origin.  All objects that are identified as SUOs by the SDSS spectral pipeline are also visually inspected.  Some objects in the SUO class may really belong to one of the known classes, but were not classified as such in the SDSS spectral pipeline because the specificity of the pipeline (its ability to identify positive instances of an object class) is limited.  The procedure for labeling an object as spectroscopically unclassified involves cross-correlation with several standard templates and the determination of the confidence level of the cross-correlation.  When this confidence level is below 0.25 then the object is labeled as an SUO.  Subsequent visual inspection was unable to provide a confident classification of these objects.  If there are a significant number of such objects and if they have similar feature vectors, a nonpredefined cluster may be learned which contains many of these objects (and whose model parameters well-describe these objects).  However, such a cluster may also well-describe (unlabeled) known class objects, and thus may contain a significant number of such objects.  Another possibility is that some objects in the known classes are mislabeled and should really be members of the SUO class or classes.  While the mixed composition of nonpredefined components is partially explained by uncertainty and/or errors associated with the spectral pipeline object labelings, another possible explanation is that the features we chose were not powerful enough to fully distinguish between objects from the different classes.  In particular, since each photometric band is essentially a weighted average of the spectrum of the object, it is clear that significantly different spectra may produce similar photometric responses.  This suggests that one should be cautious when using only photometric features for classification purposes.  Using additional features, such as the photometry in $u$, $g$, $r$, $i$, and $z$, might help differentiate between classes. Certainly, including UV or IR data in a multispectral analysis would lead to more powerful discriminators.  Working directly with the high-dimensional object spectra would also substantially enhance the potential for class discrimination.  However, there are also many spectral features that are {\\it not} class-discriminating. This indicates the need for effective feature selection, to determine the (perhaps small) subset of features that are most important for distinguishing the different object classes.  Some recent approaches have been proposed for feature selection in high-dimensional mixture modeling, e.g. \\citep{Graham06}, which we hope to exploit in the near future.  \\clearpage  \\begin{figure} \\epsscale{0.8} \\plotone{f6.eps} \\label{fig:color-color14-15} \\caption{\\small Color-Color diagrams for components 14 and 15.  The symbols (colors refer to the online version) are: green 0--SUO, blue 1--star, orange 2--galaxy, red 3--qso, cyan 4--late type star. Component 15 contains mainly bluer objects, of which 56\\% were spectroscopically identified as galaxies.  Component 14 contains redder objects, of which 58\\% are identified as galaxies Contours are for all the data with levels 5\\%, 10\\%, 20\\%, 40\\%, 60\\% and 80\\% of the maximum.} \\end{figure}  \\clearpage  As mentioned above, a number of predefined components also have significant admixtures of various types of objects.  In particular we note components 14 and 15, which together have 81\\% of all the galaxies in our sample.  These two components split the galaxies into two clusters, as shown in Figure 6.  This is very similar to the bimodal distributions found by \\citet{Strateva01} in the SDSS data and more recently seen also in combined GALEX and SDSS observations \\citep{Seibert05}.  The two components contain essentially equal fractions of galaxies (60\\%), though some fraction of the SUO class objects could also be galaxies.  They also have a low fraction of quasars.  Our modeling procedure allows for the possibility that several mixture components may be needed to describe a single class. However, in this case no other components appear to overlap this region of color-color space.  It is hard to find components for the unsupervised model which correspond well to components SS14 and SS15 that capture some of the bimodal galaxy distribution.  US5 alone captures 63\\% of the galaxies, but the remaining galaxies are fairly evenly spread among four other components.  Notably, the separation into two clusters by the semisupervised mixture model is a byproduct of our main objective of looking for subclasses in the SUOs.  However, even with a sample containing only 2000 galaxies, we find that our mixture model does a very good job of separating the galaxy distribution into two separate clusters.  It also appears to result in a better separation than that due to the unsupervised approach.  We hope to perform a similar analysis using a much larger sample of galaxies and look for a more detailed decomposition of the main galaxy class."
    },
    "0512/astro-ph0512355_arXiv.txt": {
        "abstract": "The expectation of explaining cosmological observations without requiring new energy sources is forsooth worthy of investigation. In this letter, a new kind of Cardassian models, called exponential Cardassian models, for the late-time universe are investigated in the context of the spatially flat FRW universe scenario. We fit the exponential Cardassian models to current type Ia supernovae data and find they are consistent with the observations. Furthermore, we point out that the equation-of-state parameter for the effective dark fluid component in  exponential Cardassian models can naturally cross the cosmological constant divide $w=-1$ that observations favor mildly without introducing exotic material that destroy the weak energy condition. ",
        "introduction": "The current accelerating expansion of the universe indicated by the astronomical measurements from Supernovae type Ia (SNeIa) \\cite{SNeIa} (see \\cite{SnST,SnLS} for most recent results) as well as accordance with other observations such as the cosmic microwave background (CMB) \\cite{WMAP} and galaxy power spectra \\cite{SDSS} becomes one of the biggest puzzles in the research field of cosmology. There are lots of approaches to unriddle this puzzle. One popular theoretical explanation approach is to assume that there exists a mysterious energy component, dubbed dark energy, with negative pressure, or equation of state with $w=p/\\rho < 0$ that currently dominates the dynamics of the universe (for a review see \\cite{DEReivew}).  Such a component makes up $70\\%$ of the energy density of the universe yet remains elusive in the context of general relativity and the standard model of particle physics. In recent years, many candidates for dark energy have been explored. Besides a cosmological constant, one popular candidate source of this missing energy component is a slowly evolving and spatially homogeneous scalar field, referred to as \"quintessence\" with $w>-1$ \\cite{quintessence} and \"phantom\" with $w<-1$ \\cite{phantom}, respectively. Since current observational constraint on the equation of state of dark energy lies in a relatively large range around the so-called cosmological constant divide $w_X=-1$, it is still too early to rule out any of the above candidates.  On the other hand, general relativity (GR) is very well examined in the solar system, in observation of the period of the binary pulsar, and in the early universe, via primordial nucleosynthesis. However, no one has so far tested in the ultra-large length scales and low curvatures characteristic of the Hubble radius today. Therefore, it is a priori believable that Friedmann equation is modified in the very far infrared, in such a way that the universe begins to accelerate at late time. Freese and Lewis \\cite{Freese2002} construct so-called Cardassian universe models that incarnates this hope. In Cardassian models \\cite{Freese2002,Lazkoz,Freese03}, the universe is flat and accelerating, and yet contains only matter (baryonic or not) and radiation. But the usual Friedmann equation governing the expansion of the universe is modified to be \\be \\label{friedmanEQ1} H^2\\equiv\\left(\\frac{\\dot{a}}{a}\\right)^2=\\frac{8\\pi G}{3}g(\\rho), \\ee where $\\rho$ consists only of matter and radiation, $H$ is the Hubble \"parameter\" which is a function of time, $a$ is the scale factor of the universe, and $G = 1/{m_{pl}^2}$ is the Newtonian gravitational constant. Note that as required by  inflation scenario and observations of CMB, the geometry of the universe is flat, therefore, there are no curvature terms in the above equation. Perhaps the most interesting feature of Cardassian models is that although being matter dominated, they may be accelerating and may still reconcile the indications for a flat universe ($\\Omega_T=1$) from CMB observations with clustering estimates that point consistently to $\\Omega_{m,0}=0.3$ with no need to invoke either a new dark component or a curvature term. The expectation of explaining cosmological observations without requiring new energy sources or certainly worthy of investigation.  For any suitable Cardassian model, there exist at least three requirements that should be satisfied. Firstly, the function $g(\\rho)$ should returns to the usual form $\\rho$ at early epochs in order to recover the thermal history of the standard cosmological model and the scenario for the formation of large scale structure. Secondly, $g(\\rho)$ should takes a different form  at late times $z\\sim \\mathcal{O}(1)$ in order to drive an accelerated expansion as indicated by the observation of SNeIa \\cite{SNeIa,SnST,SnLS}. Finally, the classical solution of the expansion should be stable, \\ie,  the sound speed $c_s^2$ of classical perturbations of the total cosmological fluid around homogeneous FRW solutions cannot be negative.   For the original power-law Cardassian model $ \\label{PowerLawModel} g(\\rho)=\\rho +B\\rho^n,$ where $B$ and $n<2/3$  are two constants \\cite{Freese2002}, the second term in the above equation behave like an effective cosmic dark energy component that drive the universe accelerate at late times and is negligible at early epochs. However, the sound speed of cosmological fluid in this model is not guaranteed to be positive. So this model should only be considered as an effective description at scales where the sound speed is positive \\cite{Gondolo02}. The generalized power-law Cardassian models, such as MP Cardassian model \\cite{Freese03}, satisfies all of above requirements, resorting to an additional parameter.  In this paper, we investigate another kind of Cardassian models, dubbed Exponential Cardassian models. In the next section, we construct two models that embodies the elegant idea of Cardassian universe. In section 3, we fit the parameters of the models to the type Ia supernovae observations. The issue on the sound speed of the cosmological fluid $c_s^2$ is addressed in section 4. And in the last section, a brief discussion is included. ",
        "conclusions": "\\label{} In above sections, we have investigated a new kind of Cardassian  models of the universe,dubbed Exponential Cardassian universe. Contrary to the original power-law Cardassian model \\cite{Freese2002}, the equation-of-state parameter $w$ of effective dark fluid is dependent on time that can cross the cosmological constat divide $w_{\\Lambda}=-1$ from $w_X>-1$ to $w_X<-1$ as the observations mildly indicate.  However, it is worth noting that these two models are available at much large scales where the matter density $\\rho_m \\sim \\rho_c$ and evolve in the light of $\\propto (1+z)^{-3}$. As discussed by Gondolo and Freese \\cite{Gondolo02}, the gradient of the effective Cardassian pressure $\\nabla p_{card}$ should be able to neglected compared to that of the ordinary pressure $\\nabla p_{m}$ at the scales of galaxies and galaxy clusters where the matter density is much larger than cosmic average density. In a typical galaxy where $\\rho_m \\sim 100 \\rho_c \\sim 100 \\rho_{card}$ and $p_m=\\rho_m \\Sigma^2$ with velocity dispersion $\\Sigma=300\\; \\mbox{km/s}$, if we assume that $\\nabla p_{card}\\ll\\nabla p_{m}$, we should require the parameter $n$ in both models considered above are greater than $3$. It is remarkable that this value of $n$ is compatible with the data of SNeIa observations. Therefore, these two models are both available on galactic scales.  Observationally, the exponential Cardassian models are compatible with the data of SNeIa. The further work that should do is to check it out by other cosmological and astrophysical observations such as CMB and large scale structures (LSS) as many works (for example, \\cite{koivisto,constraints}) for power-law Cardassian models, which is out of the scope of this paper. However, it is worth noting that Koivisto \\emph{et al} \\cite{koivisto} pointed out that in the MP Cardassian model, the late integrated Sachs-Wolf effect is typically very large duo to the suppression of fluctuations in Cardassian fluid at late times, and is then not compatible with the observations of CMB unless the parameters is very close to the $\\Lambda$CDM model. In the exponential Cardassian models, we assume that Cardassian fluctuations is induced by cold dark matter (i.e., case III in Ref. \\cite{koivisto}) and fluctuations in the cold dark matter and in the Cardassian fluid are related adiabatically \\begin{equation} \\delta=\\left(\\frac{d\\log \\rho_T}{d\\log\\rho}\\right)^{-1}\\delta_T=\\frac{\\delta_T}{1+w_T}, \\end{equation} where $\\delta$ and $\\delta_{T}$ denote the fluctuations in cold dark matter and Cardassian fluid, respectively. According to the expressions of $w_T$, Eq.(\\ref{PLptotalmatterradiation1}) for model I and Eq. (\\ref{PLptotalmatterradiation1.2}) for model II, it is not hard to find that the Cardassian fluctuations are both not heavily suppressed as those in MP Cardassian model  which $\\delta\\sim a^{3q}\\delta_{T}$ at late times \\cite{koivisto}. Therefore, in the exponential Cardassian models the above problem is alleviated to some extent.  Theoretically, these models, in fact as well as power-law Cardassian models, are phenomenologically  described in a fluid approach, therefore, we need further works to find a simple fundamental theory which can derive the form of $g(\\rho)$ presented above."
    },
    "0512/astro-ph0512449_arXiv.txt": {
        "abstract": "We survey a sample of 32 M5-M8 stars with distance $<$ 40pc for companions with separations between 0.1'' and 1.5'' and with $\\rm{\\Delta m_i<5}$. We find five new binaries with separations between 0.15'' and 1.1'', including a candidate brown dwarf companion. The raw binary fraction is $16^{+8}_{-4}\\%$ and the distance bias corrected fraction is $7^{+7}_{-3}\\%$, for companions within the surveyed range. No systems with contrast ratio $\\rm{\\Delta m_i>1}$ were found, even though our survey is sensitive to $\\rm{\\Delta m\\leq5}$ (well into the brown dwarf regime). The distribution of orbital radii is in broad agreement with previous results, with most systems at 1-5AU, but one detected binary is very wide at $46.8\\pm5.0$AU. We also serendipitously imaged for the first time a companion to Ross 530, a metal-poor single-lined spectroscopic binary. We used the new Lucky Imaging system LuckyCam on the 2.5m Nordic Optical Telescope to complete the 32 very low mass star SDSS i' and z' survey in only 5 hours of telescope time. ",
        "introduction": "There are compelling reasons to search for companions to nearby stars. In particular, the properties of binary systems provide important clues to their formation processes. Any successful model of star formation must be able to account for both the frequency of multiple star systems and their properties (separation, eccentricity and so forth) -- as well as variations in those properties as a function of system mass. In addition, the orbits of binary systems provide us with the means to directly measure the mass of each component in the system. This is fundamental to the calibration of the mass-luminosity relation (MLR: \\citealt{Henry_1993, Henry_1999, Segransan_2000}).  The stellar multiplicity fraction appears to decrease with decreasing primary mass (eg. \\citealt{Siegler_2005}). Around 57\\% of solar-type stars (F7--G9) have known stellar companions \\citep{Abt_1976, Duquennoy_1991}, while imaging and radial velocity surveys of early M dwarfs suggest that between 25\\% \\& 42\\% have companions \\citep{Henry_1990, Fischer_1992, Leinert_1997, Reid_1997}. Later spectral types have been studied primarily with high resolution adaptive optics imaging: \\citealt{Close_2003} and \\citealt{Siegler_2005} find binary fractions of around $10$--$20\\%$ for primary spectral types in the range M6--L1. \\citealt{Bouy_2003} and \\citealt{Gizis_2003} find that 10--15\\% of L dwarfs have companions, and \\citealt{Burgasser_2003} find that 10\\% of T dwarfs have binaries. These very low mass (VLM) \\mbox{M, L and T} systems appear to have a tighter and closer distribution of orbital separations, peaking at around 4\\,AU compared to 30\\,AU for G dwarfs \\citep{Close_2003}.  However, each of these surveys have inevitably different (and hard to quantify) sensitivities, the effect of which is especially evident in the large spread in the derived multiplicity of early M-dwarfs. In particular, high-resolution imaging surveys are sensitive only to companions wider than $\\sim$0.1'' while radial velocity surveys are much more sensitive to closer (shorter period) companions. Maxted \\& Jeffries\\,(2005), by examining a small sample of radial velocity measurements, estimate that accounting for systems with $a<3$\\,AU could increase the overall observed VLM star/BD binary frequency to 32--45\\%. In addition, each survey's target sample has a different selection of target stellar parameters (as well as different incompletenesses and biases), leading to difficulties in the comparison and pooling of results for the surveys.  This paper details the first results of an ongoing effort by our group to greatly increase the number of known VLM binary systems. A large sample size is made possible by the uniquely low observation overheads offered by the new high-resolution imaging system LuckyCam. We present results here from a trial 32-star sample, completed in only 5 hours of on-sky time. The detected binaries have been very briefly described as part of a larger sample in \\citet{Law_binaries_05}; we here undertake detailed investigations of the systems and sample.  Our Lucky Imaging system, LuckyCam, takes a sequence of images at $>$10 frames per second using a very low noise L3CCD based conventional camera. Because the atmospheric turbulence affecting the images changes on very short timescales, there are rapid ($<$100ms coherence time) variations in the image quality of the frames. To construct a high resolution long exposure image we select, align and co-add only those frames which meet a quality criterion. By varying the criterion we can trade off sensitivity against ultimate resolution. The technique is described in more detail in \\citet{Tubbs_2002, Law05}. Lucky Imaging is an entirely passive technique, allowing data to taken as soon as the telescope is pointed, and thus leading to very low time overheads. With stars brighter than $\\rm{+15m}$ in $<0.6''$ seeing, resolutions very close to the diffraction limit (Strehl ratios $>$ 0.1) of the 2.5m Nordic Optical Telescope (NOT) in i' band are regularly obtained. Figure \\ref{FIG:Profiles} shows examples of the general form of the Lucky Imaging point spread function (PSF).  In this paper we present five new VLM binaries and evaluate the utility of LuckyCam for programmes of this type. In section \\ref{SEC:Sample} we define the sample of VLM stars imaged in this survey. In section \\ref{SEC:Obs} we describe the observations, Lucky Imaging, the data reduction techniques and the survey sensitivity. Section \\ref{SEC:Results} describes the results of our survey and details the properties of the detected binaries. In section \\ref{SEC:Discussion} we discuss the results in the context of other surveys. We conclude in section \\ref{SEC:Concs}. \\begin{figure} \\centering \\resizebox{\\columnwidth}{!} { \\includegraphics{plot_good_seeing.eps} } \\caption{Typical radial Lucky Imaging profiles, azimuthally averaged from the point spread function of a star imaged at 30Hz in 0.6'' seeing. 100\\% selection corresponds to fast shift-and-add (tip/tilt correction) imaging. When selecting 1\\% of frames the Strehl ratio is almost tripled relative to the 100\\% selection, while the light from the star is concentrated into an area approximately four times smaller.} \\label{FIG:Profiles} \\end{figure} \\begin{figure} \\centering \\resizebox{\\columnwidth}{!} { \\includegraphics[angle=-90]{plot_vmk.eps} } \\caption{The observed sample (crosses), plotted in a V/V-K colour-magnitude diagram. The background distribution shows all stars in the LSPM-North catalogue. Circles show 38 spectroscopically confirmed M6 and later dwarfs from \\citealt{Cruz03}, which also appear in the LSPM-North; all but two (both M6) are recovered by our selection criteria.} \\label{FIG:Sample} \\end{figure}    \\begin{table*} \\centering \\caption{The observed sample. The quoted V \\& K magnitudes are taken from the LSPM catalogue. K magnitudes are based on 2MASS photometry; the LSPM-North V-band photometry is estimated from photographic $\\rm{B_J}$ and $\\rm{R_F}$ magnitudes and is thus approximate only, but is sufficient for spectral type estimation (figure \\ref{FIG:Sample}). Spectral types and distances are estimated from the V \\& K photometry and the young-disk photometric parallax relations described in \\citet{Leggett_1992}. Spectral types are accurate to approximately 0.5 spectral classes and distances to $\\sim30$\\%.  \\label{Tab:Sample}} \\begin{tabular}{llrrllrc} \\hline LSPM ID & 2MASS ID & V & V-K & PM / arcsec/yr& Estimated spectral type & Photom. dist/pc & Newly detected companion? \\\\ \\hline LSPM J1235+1318 & 12351726+131805 &  18.0 & 7.7 & 0.219 & M6.5 & 14 & $\\ast$ \\\\ LSPM J1235+1709 & 12351850+170937 &  19.3 & 7.5 & 0.570 & M6.5 & 29 &      \\\\ LSPM J1246+0706 & 12460939+070624 &  17.7 & 7.1 & 0.549 & M6.0 & 19 &    \\\\ LSPM J1303+2414 & 13034100+241402 &  19.6 & 7.9 & 0.370 & M7.0 & 24 &     \\\\ LSPM J1305+1934 & 13053667+193456 &  18.4 & 7.3 & 0.554 & M6.5 & 22 &     \\\\ LSPM J1314+1320 & 13142039+132001 &  15.9 & 7.2 & 0.307 & M6.0 & 7.7 & $\\ast$ \\\\ LSPM J1336+1022 & 13365393+102251 &  18.7 & 7.3 & 0.381 & M6.5 & 26 &     \\\\ LSPM J1341+0805 & 13413291+080504 &  18.5 & 7.4 & 0.269 & M6.5 & 22 &     \\\\ LSPM J1354+0846 & 13540876+084608 &  19.3 & 8.2 & 0.219 & M7.0 & 17 &     \\\\ LSPM J1423+1318 & 14231683+131809 &  17.9 & 7.3 & 0.174 & M6.5 & 18 & $\\ast$ \\\\ LSPM J1423+1426 & 14234378+142651 &  17.9 & 7.7 & 0.638 & M6.5 & 13 &     \\\\ LSPM J1428+1356 & 14280419+135613 &  18.3 & 8.3 & 0.605 & M8.0 & 10 &     \\\\ LSPM J1432+0811 & 14320849+081131 &  16.3 & 7.2 & 0.455 & M6.0 & 9.2 &     \\\\ LSPM J1440+1339 & 14402293+133923 &  19.0 & 7.7 & 0.337 & M6.5 & 22 &     \\\\ LSPM J1454+2852 & 14542356+285159 &  18.6 & 7.0 & 0.212 & M6.0 & 32 &      \\\\ LSPM J1516+3910 & 15164073+391048 &  17.1 & 7.3 & 0.224 & M6.5 & 12 &     \\\\ LSPM J1554+1639 & 15540031+163950 &  19.9 & 7.8 & 0.529 & M7.0 & 30 &     \\\\ LSPM J1605+6912 & 16050677+691232 &  19.3 & 7.5 & 0.224 & M6.5 & 29 &     \\\\ LSPM J1606+4054 & 16063390+405421 &  17.6 & 7.6 & 0.735 & M6.5 & 12 &     \\\\ LSPM J1622+4934 & 16225554+493457 &  19.4 & 7.2 & 0.316 & M6.0 & 39 &     \\\\ LSPM J1626+2512 & 16263531+251235 &  19.3 & 7.5 & 0.271 & M6.5 & 29 &     \\\\ LSPM J1646+3434 & 16463154+343455 &  16.6 & 7.0 & 0.550 & M6.0 & 13 &     \\\\ LSPM J1647+4117 & 16470576+411706 &  18.5 & 7.4 & 0.289 & M6.5 & 22 &     \\\\ LSPM J1653+0000 & 16531534+000014 &  18.6 & 7.8 & 0.287 & M7.0 & 17 &     \\\\ LSPM J1657+2448 & 16572919+244850 &  18.8 & 7.5 & 0.391 & M7.5 & 23 &     \\\\ LSPM J1703+5910 & 17031418+591048 &  18.8 & 7.0 & 0.572 & M6.0 & 35 &     \\\\ LSPM J1735+2634 & 17351296+263447 &  19.1 & 9.0 & 0.349 & M9.0 & 9.2 & $\\ast$ \\\\ LSPM J1741+0940 & 17415439+094053 &  18.7 & 7.8 & 0.435 & M7.0 & 17 &    \\\\ LSPM J1758+3157 & 17580020+315726 &  18.2 & 7.1 & 0.158 & M6.0 & 24 &    \\\\ LSPM J1809+2128 & 18095137+212806 &  18.3 & 7.1 & 0.193 & M6.0 & 25 & $\\ast$ \\\\ LSPM J1816+2118 & 18161901+211816 &  19.0 & 7.2 & 0.171 & M6.0 & 32 &      \\\\ LSPM J1845+3853 & 18451889+385324 &  19.4 & 8.4 & 0.408 & M8.0 & 16 &     \\\\  \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{SEC:Concs} LuckyCam's very low time overheads allowed a 32 VLM star sample to be imaged at high angular resolution in only 5 hours on a 2.5m telescope. Uniquely, this survey was performed in the visible, and thus complements the near-infrared surveys of \\citealt{Bouy_2003, Burgasser_2003, Close_2003, Gizis_2003, Delfosse2004} and \\citealt{Siegler_2005}. Lucky Imaging shows great promise for surveys of this type, offering a unique ground-based faint-guide-star visible-light imaging capability.  We have detected five new VLM binaries in a 32 star sample, giving a raw binary fraction of $16^{+8}_{-4}\\%$ and a distance-bias corrected fraction of $7^{+7}_{-3}\\%$. Primaries were M5.5 to M8; most secondaries were only slightly redder than the primaries. However, one newly found system is a possible M-dwarf brown-dwarf binary. The distribution of orbital radii is in broad agreement with previous results, with a peak at 1-5AU, but one newly detected binary is very wide, at $46.8\\pm5.0$AU. No systems with a high contrast ratio were detected, even though the survey is sensitive well into the brown dwarf regime.  \\begin{figure} \\centering \\resizebox{\\columnwidth}{!} { \\includegraphics{plot_sep_hist.eps} } \\caption{The $1\\sigma$ range of orbital radii for each detected binary, compared with a histogram of the previously known sample collated in \\citet{Siegler_2005} (and the 33 AU VLM binary described in \\citealt{Phan05}). Poisson $1\\sigma$ error bars are shown for the histogram - and illustrate the necessity for increased sample sizes. For reasons of clarity, the 200AU system described in \\citet{Luhman04} is not displayed.} \\label{FIG:Sep_hist} \\end{figure}"
    },
    "0512/astro-ph0512480_arXiv.txt": {
        "abstract": "We present the results of an archival {\\it XMM-Newton} study of the bright X-ray point sources (L$_X > 10^{38}$\\,erg\\,s$^{-1}$) in 32 nearby galaxies. From our list of approximately 100 point sources, we attempt to determine if there is a low-state counterpart to the Ultraluminous X-ray (ULX) population, searching for a soft-hard state dichotomy similar to that known for Galactic X-ray binaries and testing the specific predictions of the IMBH hypothesis.  To this end, we searched for ``low-state\" objects, which we defined as objects within our sample which had a  spectrum well fit by a simple absorbed power law, and ``high-state\" objects, which we defined as objects better fit by a combined blackbody and a power law.  Assuming that ``low-state'' objects accrete at approximately 10\\% of the Eddington luminosity \\citep{don03} and that ``high-state\" objects accrete near the Eddington luminosity we further divided our sample of sources into low and high state ULX sources.  We classify 16 sources as low-state ULXs and 26 objects as high-state ULXs.   As in Galactic black hole systems, the spectral indices, $\\Gamma$, of the low-state objects, as well as the luminosities, tend to be lower than those of the high-state objects.  The observed range of blackbody temperatures for the high state is 0.1-1\\,keV, with the most luminous systems tending toward the lowest temperatures.  We therefore divide our high-state ULXs into candidate IMBHs (with blackbody temperatures of approximately 0.1\\,keV) and candidate stellar mass BHs (with blackbody temperatures of approximately 1.0\\,keV).   A subset of the candidate stellar mass BHs have spectra that are well-fit by a Comptonization model, a property similar of Galactic BHs radiating in the ``very-high\" state near the Eddington limit. ",
        "introduction": "Through X-ray observations of nearby galaxies, a class of Ultraluminous X-ray (ULX) sources has emerged.  These are pointlike, non-nuclear sources with observed X-ray luminosities greater than $10^{39}$\\,erg\\,s$^{-1}$ \\citep{mil04}. Of most interest are those sources with bolometric luminosities in excess of the Eddington limit for a 20 M$_{\\sun}$ black hole, or L$_{bol} > 2.8 \\times 10^{39}$\\,erg\\,s$^{-1}$.  The true nature of these sources is unclear, and this class most likely includes several different types of objects.  Though some of these sources are located within a few parsecs of their host galaxy's dynamical center, they do not exhibit many of the characteristics of active galactic nuclei (AGN).  Because the ratio of X-ray to optical flux is a factor of 10 greater than that of AGN \\citep{and03,sto83}, these objects are fairly easy to recognize in X-ray imaging data.  Assuming that the Eddington limit is obeyed by black hole accretion, the existence of such luminous non-AGN sources presents a puzzle. Several models have been proposed to account for the high luminosities of the ULXs. Among these are relativistic and non-relativistic beaming from stellar-mass black hole systems \\citep{rey97, kin01, kord02} and accretion of matter into intermediate mass black holes (IMBHs). In several ULX systems (NGC 1313 X-2, M81 X-9, etc.), detection of emission nebulae surrounding the ULX supports isotropic emission from the central source \\citep{pak03}, which cannot be described through beaming.  Further, a number of ULX (NGC1313 X-1, etc.) X-ray spectra are best fit with combined multi-component blackbody (MCD) and power law fits, similar to Galactic black holes in their high-state.  Recently, \\citet{mill04} find that many spectral fits of ULXs require cool accretion disk temperatures of approximately 100\\,eV. The theoretical relationship between black hole mass and disk temperature (T $\\propto M^{-1/4}$) has been observed to hold true for stellar mass (typically around 1\\,keV) and supermassive (around 10-100\\,eV) black holes \\citep{mak00, por04, gie04}. Using these scaling relations, the cool accretion disk ULXs would correspond to a population of high-state IMBHs with masses of $\\approx 16 - 10^4$\\,M$_{\\sun}$.  If ULXs do not obey the Eddington limit, they could be the result of an outburst (such as can occur in low mass X-ray binaries within our own Galaxy). \\citet{jon04} find evidence for approximately 5 Galactic black hole X-ray binaries which exhibit luminosities in the ULX range during outbursts.  These sources would appear as transient ULXs.  The typical time scale for outburst of Galactic X-ray transients is a few days to rise from quiescent level with a decline from peak brightness to quiescent value of 30 - 40 days \\citep{che97}. Another possible explanation is super-Eddington emission from accretion disks surrounding stellar mass black holes \\citep{beg02, ebi03}.  Sources of this type would be expected to have soft X-ray components well modeled by hot accretion disks ($\\approx 1.3-2.0$\\,keV) similar to superluminal X-ray sources in the Galaxy (e.g. Belloni et al. 1997).  Likely, ULXs include a variety of different objects with both isotropic and non-isotropic emitters. However, if some ULXs do indeed represent a class of high-state IMBHs, similar to the high/soft (thermal dominated) state stellar mass black holes in our galaxy, we might also expect to see the low-state objects from this same population. In Galactic black hole systems, the low-state is generally characterized by lower luminosity, with L$< 0.1$\\,L$_{Edd}$ \\citep{don03}, and a power law photon spectrum, typically with index $\\Gamma \\approx 1.7$ \\citep{mcc04}.  Indeed, the existence of some ULX sources (IC 342 X-1, NGC 5204 X-1) as possible low-hard (pure power law) state IMBHs, well-fit by simple absorbed power laws, have been noted from Chandra observations by \\citet{rob04}. In this study we seek to test a direct prediction of the IMBH hypothesis; namely, whether there is a class of sources with properties consistent with what we expect of low-state IMBHs.  This requires two major assumptions: (1) that the emission from ULXs is isotropic and (2) that IMBHs exhibit the same states (whose classification was based on luminosity and spectral form) as stellar mass black holes.    Our goal is to find these ``low-state'' sources, if they exist, classify the properties of both high-state and low-state ULXs, and test whether these data are consistent or inconsistent with the predictions of the IMBH hypothesis.  We present the results of a detailed analysis of ULXs in nearby galaxies observed with the European Space Agency's {\\it XMM-Newton} observatory.  Only {\\it XMM-Newton} provides the count rates and bandpass necessary to distinguish different spectral models for most ULXs, accurately determine both the temperature of the thermal component expected for high-state objects, and determine whether this component is required in the spectral modeling of these objects.  Since the XMM-Newton X-ray spectra of ULXs are similar in quality to spectra for Galactic X-ray binaries obtained in the 1980s, our spectral classification in this paper will remain purely schematic.  Thus, our classifications as low and high state objects are a first approximation, based on the quality of the spectra available.  In Section 2, we detail the observations examined from the {\\it XMM-Newton} archives and explain the data analysis for the individual point sources. In Section 3, we discuss the spectral fitting technique as well as simulations we conducted to determine their validity.  We discuss the implications of our results in Section 4. ",
        "conclusions": "We present the results of an {\\it XMM} survey of the ULX population in nearby galaxies.  In this study, we assumed that ULXs are isotropic emitters. For our selected ULX sources (which excluded transient sources and supernovae), this assumption was supported by the finding that 37/42 of our ULXs were found to be `on' in ROSAT observations.  This implies that theses sources exhibited high luminosities for time scales of at least 10\\,years, a property that is not seen in Galactic Eddington-limit exceeding sources (such as black hole X-ray binaries undergoing an outburst). We also assumed that if some ULX sources represent a class of IMBH X-ray binaries, they would exhibit spectral states analogous to Galactic stellar mass black hole X-ray binaries.  This is the hypothesis we set out to test, classifying a source as a ULX based on (1) spectral form, (2) luminosity, and (3) coincidence of the X-ray source within the optical host galaxy. Due to the quality of spectra available for these distant X-ray sources, our classification of spectral form is really a first approximation describing the basic curvature of the spectrum.  Through this study, we have found that there exists a population of objects whose X-ray spectral properties closely match the low/hard state spectra of Galactic black holes, but whose luminosities lie in the range of L$_{bol} \\approx 2 \\times 10^{38} - 1 \\times 10^{40}$\\,erg\\,s$^{-1}$.  In the Milky Way, black holes with these spectral properties radiate at only $\\approx 0.05$ of the Eddington limit.  If this is also true for this population, it indirectly implies that these objects have a mass greater than $\\approx 30$\\,M$_{\\sun}$ ranging up to 1500\\,M$_{\\sun}$ and thus should be IMBHs.  The existence of such objects was ``predicted'' on the basis that the ULXs previously studied shared the X-ray spectral characteristics of high-state Galactic black holes; namely, an X-ray spectrum best fit by a combined blackbody and a power law \\citep{mil03}, but with much higher luminosities.  If these objects are high-state IMBHs, the corresponding low-state objects should also exist.  Our survey has also uncovered a large population of objects whose X-ray spectra are well modeled by the canonical description of Galactic black holes in the high-state (thermal dominated), a black hole with a steep power law, but whose bolometric luminosities exceed $2 \\times 10^{39}$\\,erg\\,s$^{-1}$, ranging up to $10^{41.5}$\\,erg\\,s$^{-1}$ and whose blackbody temperatures are less than 0.3\\,keV. If these objects are radiating at $\\approx 1/2$ the Eddington limit like their Milky Way counterparts their implied masses are from $30 - 3000$\\,M$_{\\sun}$, a range very similar to that implied by the low-state objects.  Using the M$^{-1/4}$ scaling of mass to temperature, the observed spectral temperatures give masses of $500 - 10^{4}$\\,M$_{\\sun}$ a considerably larger value.  In general agreement with the expectations of the IMBH hypothesis, the objects with high-state spectra are more luminous than those with low-state spectra.  We note that these results have required the high signal to noise of {\\it XMM} in order to discern the spectrum of these objects.  Many of these objects have also been observed by Chandra and their spectra have been well-fitted by simple power laws.  In addition to classification of the sources, we investigated some of the properties of the ULX sources.  We found a gap in the temperature distribution of high/soft state ULXs.  This gap may indicate a gap in mass distribution, which may provide clues to the nature of ULXs. We also found that our ULXs are persistent sources (not transients) which occupy regions on the color-color diagram of \\citet{don03} also occupied by Galactic black hole sources. Lastly, the existence of a substantial population of ULXs in nearby dwarf and other low star formation rate galaxies argues that (in agreement with \\citet{pta04,swa04}) there is more than one source term for the origin of ULXs, with at least some of them not being associated with recent star formation, at least statistically.   We conclude, from an X-ray spectral and luminosity point of view, that our data are consistent with many of these objects having the properties expected of an IMBH population.  However, we also find two other populations of objects, those whose blackbody temperature and luminosity correspond to that of stellar mass black holes with kT $\\approx 1$\\,keV and $\\log L_X$ less than $2 \\times 10^{39}$\\,erg\\,s$^{-1}$ and a small population of objects whose X-ray spectra and luminosities are consistent with that of stellar mass black holes in the very high state.  Thus, ULX selected purely on the basis of $0.3 - 10$\\,keV X-ray luminosities are a composite class with $\\approx 1/4$ being ``normal'' stellar mass black holes and the rest being consistent with a population of IMBHs.  In a follow-up paper we will discuss the environments of these objects as revealed by {\\it XMM} OM UV imaging and the implications this has for the origin of ULXs."
    },
    "0512/astro-ph0512163_arXiv.txt": {
        "abstract": "We present high-resolution ultraviolet spectra of absorption-line systems toward the low$-z$ QSO \\hs0624 ($z_{\\rm QSO} = 0.3700$). Coupled with ground-based imaging and spectroscopic galaxy redshifts, we find evidence that many of these absorbers do not arise in galaxy halos but rather are truly integalactic gas clouds distributed within large-scale structures, and moreover, the gas is cool ($T < 10^{5}$ K) and has relatively high metallicity ($Z = 0.9 Z_{\\odot}$).  {\\it HST} Space Telescope Imaging Spectrograph (STIS) data reveal a dramatic cluster of 13 \\ion{H}{1} \\lya\\ lines within a 1000 \\kms\\ interval at $z_{\\rm abs} = 0.0635$. We find 10 galaxies at this redshift with impact parameters ranging from $\\rho = 135 h_{70}^{-1}$ kpc to 1.37 $h_{70}^{-1}$ Mpc. The velocities and velocity spread of the \\lya\\ lines in this complex are unlikely to arise in the individual halos of the nearby galaxies; instead, we attribute the absorption to intragroup medium gas, possibly from a large-scale filament viewed along its long axis.  Contrary to theoretical expectations, this gas is not the shock-heated warm-hot intergalactic medium (WHIM); the width of the \\lya\\ lines all indicate a gas temperature $T \\ll 10^{5}$ K, and metal lines detected in the \\lya\\ complex also favor photoionised, cool gas.  No \\ion{O}{6} absorption lines are evident, which is consistent with photoionisation models.  Remarkably, the metallicity is near-solar, [M/H] $= -0.05 \\pm 0.4$ ($2\\sigma$ uncertainty), yet the nearest galaxy which might pollute the IGM is at least 135~$h_{70}^{-1}$ kpc away.  Tidal stripping from nearby galaxies appears to be the most likely origin of this highly enriched, cool gas.  More than six Abell galaxy clusters are found within $4^{\\circ}$ of the sight line suggesting that the QSO line of sight passes near a node in the cosmic web.  At $z \\approx$ 0.077, we find absorption systems as well as galaxies at the redshift of the nearby clusters Abell 564 and Abell 559. We conclude that the sight line pierces a filament of gas and galaxies feeding into these clusters.  The absorber at $z_{\\rm abs}$ = 0.07573 associated with Abell 564/559 also has a high metallicity with [C/H] $> -0.6$, but again the closest galaxy is relatively far from the sight line ($\\rho = 293\\,h_{70}^{-1}$ kpc).  The Doppler parameters and \\ion{H}{1} column densities of the Ly$\\alpha$ lines observed along the entire sight line are consistent with those measured toward other low$-z$ QSOs, including a number of broad ($b>40$~\\kms) \\lya\\ lines. ",
        "introduction": "In cold  dark matter cosmology,  the initially smooth  distribution of matter in the universe is  expected to collapse into a complex network of filaments and voids, structures which have been termed the ``cosmic web''. The filamentary distribution of galaxies in the nearby universe has been  revealed in detail  by recent large galaxy  redshift surveys such as  the 2dFGRS (Colless et  al.  2001, Baugh et  al.  2004), the Sloan Digital  Sky Survey (SDSS, Stoughton et  al.  2002, Doroshkevich et al.  2004) and the  2$\\mu m$ All  Sky Survey (2MASS, Maller  et al. 2002).  Numerical  simulations  successfully  reproduce  this  network (Jenkins et al. 1998; Colberg  et al. 2004) and indicate that galaxies are only the tip of the iceberg  in this cosmic web (Katz et al. 1996; Miralda-Escud\\'{e}  et al.  1996).   Hydrodynamic simulations  suggest that at  the present epoch, in  addition to dark  matter and galaxies, the filaments are also composed of a mixture of cool, photoionised gas (the  low$-z$  remnants of  the  \\lya\\  forest)  and a  shock  heated, low-density  gaseous  phase   at  temperatures  between  $10^5$~K  and $10^7$~K  that contains most  of the  baryonic mass,  the ``warm-hot'' intergalactic  medium  (WHIM,  Cen   \\&  Ostriker  1999;  Dav\\'{e}  et al. 1999).  Observational constraints  on the physical  conditions, distribution,a nd  metal  enrichment  of  gas  in the  low-redshift  cosmic  web  are currently quite  limited. The  existence of the  WHIM appears to  be a robust prediction  of cosmological  simulations (Dav\\'e et  al. 2001). Thus,  observational efforts  are increasingly  being invested  in the search  for WHIM  gas  and, more  generally,  the gaseous  filamentary structures predicted by the models. Large-scale gaseous filaments have been detected in X-ray emission (Wang et al. 1997; Scharf et al. 2000; Tittley \\& Henriksen 2001; Rines  et al 2001). However, X-ray emission studies  with current  facilities  predominantly reveal  gas which  is hotter  and denser  than  the WHIM;  this  X-ray emitting  gas is  not expected to contain a substantial portion of the present-epoch baryons (Dav\\'{e} et al.  2001).  The  most promising method for observing the WHIM in  the near term is  to search for  UV (\\ion{O}{6}, \\ion{Ne}{8}) and X-ray  (\\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9})  absorption lines due to  WHIM gas  in the  spectra of  background QSOs/AGNs  (Tripp  et al. 2000, 2001; Savage et al.   2002,2005; Nicastro et al.  2002; Bergeron et al.  2002; Richter et al.  2004; Sembach et al.  2004; Prochaska et al. 2004; Danforth  \\& Shull 2005).  While absorption  lines provide a sensitive and  powerful probe of  the WHIM, the pencil-beam  nature of the measurement along a sight  line provides little information on the context  of  the absorption,  e.g.,  whether  the  lines arise  in  an individual galaxy disk/halo, a  galaxy group, or lower-density regions of a large-scale filament or void.  Thus, to understand the nature of highly ionised absorbers at low redshifts, several groups are pursuing deep galaxy redshift surveys and observations of QSOs behind well-defined galaxy groups or clusters.  For example, to study gas located in large-scale filaments, Bregman et al. (2004) have searched for absorption lines indicative of the WHIM in regions between galaxy clusters/superclusters and have identified some candidates.  In this paper, we carry out a similar search as part of a broader program that combines a large {\\it HST} survey of low$-z$ \\ion{O}{6} absorption systems observed on sight lines to low$-z$ quasars (Tripp et al. 2004) and a ground based survey to measure the redshifts and properties of the galaxies foreground to the background QSOs.  The ground based survey is done in two steps: first, multi-band (U,B,V,R and I) imagery is obtained to identify the galaxies and to estimate their photometric redshifts. Then, spectroscopic redshifts are obtained for the galaxies that are potentially (according to the photometric redhshifts) at lower redshift that the background object.  As part of the large {\\it HST} survey, we have observed the quasar HS0624+6907 ($z_{\\rm QSO}$ = 0.3700) with the E140M echelle mode of the Space Telescope Imaging Spectrograph (STIS) on board the {\\it Hubble Space Telescope}. We have also obtained multiband images and spectroscopic redshifts of galaxies in the \\hs0624 field.  The sight line to \\hs0624 passes by several foreground Abell clusters (\\S~\\ref{sec:abell_clusters}) and provides an opportunity to search for gas in large-scale filaments.  We shall show that gas (absorption systems) and galaxies are detected at the redshifts of the structures delineated by the Abell clusters in this direction.  While the absorbing gas is intergalactic, and it is likely that we are probing gas in cosmic web filaments, the properties of these absorbers are surprising.  Instead of low-metallicity WHIM gas, we predominantly find cool, photoionised, and high-metallicity gas in these large-scale structures.  This paper is organized as follows. The observations and data reduction procedures are described in \\S2, including {\\it HST}/STIS and {\\it Far Ultraviolet Spectroscopic Explorer} (\\FUSE) observations as well as ground-based imaging and galaxy redshift measurements. In \\S3, we present information on the foreground environments probed by the \\hs0624 sight line, derived from the literature on Abell clusters and from our new galaxy redshift survey.  The absorption-line measurement methods are described in \\S4, and we investigate the physical state and metallicity of the absorbers in \\S5. Section 6 reviews the properties of the full sample of Ly$\\alpha$ lines derived from the STIS spectrum with emphasis on the search for broad Ly$\\alpha$ lines. Section 7 discusses the implications of this study, and we summarize our conclusions in \\S8. Throughout this paper, we use the following cosmological parameters: $h_{70}=H_0/70$~\\kms, $\\Omega_m=0.3$ and $\\Lambda_o=0.7$. ",
        "conclusions": "\\label{sec:discussion}  We have acquired detailed information about the abundances, physical conditions, and galaxy proximity of absorption systems in the direction of \\hs0624.  What are the implications of these measurements for broader questions of galaxy evolution and cosmology?  The processes that add gas to galaxies (e.g. accretion) and remove gas from galaxies (e.g., winds, dynamical stripping) can have profound effects on galaxy evolution, and the ``feedback'' of matter and energy from galaxies into the IGM is now believed to play an important role in shaping structures that subsequently grow out of the IGM (Voit G. M., 2005).  The quantity and implications of the $10^{5} - 10^{7}$ K WHIM gas is a topic of particular interest currently.  The galaxies and absorption systems in the direction of \\hs0624, particularly the galaxy group and Ly$\\alpha$ complex at $z$ = \\zva , have some interesting, and perhaps surprising, implications regarding these questions, which we now discuss.  \\begin{figure} \\begin{center} \\resizebox{1.0\\hsize}{!}{\\includegraphics{figure_16.ps}} \\caption{Images of galaxies found in the $z$ = \\zva\\ group in the field of \\hs0624, recorded in the $R$ band with the KPNO 4m MOSA camera. Each box in the montage spans $50''\\times 50''$, and the galaxy name from Table~\\ref{tab:spec_red} is listed in the lower left corner. The galaxy is in the center of each panel. SE6 and SE7 are in close proximity; SE6 is the more extended spiral galaxy.}\\label{fig:galpics} \\end{center} \\end{figure}  ROSAT observations of diffuse X-ray emission have established that galaxy groups that are dominated by early-type galaxies often contain diffuse, hot intragroup gas (Mulchaey 2000, and references therein). Based on the observed relation between intragroup gas temperature and velocity dispersion $\\sigma$ in X-ray bright groups ($T \\propto \\sigma ^{2}$) and the fact that spiral-rich groups have lower velocity dispersions than elliptical-rich groups, Mulchaey et al. (1996) have hypothesized that spiral-rich groups might have somewhat cooler intragroup media that could give rise to QSO absorption lines at WHIM temperatures (e.g., \\ion{O}{6}).  However, the galaxy group at $z$ = \\zva\\ appears to have properties that are not consistent with the elliptical-rich groups detected with ROSAT nor with the idea that spiral-rich groups contain warm-hot intragroup gas. It is unclear if the galaxy group at $z$ = \\zva\\ is a spiral-rich group.  Figure~\\ref{fig:galpics} shows $R$ images from the MOSA data of the 10 galaxies that we have found in this group. Most of the galaxies in the group show evidence of disks and bulges (both in the direct images and in radial brightness profiles). We find from visual inspection that 4-5 of the 10 galaxies have indications of spiral structure (SW3, SE1, SE8, SE6, and possibly SE4). However, the remaining galaxies could be early-type S0 galaxies, and therefore the early-type fraction might be comparable to groups that show diffuse X-ray emission (see, e.g, Figure 7 in Zabludoff \\& Mulchaey 1998). The NE3, SE13, and SE5 galaxies, which morphologically appear to be early-type galaxies, have colours and magnitudes consistent with the ``red sequence'' colour-magnitude relation observed in clusters (e.g., Bower, Lucey, \\& Ellis 1992; McIntosh et al. 2005); these galaxies are likely S0s (the other 7 galaxies have blue colours characteristic of late types). The velocity dispersion of the group at $z$ = \\zva , albeit uncertain, is more comparable to those of elliptical-dominated groups than spiral-rich groups (Mulchaey et al. 1996; Zabludoff \\& Mulchaey 1998).  Regardless of whether the group is elliptical- or spiral-rich, it is surprising that we find a large number of cool, photoionised clouds in the intragroup medium (\\S \\ref{ss:z064}).  In the hot intragroup medium of an elliptical-rich group, \\ion{H}{1} lines should be extremely broad and weak, but instead we find strong, narrow lines (see Figure~\\ref{fig:hi00635}).  Even in the cooler gas predicted to be found in late-type dominated groups, the \\ion{H}{1} lines should be broader. We could entertain models of cooling intragroup gas, but in such models \\ion{O}{6} is expected to be stronger.  Likewise, if the intragroup gas is a multiphase medium with cooler clouds (which cause the \\ion{H}{1} absorption lines) embedded in a hotter phase, then we might expect to detect \\ion{O}{6} from the interface between the phases (Fox et al. 2005), unless conduction is somehow suppressed.  The lack of evidence of hot gas leads us to question whether this group is a bound, virialized system.  An alternative possibility is that our sight line passes along the long axis of a large-scale filamentary structure in the cosmic web. In this case, the projection of the galaxies and Ly$\\alpha$ clouds along the sight line could give a false impression of a group in which hot gas would be expected. However, in cosmological simulations of large-scale filaments, WHIM gas is expected to be widespread at the present epoch, even in modest-overdensity regions (see, e.g., Figure 4 in Cen \\& Ostriker 1999b), so it is interesting that we find a substantial number of cool clouds at $z \\approx$ \\zva , somewhat contrary to theoretical expectations. As noted above, our data do not preclude the presence of WHIM gas at $z \\approx$ \\zva , but we find no clear evidence of it. Other sight lines show similar clusters of Ly$\\alpha$ lines, e.g., the Ly$\\alpha$ complex at $z_{abs} \\approx$ 0.057 toward PKS2155-304 (Shull et al. 1998; Shull, Tumlinson, \\& Giroux 2003) or the Ly$\\alpha$ lines at $z_{abs} \\approx$ 0.121 toward H1821+643 (Tripp et al. 2001).  However, unlike the \\hs0624 Ly$\\alpha$ complex, the PKS2155-304 and H1821+643 examples both show evidence of warm-hot intragroup gas. To test whether the observations and simulations are in accord, it would be useful to assess the frequency and physical properties of these Ly$\\alpha$ complexes in cosmological simulations for comparison with the observations.  It is also interesting that the two systems for which we have obtained abundance constraints (at $z$ = \\zva\\ and \\zvb ) both indicate relatively high metallicities, but both of these systems are at least 100 kpc away (in projection) from the nearest known galaxy. This naturally raises a question: how did gas that is so far from a galaxy attain such a high metallicity?  The gas could have been driven out of a galaxy by a galactic wind; some wind models predict that the outflowing material will have a high metallicity (Mac Low \\& Ferrara 1999), The difficulty with this interpretation is that winds from nearby galaxies are usually observed to contain substantial amounts of hot gas (e.g., Strickland et al. 2004), which seems to be inconsistent with the absorption line properties as we have discussed.  A more likely explanation is that the high-metallicity gas we have detected in absorption has been tidally stripped out of one of the nearby galaxies. There are indications that tidal stripping could be a more gentle process for removing gas from galaxies, and a tidally stripped origin can therefore more easily accommodate the observed low-ionisation state of the gas. For example, in the direction of NGC3783, the Galactic high-velocity cloud (HVC) at $v_{\\odot} =$ 247 \\kms\\ is now recognized to be tidally stripped debris from the SMC. While this tidally stripped material shows a wide array of low-ionisation absorption lines, it has little or no associated high-ion absorption (Lu et al. 1994; Sembach et al. 2001a).  Moreover, the tidally-stripped HVC contains molecular hydrogen, which Sembach et al. argue formed in the SMC and survived the rigors of tidal stripping (as opposed to forming in situ in the stream). Both the absence of high ions and the survival of H$_{2}$ suggest that the stripping process did not substantially ionise and heat this HVC. Several galaxies are close enough to the \\hs0624 sight line to be plausible sources of the gas in a tidal stripping scenario. One of the nearby galaxies, SE1, has a distorted spiral morphology.  This galaxy is a plausible source of tidally stripped matter."
    },
    "0512/astro-ph0512213_arXiv.txt": {
        "abstract": "We argue that combined observations of galaxy rotation curves and gravitational lensing not only allow the deduction of a galaxy's mass profile, but also yield information about the pressure in the galactic fluid. We quantify this statement by enhancing the standard formalism for rotation curve and lensing measurements to a first post-Newtonian approximation. This enhanced formalism is compatible with currently employed and established data analysis techniques, and can in principle be used to reinterpret existing data in a more general context. The resulting density and pressure profiles from this new approach can be used to constrain the equation of state of the galactic fluid, and therefore might shed new light on the persistent question of the nature of dark matter. ",
        "introduction": "One of the most compelling issues of modern astrophysics is the open question concerning the nature of the dark matter which dominates the gravitational field of individual galaxies and galaxy clusters.  [See for instance \\citet{Persic:1996}, \\citet{Borriello:2001}, and \\citet{Salucci:2003}.] While the current consensus in the astrophysics community is to advocate the cold dark matter (CDM) paradigm, no \\textit{direct} observations of the equation of state have been carried out to confirm this widely adopted assumption. Efforts to confirm this assumption include attempts to detect elementary particles that have been suggested as cold dark matter candidates. However, experiments that aim (for instance) to detect massive axions with Earth-based detectors \\citep[][\\S 22.2.2]{Eidelman:2004} do not yet yield a positive result.  A different approach to analysing the nature of dark matter has been suggested by \\citet{Bharadwaj:2003} who first proposed that combined measurements of rotation curves and gravitational lensing could be used to determine the equation of state of the galactic fluid. Whereas their analysis made particular assumptions on the form of the rotation curve, and is restricted to a certain type of equation of state, herein we provide a general formalism that allows us to deduce the density and pressure profiles without any prior assumptions about their shape or the equation of state.  Analytic galaxy halo models that predict a significant amount of pressure or tension in the dark matter fluid include ``string fluid'' \\citep{Soleng:1993}, or some variations of scalar field dark matter (SFDM). See for instance \\citet{Schunck:1999}, \\citet*{Matos:1999et}, \\citet*{Matos:2000ki}, \\citet{Peebles:2000}, and \\citet*{Arbey:2003}. Our method provides a means of observing, or at least constraining, the pressure distribution in a galactic halo. Therefore it is in principle able to give evidence for or against specific proposed dark matter candidates.  The key point is that in general relativity, density and pressure \\textit{both} contribute to generating the gravitational field \\emph{separately}.  Furthermore, the perception of this gravitational field depends on the velocity of probe particles. These effects become especially important when one compares rotation curve and gravitational lensing measurements, where the probe particles are fundamentally different: interstellar gas or stars at subluminal velocities for rotation curves, and photons which travel at the speed of light for lensing measurements. Our formalism accounts for these crucial differences between the probe particles, and relates observations of both kinds to the the density and pressure profile of the host galaxy. Although we (mainly) consider static spherically symmetric galaxies in a first post-Newtonian approximation, the basic concept is fundamental and can be extended to more general systems with less symmetry. A suitable framework for considering most exotic weak gravity scenarios is provided by the effective refractive index tensor, as introduced by \\citet*{Boonserm:2005}.  The present approach might also help to shed some light on prevailing problems that arise when combining rotation curve and lensing observations. For example, an unresolved issue exists when measuring the Hubble constant from the time delay between gravitationally lensed images: Using the standard models for matter distribution in the lens galaxy, the resulting Hubble constant is either too low compared to its value from other observations, or the dark matter halo must be excluded from the galaxy model to obtain the commonly accepted value of $H_0$ \\citep{Kochanek:2004}. A possible explanation of this trend might lie in a disregarded pressure component of the dark matter halo.  We organise this article in the following manner: First we introduce the minimal necessary framework of general relativity concepts, and point out the important conditions required to obtain the Newtonian gravity limit.  Next, we elaborate on the post-Newtonian extension of the currently employed rotation curve and gravitational lensing formalisms.  Consequently, we show how to combine rotation curve and lensing measurements to make inferences about the density and pressure profile of the observed galaxy. We then examine how noticeable the effects of non-negligible pressure could be in the measurements. Lastly, we discuss how the formalism adapts to non-spherically symmetric galaxies and comment on the current observational situation and issues arising with the new formalism. ",
        "conclusions": "We have argued that the standard formalism of rotation curve measurements and gravitational lensing make an \\textit{a priori} Newtonian assumption that is based on the CDM paradigm. We introduce a post-Newtonian formalism that does not rely on such an assumption, and furthermore allows one to deduce the density- and pressure-profiles in a general relativistic framework. In this framework, rotation curve measurements provide a pseudo-mass profile $m_\\rmn{RC}(r)$ and gravitational lensing observations yield a different pseudo-mass profile $m_\\rmn{lens}(r)$. Combining both pseudo-masses allows one to draw conclusions about the density- and pressure profiles\\footnote{These formulae are given in SI units, hence the factor of $c^2$.} in the lensing galaxy, \\begin{eqnarray} \\rho(r) &=& \\frac{1}{4\\upi\\,r^2} \\, \\left[ 2\\,m_\\rmn{lens}'(r) - m_\\rmn{RC}'(r) \\right] \\, , \\\\ p_r(r)+2p_t(r) &\\approx& \\frac{2\\,c^2}{4\\upi\\,r^2} \\, \\left[m_\\rmn{RC}'(r)-m_\\rmn{lens}'(r) \\right] \\, . \\end{eqnarray} In the case of absent or negligible pressure, this could be used to observationally confirm the CDM paradigm of a pressureless galactic fluid. Conversely, if significant pressure is detected, a decomposition of the galaxy morphology would allow new insight into the equation of state of dark matter.  For instance, detailed observation of the recently discovered closest known strong lensing galaxy ESO 325-G004 \\citep{Smith:2005} could provide satisfactory data to allow the decomposition of density and pressure of the galactic fluid, as outlined in this article. The system consists of an isolated lensing galaxy at redshift $z\\approx 0.035$ with an effective radius of $R_\\rmn{eff}=12\\,\\farcs 5$ and arc-shaped images of the background object at $R\\approx 3\\arcsec$, and possible arc candidates at $R\\approx 9\\arcsec$. \\citet{Smith:2005} intend to collect more detailed data that hopefully will include extended stellar dynamics and hence, allow for a direct comparison of the rotation curve and lensing data, if the arc candidates at $R\\approx 9\\arcsec$ turn out to contribute to the measurements.  Since the formalism presented is based on a first-order weak field approximation, we suggest that to confirm the findings, one should re-insert the obtained density and pressure profiles into the metric \\eref{ss_metric}. The actual observed quantities can then be extracted numerically for comparison from the exact field equations \\eref{Einstein_tt}--\\eref{Einstein_transverse} and the geodesic equations.  Finally, even though data might not yet be available to constrain the dark matter equation of state noticeably, one should note that the possibility of non-negligible pressure in the galactic fluid introduces a new free parameter into the analysis of combined rotation curve and lensing observations."
    },
    "0512/astro-ph0512025_arXiv.txt": {
        "abstract": "Unified schemes of active galactic nuclei (AGN) require an obscuring dusty torus around the central engine. Torus sizes of hundreds of parsecs were deduced from early theoretical modeling efforts, but high-resolution IR observations now show that the torus size is no more than a few parsecs. This conflict is resolved when the clumpy nature of the torus is taken into account. The compact torus may be best understood when identified with the dusty, optically thick region of the wind coming off the central accretion disk. ",
        "introduction": "\\label{sec:introduction}  The great diversity of AGN classes has been explained by a single unified scheme (e.g.\\ Antonucci 1993; Urry \\& Padovani 1995). The nuclear activity is powered by a super\\-massive (\\about\\E6--\\E{10} \\Mo) black hole and its accretion disk. This central engine is surrounded by dusty clouds, which are individually optically thick, in a toroidal structure (Krolik \\& Begelman 1988). Much of the observed diversity is simply the result of viewing this axi\\-symmetric geometry from different angles. The clumpy torus provides anisotropic obscuration of the central region so that sources viewed face-on are recognized as type 1 objects, those observed edge-on are type 2. The fraction of the sky obscured by the torus determines the relative numbers of type 1 and 2 sources. From the statistics of Seyfert galaxies, Schmitt et al (2001) find that the torus height and radius obey $H/R$ \\about\\ 1. In the ubiquitous sketch by Urry \\& Padovani (1995), the torus is depicted as a large doughnut-like object, presumably populated by molecular clouds accreted from the galaxy. Gravity controls the orbital motions of the clouds, but the origin of cloud vertical motions capable of sustaining the ``doughnut\" in a steady-state with $H \\sim R$ was recognized as a problem by Krolik \\& Begelman (1988). This problem has eluded solution to this date. ",
        "conclusions": ""
    },
    "0512/astro-ph0512031_arXiv.txt": {
        "abstract": "We demonstrate that gravitationally lensed quasars are easily recognized using image subtraction methods as time variable sources that are spatially extended.  For Galactic latitudes $|b|\\gtorder 20^\\circ$, lensed quasars dominate the population of spatially extended variable sources, although there is some contamination from variable star pairs, variable star-quasar pairs and binary quasars that can be easily controlled using other information in the survey such as the object light curves and colors.  This will allow planned large-scale synoptic surveys to find lensed quasars almost down to their detection limits without the need for extensive follow-up observations. ",
        "introduction": "\\label{sec:introduction}  In theory, gravitational lenses can be used to address astrophysical problems such as the cosmological model, the structure and evolution of galaxies, and the structure of quasar accretion disks (see the reviews by Kochanek~(\\citeyear{Kochanek2004saas}) of strong lensing and Wambsganss (\\citeyear{Wambsganss2004saas}) of microlensing). One of the main challenges in using lenses for any of these applications is discovering large numbers of lenses efficiently (see the review of lens surveys in Kochanek~(\\citeyear{Kochanek2004saas})).  Most known lenses have been found either from optical imaging surveys of known quasars (see  Pindor et al.~\\citeyear{Pindor2003p2340} for a recent study), radio imaging surveys of flat-spectrum radio sources (see Browne et al.~\\citeyear{Browne2003p13}), or searches for anomalous, higher redshift emission lines in galaxy spectra (see Bolton et al.~\\citeyear{Bolton2005}).  Imaging surveys of all radio sources (Burke~\\citeyear{Burke1990}) proved difficult because of the confusing array of structures observed for steep spectrum radio sources.  Haarsma et al.~(\\citeyear{Haarsma2005}) proposed improving the efficiency of searches for lensed steep-spectrum sources by looking for radio lobes with optical counterparts, but the approach is limited by the resolution and sensitivity of existing all-sky radio surveys.  None of these methods is easily applied to the next generation of large scale imaging surveys such as the SDSS Supernova Survey (Sako et al.~\\citeyear{Sako2005}), the Dark Energy Survey (DES, Abbott et al.~\\citeyear{Abbott2005}), Pan-STARRS (Kaiser~\\citeyear{Kaiser2004}) and the Large Synoptic Survey Telescope (LSST, Tyson et al.~\\citeyear{Tyson2003}).    One possibility is to use a combination of color and morphology to identify quasar lens candidates (Morgan et al.~\\citeyear{Morgan2004}).  This strategy can be effective as long as emission (or absorption) by the lens galaxy does not significantly change the color of the system from that of the quasars, which restricts its applicability to systems in which the quasar images are significantly brighter than the lens galaxy.  A new feature of all these projects, however, is that they are synoptic surveys which obtain light curves for variable sources.  Pindor (\\citeyear{Pindor2005p649}) suggested that the synoptic data could be used to find lenses by cross-correlating the light curves of closely separated sources to search for the time delays present in the lensed systems.  This approach may be problematic as a search method because it requires the automated extraction of light curves for the individual lensed images, some of which may also be distorted by the effects of microlensing.  However, it will be an essential component of verifying lens candidates in the synoptic surveys.  In this paper we introduce a far simpler strategy.  Unlike almost any other source, lensed quasars are ``extended'' variable sources because the variable flux is spread out over the scale of the image separations.  As we discuss in \\S2, restricting the search to extended variable sources is an extraordinarily powerful means of eliminating sources other than gravitational lenses. In \\S3 we demonstrate the method using data we have been acquiring to measure time delays and microlensing variability in known lensed quasars (Kochanek et al.~\\citeyear{Kochanek2005}).  We summarize our proposed search in \\S4. ",
        "conclusions": "\\label{sec:summary}  Almost all new, large scale imaging surveys will be synoptic surveys that monitor the time variability of sources in the survey area.  By using difference imaging to search for extended variable sources, these surveys can easily identify lensed quasars because almost all other variable sources are either point sources or orbital tracks created by solar system objects.  We estimate that gravitational lenses are the most common extended variable sources for Galactic latitudes $|b| \\gtorder 20^\\circ$, with modest contamination from pairs of variable stars, variable star/quasar pairs and binary quasars.  Limiting the search to extended variable sources reduces the number of non-lens background objects by more than $10^5$. Thus, it should be relatively straight forward for SDSS, Pan-STARRS, DES or LSST to identify the lensed quasars in their respective variability survey areas.  For LSST, this should amount to roughly $10^3$ lensed quasars to V$<23$~mag.  Note that the criterion of being an extended variable source can be combined with other criteria to further reduce the rate of false positives based on other information available from the same survey.  For example, the quasars will have slowly varying aperiodic light curves, while many stars will show more rapid variability or periodic variability.  Where the source is resolved, the light curves of lensed quasars should be similar and can be cross-correlated to measure a time delay and verify that the source is a lens.  Note, however, that our discovery method depends only on the existence of variability rather than the measureability of the delay.  The colors of stars and quasars are different at most redshifts (e.g. Richards et al.~\\citeyear{Richards2004}), and the colors of lensed images should be similar, up to concerns about differential extinction in the lens (e.g. Falco et al.~\\citeyear{Falco1999}).   Finally, in the average image it should be possible to detect a lens galaxy, potentially using difference imaging methods to accurately subtract the quasar contribution.  In general, the background of non-lens sources can be so greatly suppressed that we suspect the only significant issue for candidate selection will be systematic errors in identifying extended variable sources that are presently difficult to quantify."
    },
    "0512/astro-ph0512207_arXiv.txt": {
        "abstract": "We analyze a 19-night photometric search for transiting extrasolar planets in the open cluster NGC 1245.  An automated transit search algorithm with quantitative selection criteria finds six transit candidates; none are bona fide planetary transits. We characterize the survey detection probability via Monte Carlo injection and recovery of realistic limb-darkened transits. We use this to derive upper limits on the fraction of cluster members with close-in Jupiter-radii, $R_{J}$, companions.   The survey sample contains $\\sim$870 cluster members, and we calculate 95\\% confidence upper limits on the fraction of these stars with planets by assuming the planets have an even logarithmic distribution in semimajor axis over the Hot Jupiter (HJ - 3.0$<P/\\days<$9.0), Very Hot Jupiter (VHJ - 1.0$<P/\\days<$3.0), and an as of yet undetected Extremely Hot Jupiter (EHJ - $P_{\\rm Roche}<P<$1.0 day) period ranges.  For 1.5 $R_{J}$ companions we limit the fraction of cluster members with companions to $<$1.5\\%, $<$6.4\\%, and $<$52\\% for EHJ, VHJ, and HJ companions, respectively.  For 1.0 $R_{J}$ companions, we find $<$2.3\\% and $<$15\\% have EHJ and VHJ companions, respectively.  We do not reach the sensitivity to place any meaningful constraints on 1.0 $R_{J}$ HJ.  From a careful analysis of the random and systematic errors of the calculation, we show that the derived upper limits contain a $\\pm^{13\\%}_{7\\%}$ relative error.  For similar photometric noise and weather properties as this survey, observing NGC 1245 twice as long results in a tighter constraint on HJ companions than observing an additional cluster of similar richness as NGC 1245 for the same length of time as this survey.  If 1\\% of stars have 1.5 $R_{J}$ HJ companions, we expect to detect one planet for every 5000 dwarf stars observed for a month.  To reach a $\\sim$ 2\\% upper limit on the fraction of stars with 1.5 $R_{J}$ companions in the 3.0$<P/\\days<$9.0 range, we conclude a total sample size of $\\sim 7400$ dwarf stars observed for at least a month will be needed.  Results for 1.0 $R_{J}$ companions, without substantial improvement in the photometric precision, will require a small factor larger sample size. ",
        "introduction": "Extrasolar planet detections and analysis of the non-detections further our knowledge of the planet formation process and contribute to an empirical determination of the typical planetary system.  These empirical constraints will eventually decide the ubiquity or rarity of planetary bodies in the Universe.  A variety of techniques exist to detect extrasolar planets \\citep{PER00}, and there are currently 168 extrasolar planet discoveries\\footnote{http://www.obspm.fr/encycl/catalog.html}.  Most of the extrasolar planets have been discovered using the radial velocity technique.  Radial velocity detections indicate that $1.2\\%\\pm 0.3\\%$ of FGK main-sequence stars in the solar neighborhood have a ``Hot'' Jupiter-mass planet (HJ) orbiting within 0.1 AU \\citep{MAR05}.  At this small separation from the central star, the high temperatures and low disk column densities prevent in situ formation of HJ planets \\citep{BOD00}.  Several mechanisms exist to exchange angular momentum between the protoplanet and natal disk, enabling the protoplanet to migrate from a more likely formation separation (several AU) to within 0.1 AU \\citep{TER00}.  Due to tidal circularization, HJ have nearly circular orbits with an observed median eccentricity, $<e>=0.07$, whereas planets with larger separations have a median eccentricity of $<e>=0.25$.  In addition to the detection statistics and planet properties, the extrasolar planet detections indicate several physical relationships between the stellar host properties and the frequency of extrasolar planets.  The most striking of these is that the probability for hosting an extrasolar planet increases rapidly with stellar metal abundance, consistent with $P\\propto N_{\\rm Fe}^{2}$ \\citep{FIS05}.  The frequency of planets may also depend on the stellar mass. \\citet{BUT04} and \\citet{BON05} point out that there exists a deficit of $M_{J}$ planets orbiting M dwarf stars.  However, the increasing number of short-period Neptune-mass planets being discovered around M dwarfs suggests that the overall frequency of planets (of all masses) orbiting M dwarfs may be similar to FGK dwarfs, but the typical planet mass is less, thereby escaping detection given the detection limitations of the current radial velocity surveys \\citep{BON05}. Additionally, none of the M dwarfs harboring planets are metal rich \\citet{BON05}.  A coherent theory of planet formation and survival requires not only reproducing the physical properties of the planets, but reproducing any trends in the physical properties on the host environment. Despite the knowledge and constraints on extrasolar planets that radial velocity surveys provide, radial velocity surveys have their limitations.  The high resolution spectroscopic requirements of the radial velocity technique limit its use to the solar neighborhood and orbital periods equivalent to the lifetime of the survey.  A full consensus of the planetary formation process requires relying on additional techniques to detect extrasolar planets in a larger variety of conditions prevalent in the Universe.  For instance, microlensing surveys are sensitive to extrasolar planets orbiting stars in the Galactic disk and bulge with distances of many kpc away (\\citealt{MAO91,GOU92}).  Two objects consistent with Jupiter-mass companions have been detected via the microlensing technique \\citep{BON04,UDA05}.  Additional information is obtained from studying the microlensing events that did not result in extrasolar planet detections.  Microlensing surveys limit the fraction of M dwarfs in the Galactic bulge with Jupiter-mass companions orbiting between 1.5 AU to 4 AU to $<33\\%$ (\\citealt{ALB01,GAU02}).  Although limited to the solar neighborhood, attempts to directly image extrasolar planets are sensitive to planets with semimajor axis beyond 20 AU.  The light from the parent star limits detecting planets interior to the seeing disk.  Adaptive optics observations of young ($\\sim 1$ Myr) stars provide the best opportunity to directly image extrasolar planets since the young planets are still relatively bright while undergoing a rapid, cooling contraction.  Although the interpretation relies on theoretical modeling of these complex planetary objects, three sources in nearby star forming regions have been detected whose broad-band colors and spectra are consistent with those expected from 1-42 Jupiter-mass objects \\citep{NEU05,CHA05A,CHA05B}.  The contrast ratios necessary for extrasolar planet detection are difficult to reach, and results for detecting higher mass brown dwarfs are more complete.  An analysis of the Cornell High-Order Adaptive Optics Survey (CHAOS) derives a brown dwarf companion upper limit of 10\\% orbiting between 25 and 100 AU of the parent star \\citep{CAR05}.  \\citet{MCC04} estimate $1\\%\\pm 1\\%$ of G,K, and M stars have brown dwarf companions orbiting between 75 and 300 AU, but this estimate may not account for the full range of orbital inclination and eccentricities possible \\citep{CAR05}.  At greater separations, $>$ 1000 AU, brown dwarf companions to F-M0 main-sequence stars appear to be as common as stellar companions \\citep{GIZ01}.  After the radial velocity technique, the transit technique has had the most success in detecting extrasolar planets \\citep{KON05}.  The transit technique can detect $R_{J}$ transits in any stellar environment where $\\lesssim$1\\% photometry is possible.  Thus, it provides the possibility of detecting extrasolar planets in the full range of stellar conditions present in the Galaxy: the Solar neighborhood, the thin and thick disk, open clusters, the halo, the bulge, and globular clusters are all potential targets for transit surveys.  A major advantage of the transit technique is the current large-format mosaic CCD imagers which provide multiplexed photometric measurements with sufficient accuracy across the entire field of view.  The first extrasolar planet detections via the transit technique began with the candidate list provided by the OGLE collaboration \\citep{UDA02}.  However, confirmation of the transiting extrasolar planet candidates requires radial velocity observations.  Due to the well known equation-of-state competition between electron degeneracy and ionic Coulomb pressure, the radius of an object becomes insensitive to mass across the entire range from below $M_{J}$ to the hydrogen-burning limit \\citep{CHA00}.  Thus, objects revealing a $R_{J}$ companion via transits may actually have a brown-dwarf mass companion when followed up with radial velocities. This degeneracy is best illustrated by the planet-sized brown dwarf companion to OGLE-TR-122 \\citep{PON05}.  The first radial-velocity confirmations of planets discovered by transits \\citep{KON03,BOU04} provided a first glimpse at a population of massive, very close-in planets with $P<$ 3 days and $M_{p}>M_{J}$ (``Very Hot Jupiters'' - VHJ) that had not been seen by radial velocity surveys. \\citet{GAU05A} demonstrated that, after accounting for the strong sensitivity of the transit surveys to the period of the planets, the transit detections were likely consistent with the results from the radial velocity surveys, implying that VHJs were intrinsically very rare.  Subsequently, in a metallicity-biased radial velocity survey, \\citet{BOU05B} discovered a VHJ with $P=2.2$ day around the bright star HD189733 that also has observable transits.  Despite the dependence of transit detections on radial velocity confirmation, radial velocity detections alone only result in a lower limit on the planetary mass, and thus do not give a complete picture of planet formation.  The mass, radius information directly constrains the theoretical models, whereas either parameter alone does little to further constrain the important physical processes that shape the planet properties \\citep{GUI05}.  For example, the mass-radius relation for extrasolar planets can constrain the size of the rocky core present (e.g., \\citealt{LAU05}).  Also, the planet transiting across the face of its parent star provides the exciting potential to probe the planetary atmospheric absorption lines against the stellar spectral features \\citep{CHA02,DEM05,NAR05}.  Or, in the opposite case, emission from the planetary atmosphere can be detected when the planet orbits behind the parent star \\citep{CHA05,DEM05B}.  Despite these exciting results, the transit technique is significantly hindered by the restricted geometrical alignment necessary for a transit to occur. As a result, a transit survey necessarily contains at least an order of magnitude more non-detections than detections.  In addition, null results themselves can provide important constraints.  For example, the null result in the globular cluster 47 Tucanae adds important empirical constraints to the trend of increasing probability of having a planetary companion with increasing metallicity \\citep{GIL00,SAN04}. Thus, understanding the sensitivity of a given transit survey, i.e.\\ the expected rate of detections and non-detections, takes on increased importance. Several studies have taken steps toward sophisticated Monte Carlo calculations to quantify detection probabilities in a transit survey \\citep{GIL00,WEL05,MOC05,HID05,HOO05}. Unfortunately, these studies do not fully characterize the sources of error and systematics present in their analysis, and therefore the reliability of their conclusions is unknown.  Furthermore, essentially all of the previous studies have either (1) not accurately determined the number of dwarf main-sequence stars in their sample, or (2) made simplifying assumptions which may lead to misestimated detection probabilities, or (3) contained serious conceptual errors in the procedure with which they have determined detection probabilities, or (4) some combination of the above.  As a specific and important example, most studies do not apply identical selection criteria when searching for transits amongst the observed light curves and when recovering injected transits as part of determining the survey sensitivity.  Removal of false-positive transit candidates arising from systematic errors in the light curve has typically involved subjective visual inspections, and these subjective criteria have not been applied to the recovery of injected transits when determining the survey sensitivity.  This is statistically incorrect, and can in principle lead to overestimating the survey sensitivity. Even if identical selection criteria are applied to the original transit search and in determining the survey sensitivity, some surveys do not apply conservative enough selections to fully eliminate false-positive transit detections.  In this paper, we address these shortcomings of previous studies in our analysis of a 19-night photometric search for transiting extrasolar planets in the open cluster NGC 1245.  An automated transit search algorithm with quantitative selection criteria finds six transit candidates; none are bona fide planetary transits. We describe our Monte Carlo calculation to robustly determine the sensitivity of our survey, and use this to derive upper limits on the fraction of cluster members with close-in, Jupiter-radii, $R_{J}$, companions.  Leading up to the process of calculating the upper limit, we develop several new analysis techniques.  First, we develop a differential photometry method that automatically selects comparison stars to reduce the systematic errors that can mimic a transit signal.  In addition, we formulate quantitative transit selection criteria, which completely eliminate false positives due to systematic light-curve variability without human intervention.  We characterize the survey detection probability via Monte Carlo injection and boxcar recovery of transits. Distributing the Monte Carlo calculation to multiple processors enables rapid calculation of the transit detection probability for a large number of stars.  The techniques developed here enable combining results from transit surveys in a statistically meaningful way.  This work is part of the Survey for Transiting Extrasolar Planets in Stellar Systems (STEPSS). The project concentrates on stellar clusters since they provide a large sample of stars of homogeneous metallicity, age, and distance \\citep{BUR03,BUR04}.  Overall, the project's goal is to assess the frequency of close-in extrasolar planets around main-sequence stars in several open clusters.  By concentrating on main-sequence stars in open clusters of known (and varied) age, metallicity, and stellar density, we will gain insight into how these various properties affect planet formation, migration, and survival.  The survey characteristics and the photometric procedure are given in \\S\\ref{OBS}.  We explain the automated algorithm to calculate the differential light curves and describe the light curve noise properties in \\S\\ref{LC}.  In \\S\\ref{trandetect} we describe our implementation of the box-fitting least squares (BLS) method \\citep{KOV02} for transit detection.  In \\S\\ref{sec:selcrit} we present a thorough discussion of the quantitative selection criteria for transit detection, followed by a discussion of the objects with sources of astrophysical variability that meet the selection criteria in \\S\\ref{trncands}.  We outline the Monte Carlo calculation for determining the detection probability of the survey in \\S\\ref{effcalc}.  We present upper limits for a variety of companion radii and orbital periods in \\S\\ref{results}.  A discussion of the random and systematic errors present in the technique is given in \\S\\ref{uplimiterrsec}.  We compare the final results of this study to our expected detection rate before the survey began and discuss the observations necessary to reach sensitivities similar to radial velocity detection rates in \\S\\ref{discussion}. Finally, \\S\\ref{conclusion} briefly summarizes this work. ",
        "conclusions": "In this study we complete the analysis of a 19-night search for transiting extrasolar planets orbiting members of the open cluster NGC 1245.  An automated transit search algorithm with quantitative selection criteria finds six transit candidates; none are bona fide planetary transits.  Thus, this work also details the procedure for analyzing the null-result transit search in order to determine an upper limit on the fraction of stars in the cluster harboring close-in $R_{J}$ companions.  In addition, we outline a new differential photometry technique that reduces the level of systematic errors in the light curve.  A reliable upper limit requires quantifiable transit selection criteria that do not rely on visual, qualitative judgments of the significance of a transit.  Thus, we develop completely quantitative selection criteria that enable us to calculate the detection probability of the survey via Monte Carlo techniques.  We inject realistic limb-darkened transits in the light curves and attempt their recovery.  For each star we inject 100,000 transits at a variety of semimajor axes, orbital inclination angles, and transit phases, to fully map the detection probability for 2700 light curves consistent with cluster membership based on their position in the CMD. After characterizing the field contamination, we conclude the sample contains $\\sim$870 cluster members.  When calculating a 95\\% confidence upper limit on the fraction of stars with planets, we assume companions have an even logarithmic distribution in semimajor axis over several ranges of orbital period. We adopt the period ranges as outlined by \\citet{GAU05A}, for HJ and VHJ companions, and an as of yet undetected population with P$<$1.0 day, which we denote as Extremely Hot Jupiters (EHJ).  For NGC 1245, we limit the fraction of cluster members with 1.0 $R_{J}$ companions to $<$3.2\\% and $<$24\\% for EHJ and VHJ companions, respectively.  We do not reach the sensitivity to place any meaningful constraints on 1.0 $R_{J}$ HJ companions.  For 1.5 $R_{J}$ companions we limit the fraction of cluster members with companions to $<$1.5\\%, $<$6.4\\%, and $<$52\\% for EHJ, VHJ, and HJ companions, respectively.  We also fully characterize the errors associated with calculating the upper limit.  We find the overall error budget separates into two equal contributions from error in the total number of single dwarf cluster members in the sample and the error in the detection probability.  After correcting the detection probability for systematic overestimates that become increasingly important for detecting transits toward longer orbital periods (see \\S\\ref{effcalc}), we conclude that random and systematic errors in determining the number single dwarf stars in the sample dominate the error budget.  \\S\\ref{results} details the error analysis, and overall, we assign a $^{+13\\%}_{-7\\%}$ fractional error in the upper limits.  In planning future transit surveys, we demonstrate that observing NGC 1245 for twice as long will reduce the upper limits for the important HJ period range more efficiently than observing an additional cluster of similar richness as NGC 1245 for the same length of time as this data set.  To reach a $\\sim$ 2\\% upper limit on the fraction of stars with 1.5 $R_{J}$ HJ companions, where radial velocity surveys currently measure 1.3\\% \\citep{MAR05}, we conclude a total sample size of $\\sim 7400$ dwarf stars observed for a month will be needed.  If 1\\% of stars have 1.5 $R_{J}$ HJ extrasolar planets, we expect to detect one planet every 5000 dwarf stars observed for a month.  Results for 1.0 $R_{J}$ companions without substantial improvement in the photometric precision likely will require a small factor larger sample size."
    },
    "0512/astro-ph0512388_arXiv.txt": {
        "abstract": "In the course of our search for double degenerate binaries as potential progenitors of type Ia supernovae with the ESO VLT several new subdwarf B (sdB) binaries were discovered. In this paper, we present detailed analyses of six radial velocity variable sdB stars. Radial velocity curves have been measured. From the mass functions we derive lower limits to the masses of the unseen companions and we discuss their nature. In addition, stellar parameters like effective temperatures, surface gravities and helium abundances were determined as well as metal abundances. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0512/astro-ph0512341_arXiv.txt": {
        "abstract": "We have conducted a search for giant pulses from four millisecond pulsars using the 100\\,m Green Bank Telescope.  Coherently dedispersed time-series from PSR~J0218+4232 were found to contain giant pulses of very short intrinsic duration whose energies follow power-law statistics.  The giant pulses are in phase with the two minima of the radio integrated pulse profile but are phase aligned with the peaks of the X-ray profile.  Historically, individual pulses more than 10-20 times the mean pulse energy have been deemed to be ``giant pulses''. As only 4 of the 155 pulses had energies greater than 10 times the mean pulse-energy, we argue the emission mechanism responsible for giant pulses should instead be defined through: (a) intrinsic timescales of microsecond or nanosecond duration; (b) power-law energy statistics; and (c) emission occurring in narrow phase-windows coincident with the phase windows of non-thermal X-ray emission.  Four short-duration pulses with giant-pulse characteristics were also observed from PSR~B1957+20.  As the inferred magnetic fields at the light cylinders of the \\msps that emit giant pulses are all very high, this parameter has previously been considered to be an indicator of giant pulse emissivity.  However, the frequency of giant pulse emission from PSR~B1957+20 is significantly lower than for other millisecond pulsars that have similar magnetic fields at their light cylinders.  This suggests that the inferred magnetic field at the light cylinder is a poor indicator of the rate of emission of giant pulses. ",
        "introduction": "\\label{sec:introduction}  The Crab radio pulsar was discovered through the direct detection of strong individual pulses \\citep{sr68}.  Further studies revealed that the strongest pulses followed power-law energy statistics \\citep{ag72} distinct from the Gaussian statistics of the general pulse population \\citep{cor76a}.  In an observation by \\citet{lcu+95} around one in 1200 pulses had an energy greater than 20 times the mean pulse energy, $\\eav$.  Despite this, \\citet{cbh+04b} found that at all radio frequencies phase-coherent summation of the \\gps gives a higher signal-to-noise ratio than summation of all the pulses. Extraordinarily, the \\gps also have structure that is significantly narrower than the mean pulse.  \\citet{hkwe03} observed pulses with that had structure persisting for less than 2\\,ns and inferred that the brightness temperatures of these pulses are $T_{\\rm B}\\sim 10^{37}$K.  The young Crab-like pulsar B0540$-$69 in the Large Magellanic Cloud also emits giant pulses \\citep{jr03}.  In 31.2\\,hr of observations at a center frequency of 1390\\,MHz \\citet{jrmz04} only detected the integrated emission-profile of PSR~B0540$-$69 at a very low level of significance.  Despite their difficulty in detecting the integrated emission, Johnston \\etal were able to detect and analyse 141 individual pulses.  The relative ease with which giant pulses can be seen over large distances has led several authors to advocate their detection as a way to find extra-galactic pulsars \\citep[see e.g.][]{jr03,cbh+04b}.  To date, no other young pulsars have been found to emit pulses with the high energies and extremely short durations characteristic of the giant pulses from the Crab pulsar.  Three young pulsars have been found to emit narrow pulses of emission showing power-law statistics \\citep{jvkb01,jr02,cjd04}.  However, it is not clear that the pulses should be classed as true ``giant pulses'' because the power-law tails have only been seen to extend to low energies.  In addition, the structure of these events has thus far not been shown to have timescales as short as those of giant pulses from the Crab pulsar.  The recycled pulsars B1937+21, B1821$-$24, and J1823$-$3021A also emit giant pulses despite having markedly different periods ($P$) and period derivatives (\\pdot) to the Crab pulsar \\citep{cstt96,rj01,kbmo05}.  One common factor between these \\msps, PSR~B0540$-$69, and the Crab pulsar is that they all have very high magnetic fields inferred at their light cylinders\\footnote{The ATNF Pulsar Catalogue has been used to obtain the pulsar parameters and statistics used in this paper.  See: http://www.atnf.csiro.au/research/pulsar/psrcat.}  $B_{\\rm LC} \\propto P^{-2.5}\\dot{P}^{0.5}$.  When viewed in the context of the known \\msp population, the three \\gp emitters also have very low characteristic ages $\\tau = P/(2\\dot{P})$ and very high spin-down luminosities $\\dot{E} \\propto P^{-3}\\dot{P}$.  PSRs B1821$-$24, B1937+21, and J0218+4232 have some of the highest X-ray luminosities of all \\msps \\citep{bt99,gch+02,cus04,hge+05}.  The emission from all three pulsars is non-thermal, and the X-ray profiles of PSRs B1821$-$24 and B1937+21 align in phase with their \\gp emission \\citep{rj01,chk+03}.  Another field pulsar with a high X-ray luminosity is PSR~B1957+20 \\citep{bt99}.  However, no X-ray pulsations have been detected from this source and it is unclear how much of the emission originates from the bow-shock between the pulsar wind and the companion wind \\citep{sgk+03}.  In this paper we present the results of a sensitive baseband search for microsecond-timescale emission from four millisecond pulsars. Upper limits are placed on emission from PSRs J1843$-$1113 and J1012+5307, and a new population of short-duration pulses is reported for PSR B1957+20.  A previously unknown population of giant pulses from PSR~J0218+4232 is characterized and the results used to clarify the defining characteristics of giant pulse phenomenology. ",
        "conclusions": "We have searched four millisecond pulsars for individual pulses of emission with microsecond timescales and have found such emission from two of them.  Only four individual pulses were detected from PSR~B1957+20 in 8003\\,s of observations centered at 825\\,MHz.  As these pulses are exceptionally narrow there is little scattering-induced pulse-broadening at least some orbital phases. Although it is debatable whether or not these strong pulses are true ``giant pulses'', we can say that the giant pulse emission rate from PSR~B1957+20 is significantly less than the rates for other pulsars with similar values of magnetic field at the light cylinder.  Although \\blc can be used as a rough guide to whether a pulsar emits giant pulses, we suggest it is a poor indicator of the emission rate.  PSR~J0218+4232 emits giant pulses at a low rate that is inconsistent with the findings of \\citet{jkl+04a}.  It is most likely that the pulses reported by Joshi \\etal are spurious.  The giant pulses of PSR~J0218+4232 are confined to two narrow phase regions separated by roughly 50\\% of phase which align in phase with the peaks of the X-ray profile and roughly coincide with the minima of the integrated pulse-profile in the radio band.  This strong correlation between X-ray and radio properties confirms that the two emission processes originate in similarly defined regions of the pulsar magnetosphere.  Most of the 139 giant pulses observed from PSR~J0218+4232 at a center frequency of 857\\,MHz had relatively low energies, typically only a few times the mean pulse energy.  Only three had energies above $10\\eav$ and none had energies above $20\\eav$.  The pulses exhibit power-law statistics, are only found in narrow phase windows that coincide in phase with the X-ray pulse-components, and are very narrow just like the giant pulses of PSR~B1937+21; it is apparent then that ``giant'' pulses should be defined not through large flux densities, but by these three properties.  The brightest pulse seen at a center frequency of 1373\\,MHz seems to be around 500\\,ns in duration when viewed at 125\\,ns time resolution.  At higher time resolution finer features become apparent, but it is unclear whether these are significant.  PSR~J0218+4232 is the fourth \\msp found to emit giant pulses after PSRs B1937+21, B1821$-$24, and J1823$-$3021A.  All four have low characteristic ages and steep radio spectra.  With the exception of PSR~J1823$-$3021A which has not been observed in X-rays, the four pulsars all have high X-ray luminosities and exhibit power-law spectra.  The presence of X-ray emission with a steep power-law spectrum therefore seems to be the best indicator of whether a \\msp emits giant pulses.  Radio observations would be expected to show that narrow giants will be present at the phase of the X-ray emission."
    },
    "0512/astro-ph0512427_arXiv.txt": {
        "abstract": "The Ly$\\alpha$ forest at $z \\gtrsim 5.5$  shows strong scatter in the mean transmission even when smoothed over very large spatial scales, $\\gtrsim 50$ Mpc/$h$. This has been interpreted as a signature of strongly fluctuating radiation fields, or patchy reionization. To test this claim, we calculate the scatter arising solely from density fluctuations, with a uniform ionizing background, using analytic arguments and simulations. This scatter alone is comparable to that observed. It rises steeply with redshift and is of order unity by $z \\sim 6$, even on $\\sim 50$ Mpc/$h$ scales. This arises because: i) at $z \\sim 6$, transmission spectra, which are sensitive mainly to rare voids, are highly biased (with a linear bias factor $b \\geq 4-5$) tracers of underlying density fluctuations, and ii) projected power from small-scale transverse modes is aliased to long wavelength line-of-sight modes. Inferring patchy reionization from quasar spectra is therefore subtle and requires much more detailed modeling. Similarly, we expect order unity transmission fluctuations in the $z \\sim 3$ HeII Ly$\\alpha$ forest from density fluctuations alone, on the scales over which these measurements are typically made. ",
        "introduction": "\\label{intro}  At redshifts close to $z \\sim 3$, the structure in the Ly$\\alpha$ forest has been shown to arise naturally from density fluctuations in the cosmic web (e.g. Miralda-Escud\\'e et al. 1996). At sufficiently high redshift, however, the structure in the Ly$\\alpha$ forest may instead largely reflect the topology of reionization, and/or a strongly fluctuating radiation field. In high-redshift quasar spectra with extended opaque regions, significant gaps of substantial transmission occur (e.g. Becker et al. 2001, White et al. 2003, White et al. 2005). This has previously been attributed to a strongly fluctuating UV background, as expected at the tail end of reionization (Wyithe \\& Loeb 2005, Fan et al 2005).  Could these transmission gaps simply arise from underdense regions where the neutral hydrogen fraction is lower? The transmission in the $z \\sim 6$ quasar spectra differs significantly from sightline to sightline, even when one averages over co-moving length scales of $\\sim 50 - 100$ Mpc/$h$. Since the density variance is small over such large scales, one might naively expect that the reionization of the IGM must be {incomplete} near $z \\sim 6$.  In this {\\it Letter}, we critically examine this naive intuition. Is rapidly increasing scatter in sightline to sightline flux transmission a good diagnostic for patchy reionization (Fan et al. 2002, Lidz et al. 2002, Sokasian et al. 2003, Paschos \\& Norman 2005)? We find that in fact the fractional scatter in the mean transmissivity of the IGM will be large at high redshift, even for a completely uniform ionizing background. An analogous calculation applies to the case of the HeII Ly$\\alpha$ forest near $z \\sim 3$. ",
        "conclusions": "\\label{conclusion}  We do not mean to imply that the radiation field is indeed uniform at $z \\sim 6$. Our intent is only to show that ruling out the {\\em null hypothesis} that the scatter in the high redshift Ly$\\alpha$ forest results solely from density fluctuations is subtle. Indeed, we find that transmissivity fluctuations should be {\\em of order unity} at $z \\sim 6$ on scales of $L \\sim 50$ Mpc/$h$, from density fluctuations alone. We therefore caution against over-interpreting the large scatter in $\\tau_{\\rm eff}$ seen in the spectra of $z \\sim 6$ quasars or the analogous scatter seen in the $z \\sim 3$ HeII Ly$\\alpha$ forest. In future work, we intend to model the effect of inhomogeneous reionization on the statistics of the Ly$\\alpha$ forest, following up on the work of Furlanetto \\& Oh (2005), and to design statistical measures to discern the presence or absence of these fluctuations."
    },
    "0512/astro-ph0512611_arXiv.txt": {
        "abstract": "Motivated by upcoming data  from astrometric and spectroscopic surveys of  the  Galaxy, we   explore  the  chemical abundance  properties and phase-space distributions   in  hierarchically-formed   stellar   halo simulations set  in a $\\Lambda$CDM Universe.   Our sample of Milky-Way type stellar  halo  simulations result  in average metallicities that range  from  \\FeH $\\simeq -1.3$ to $-0.9$,  with  the most metal  poor  halos resulting from  accretion   histories  that lack destructive mergers with massive   (metal    rich) satellites.   Our stellar halo metallicities increase with stellar halo mass. The slope of the \\FeH $- M_{*}$ trend  mimics   that  of  the  satellite  galaxies that were destroyed to build the halos, implying that  the relation propagates  hierarchically.  All simulated halos  contain a  significant   fraction of  old  stellar  populations accreted more  than 10~Gyr ago and  in a  few cases, some intermediate age populations exist.   In contrast with  the Milky Way,  many of our simulated  stellar  halos  contain old stellar   populations which are metal  rich, originating in  the early accretion of massive satellites ($M_{*}  \\sim 10^{9}  M_{\\odot}$).  We  suggest  that  the  (metal rich) stellar halo of  M31 falls into this  category,  while the  more metal poor halo  of  the Milky Way  is  lacking in   early massive accretion events.  Interestingly, our  hierarchically-formed stellar halos often have  non-negligible   metallicity   gradients  in    both  \\FeH   and \\alphaFet. These gradients extend a few tens of kpc, and can be as large as 0.5~dex in \\FeH and 0.2~dex in \\alphaFet, with the most metal poor halo stars typically buried within the central  $\\sim 5$~kpc of the galaxy. Thus evidence for metallicity gradients alone in the  Milky Way stellar halo would not preclude   its formation via a   hierarchical process.  Only coupled with phase-space data  can metallicity information be utilized to test ideas about accreted versus in situ formation. Finally, we  find that chemical abundances can act as a rough substitute for  time of accretion of satellite galaxies and, based on   this finding, we  propose a  criterion for identifying tidal streams spatially by  selecting stars with \\alphaFe ratios below solar. ",
        "introduction": "\\label{sec:introduction}  The  stars in the halo  hold important information about the formation history of the Galaxy. Their spatial and velocity distributions can be used  to retrace their dynamic origin, and their chemical   abundances can  constrain   the  star  formation histories of their constituents. A major goal over the next decade is to obtain kinematic and chemical information for a large number of stars in our own Galaxy, thus providing a detailed reconstruction of its formation history and a link to the underlying cosmology \\citep[e.g.][]{freeman02}.  It is the goal of this work to provide a fist step towards the modeling needed to exploit these observations to their fullest potential.  The combination of spatial, kinematic and chemical data will provide a powerful discriminate   for   testing  different   Galaxy  formation    models. Historically,  metallicity gradients have acted a primary motivator in formulating hypotheses:   large-scale metallicity gradients are associated with a rapid, semi-continuous collapse \\citep[][ELS]{eggen62}; homogeneous metallicity distributions are associated with  chaotic assembly from fragments \\citep[][SZ]{searle77,searle78}.  While there are no well-motivated models that predict these opposite extremes from first principles, the poles of debate set by ELS and SZ serve as important straw-man theories for shaping ideas and testing data. Nevertheless, realistic models will always fall into less idealized categories. For example, our $\\Lambda$CDM-based model relies entirely on a hierarchical origin for the stellar halo, yet predicts some non-negligible metallicity gradients (see below).  Likewise models that allow significant ``rapid'' in situ star formation necessarily contain some accreted stellar material as a result of having a background hierarchical cosmology \\cite[e.g.][]{renda05}.  Clearly, only in combination with spatial and  kinematic data can chemical abundance data be used to its full potential.  Numerous  results over the  past decades support  the  validity of the hierarchical   model of  structure   formation \\citep{white78,blumenthal84}. In  this  model,  the expectation for stellar halos of galaxies like the Milky Way is that they form in large part by   phase-mixing of    tidal debris   from  the numerous     accreted satellites. Although the dynamical  evolution of galaxy halos has been extensively modeled, and we have  now predictions for the  phase-space distribution of both dark matter and  stars in the  halo of the Galaxy \\citep[e.g.][]{johnston98,helmi99,bullock01,helmi03,abadi05,bullock05,diemand05,moore05}, currently    there are  no    clear   predictions for  the  associated phase-space distribution of chemical elements.  Intuitively, one expects   stochastic  accretions  to leave   peculiar signatures in both the  kinematics and chemical abundances of  present day stars. Some observational evidence in that respect is indeed found in  our  Galaxy.  The halo contains  stars  which  stand out  in  both metallicity                       and                         velocity \\citep{carney96,nissen91,nissen97,majewski96,chiba00,altmann05},    as does the  disk \\citep{helmi99,navarro04,helmi05}. Similar evidence has been found in M31 $-$ the only  other spiral galaxy whose stellar halo has been studied at a comparable level of detail  as that of the Milky Way \\citep{ferguson02,reitzel02,bellazzini03,ferguson05,guhathakurta05a}. The giant stellar stream   in M31, a merger   debris  extending more  than 100~kpc to the SE of  M31's disk, stands  out from the background halo not      only  in       stellar     over-densities     and   kinematics \\citep{ibata01a,ibata04,guhathakurta05b},  but  also   in  metallicity \\citep{ferguson02}.   Over  the next decade an immense  amount of data will become available for stars   in the  Galaxy  from  both  astrometric  and spectroscopic surveys     $-$  e.g.    from    the    astrometric  satellite    GAIA \\citep[e.g.][]{perryman01},  the ``Radial Velocity Experiment'' (RAVE) \\citep[e.g.][]{steinmetz03},  and  the ``Sloan  Extension for Galactic Underpinnings                  and                      Exploration'', SEGUE\\footnote[1]{http://www.sdss.org}.  Positions,  line     of sight velocities  and proper motions will  be measured for millions of stars in      the   halo,    enabling  the   full       phase-space  to   be reconstructed.  Similarly,  high accuracy, wide  field measurements of stellar spectra will  allow  the mapping   of the  Galaxy  in chemical abundances.  These studies will not only constrain cosmology and galaxy formation \\citep[e.g.][]{bullock01} but also test ideas about first light and reionization \\citep[e.g.][]{tumlinson04} With the upcoming observational data, the goal of putting together the history of our Galaxy is finally becoming feasible. It is therefore important to have theoretical  predictions for the  combined kinematic  and chemical abundance distributions, as  well as to have a coherent theoretical framework for interpreting the upcoming results.  Although  future surveys promise  to   provide  detailed maps of   the chemical and phase-space  structure of the Galaxy,  we may worry  that ours is but one among many galaxies  of its size.   How typical is our galaxy?  How can it be used to constrain general ideas about cosmology and galaxy formation  if it is  atypical in some way?  Fortunately, in the context of  $\\Lambda$CDM,  we have  well  defined expectations  for variations in the formation times and accretion histories of Milky-Way size galaxies and this idea is  at the heart  of our exploration.  For example \\cite{mouhcine05a,  mouhcine05b}  recently have  estimated the metallicities of the inner halos of several nearby spiral galaxies and find that they  are quite metal  rich compared to  the stellar halo of the Milky Way.  They  go  on to suggest  that  the  Milky Way  is  not typical for a normal  spiral galaxy of  its luminosity.   As discussed below, if this result holds then our  models provide a straightforward interpretation: that the Milky Way  has experienced fewer than average major accretion events  with  LMC-size objects.   Indeed, phase  space information could test  ideas of this  kind specifically by  revealing few major accretion events and many more lower-mass accretions.   The  modeling of the  chemical evolution of  the Galaxy in a realistic cosmological   context    is     just  beginning  to     be   explored \\citep[e.g.][]{bekki01,brook03,brook04}. Recent advances in   modeling techniques, like the advent of ``hybrid'' methods, allow us to address this   problem with  greater  resolution.  Hybrid  methods enable  the modeling  of star formation   and  chemical enrichment  with  detailed semi-analytical prescriptions,  and   at  the same  time  provide  the detailed     dynamics     via     the coupled      N-body  simulations \\citep[e.g.][]{kauffmann99,springel01}.  In   this  study   we  use  a previously developed hybrid  method \\citep{bullock05,robertson05,font05} to explore the  phase-space distribution of  chemical abundances  in a series  of  Milky  Way-type galaxy halos    formed in  a  $\\Lambda$CDM Universe   ($\\Omega_{m}=0.3,  \\, \\Omega_{\\Lambda}=0.7, \\,  h=0.7,$ and $\\sigma_{8}=0.9$).     A description of our sample of stellar galaxy models is given in Section \\ref{sec:sample}. In Section \\ref{sec:results} we present results on the spatial distribution of \\FeH and \\alphaFe chemical abundances in these stellar halos and propose a criterion for detecting cold stellar streams based on their \\alphaFe abundances. In Section \\ref{sec:comparisons} we discuss our results and compare them with the available observations and in Section \\ref{sec:conclusions} we conclude. ",
        "conclusions": "\\label{sec:conclusions}  The  motivation  of   this work was   to provide   predictions of  the hierarchical formation paradigm that can be  tested with the data from current and  upcoming astrometric and  spectroscopic observations.  In this  study we have explored  the global properties as the phase-space distribution of the \\FeH and  \\alphaFe chemical abundances in a series of Milky  Way-type stellar  halos  formed in a  $\\Lambda$CDM Universe. Our results can be summarized as follows:  \\vskip 4pt - For a fixed dark matter halo mass, we predict that stellar halo metallicities increase with stellar halo mass. The resultant \\FeH $- M_{*}$ relation for stellar halos arises because high-mass stellar halos tend to have accreted a larger number of massive satellite progenitors.  The relation propagates hierarchically because massive progenitors are more metal rich than their low-mass counterparts.   \\vskip 4pt - Chemical abundance distributions provide important links to the minor  merger histories of  galaxies.  For    example,  the observed  metallicity differences between  the (metal poor) Milky Way  and (metal rich) M31 stellar halos may arise because M31 destroyed one or two more massive satellites than the Milky Way.  \\vskip 4pt - In agreement  with  observations, we  find that  all simulated halos contain an underlying population of old metal poor stars.  More massive halos contain additional intermediate age populations accreted more recently.  \\vskip 4pt - In contrast with the Milky Way, some of  our halo models contain old stellar populations which   are  metal rich,  originating   in massive satellites ($M_{*}  \\sim  10^{9} M_{\\odot}$) accreted   early on.  The metallicity constraints  seems  to suggest that the  Milky  Way has an atypical   mass accretion  history, lacking in massive early accretion events,   compared with  other  galaxies of a similar mass.  \\vskip 4pt - We predict metallicity gradients in \\FeH  and \\alphaFe in our {\\it hierarchically-formed} stellar halos that are often non-negligible. Gradients extend typically  over a few tens  of kpc and may be as large as $0.5$~dex in \\FeH and $0.2$~dex in \\alphaFet.   \\vskip 4pt - We predict that most metal poor stars in the Galactic halo are buried within the central $\\sim 5$kpc of the Galaxy.  This will likely be important for efforts to use the observed count of low metallicity stars to constrain models of cosmic reionization and the transition from Pop III to Pop II star formation \\citep[e.g.][]{tumlinson04}.   \\vskip 4pt - We suggested a method to  identify cold tidal streams  in stellar halos  by selecting stars   with \\alphaFe $  \\le 0.05$,  which trace streams from recently accreted satellites. If used to complement the common kinematical  identification methods, the proposed selection criterion may improve the detection of tidal streams.   \\vskip4pt Combining kinematic and chemical abundance information  proves to be a powerful  tool  for understanding how   galaxies  form and  evolve. In anticipation   to the wealth  of data  to be  delivered by current and future Galactic  surveys, it is  important  to develop and  improve on numerical simulations of  the chemical  evolution of  the Galaxy in  a realistic  cosmological context $-$  such as   those presented in  our study $-$ and to devise criteria to identify more tidal streams.  While our methods successfully reproduce the gross chemical abundance properties of the Galactic halo and satellite dwarf galaxies, our treatments of dissipational physics including star formation, gas cooling, and energetic feedback mechanisms may be too simplistic to recover the detailed properties of the modeled systems with high accuracy.  For instance, the assumed truncation of star formation in dwarf systems after they accrete into their host galaxy is very approximate. Gravitational torques applied by the central galaxy and other satellite systems may induce episodes of efficient star formation in the dwarfs we model, especially near pericentric passage. Such bursts occurred in the LMC and SMC, as inferred from their stellar populations \\citep[e.g.][]{smecker-hane02,harris04}, and should be modeled if more accurate chemical enrichment histories for dwarf systems are to be achieved.  These complications speak to the need for high-resolution hydrodynamical simulations of the formation of the Galactic stellar halo and dwarf satellite population.  While recent hydrodynamical simulations have begun to address this issue \\citep[e.g.][]{brook04},  they have not yet matched the mass resolution necessary to capture the entire population of dwarfs galaxies we follow in our collisionless simulations nor the full complexity of our chemical enrichment model.  We acknowledge the importance of the previous work and look forward to increasingly  sophisticated hydrodynamical simulations of the cosmological formation of the Galactic stellar halo."
    },
    "0512/astro-ph0512082_arXiv.txt": {
        "abstract": "We study the evolution of the ionization state of the intergalactic medium~(IGM) at the end of the reionization epoch using moderate resolution spectra of a sample of nineteen quasars at $5.74 < z_{\\rm em} < 6.42$ discovered in the Sloan Digital Sky Survey. Three methods are used to trace IGM properties: (a) the evolution of the Gunn-Peterson (GP) optical depth in the Ly$\\alpha$, $\\beta$, and $\\gamma$ transitions; (b) the distribution of lengths of dark absorption gaps, and (c) the size of HII regions around luminous quasars. Using this large sample, we find that the evolution of the ionization state of the IGM accelerated at $z>5.7$: the GP optical depth evolution changes from $\\tau^{\\rm eff}_{\\rm GP} \\sim (1+z)^{4.3}$ to $(1+z)^{\\gtrsim 11}$, and the average length of dark gaps with $\\tau>3.5$ increases from $<10$ to $>80$ comoving Mpc. The dispersion of IGM properties along different lines of sight also increases rapidly, implying fluctuations by a factor of $\\gtrsim 4$ in the UV background at $z>6$, when the mean free path of UV photons is comparable to the correlation length of the star forming galaxies that are thought to have caused reionization. The mean length of dark gaps shows the most dramatic increase at $z\\sim 6$, as well as the largest line-of-sight variations. We suggest using dark gap statistics as  a powerful probe of the ionization state of the IGM at yet higher redshift. The sizes of HII regions around luminous quasars decrease rapidly towards higher redshift, suggesting that the neutral fraction of the IGM has increased by a factor of $\\gtrsim 10$ from $z=5.7$ to 6.4, consistent with the value derived from the GP optical depth. The mass-averaged neutral fraction is $1 - 4$\\% at $z\\sim 6.2$ based on the GP optical depth and HII region size measurements. The observations suggest that $z\\sim 6$ is the end of the overlapping stage of reionization, and are inconsistent with a mostly neutral IGM at $z\\sim 6$, as indicated by the finite length of dark absorption gaps. ",
        "introduction": "After the recombination epoch at $z\\sim 1100$, the universe became mostly neutral, until the first generation of stars and quasars reionized the intergalactic medium (IGM) and ended the cosmic ``dark ages'' (e.g., Rees 1998). Cosmological models predict reionization at redshifts between 6 and 20 (e.g., Gnedin \\& Ostriker 1997, Gnedin 2000, 2004, Ciardi et al. 2001, Benson et al. 2001, Razoumov et al. 2002, Cen 2003a,b, Ciardi, Ferrara \\& White 2003, Haiman \\& Holder 2003, Wyithe \\& Loeb 2003a, Somerville et al. 2003, \\cite{choudhury05}). When and how the universe reionized remains one of the fundamental questions of modern cosmology (for reviews of theoretical models, see Barkana \\& Loeb 2001, Loeb \\& Barkana 2001, and Ciardi \\& Ferrara 2005).  The last few years have witnessed the first direct observational constraints on the history of reionization. Lack of complete Gunn-Peterson (GP, 1965) absorption at $z<6$ indicates that the IGM is highly ionized by that epoch (e.g. Fan et al. 2000, 2001, Becker et al. 2001, Djorgovski et al. 2001, Songaila \\& Cowie 2002). Spectroscopic observations of the highest-redshift quasars and galaxies provide a number of observational constraints suggesting that $z\\sim 6$ indeed marks the end of the reionization epoch: \\begin{itemize} \\item The detection of complete GP absorption troughs in the spectra of quasars at $z>6$ suggests that the neutral fraction of the IGM is increasing rapidly with redshift (Cen \\& McDonald 2002, Fan et al. 2002, \\cite{lidz2002}, White et al. 2003, \\cite{gnedin2004}). This feature is consistent with the phase of IGM evolution at which  individual HII regions overlapped; at this point,  the IGM ionization state evolved rapidly, similar to a phase transition (cf. Songaila 2004 and discussion below). The lower limit of the volume averaged neutral fraction is $\\sim 10^{-3}$ based on Gunn-Peterson measurements, although it could be much higher in the deepest troughs. \\item The sizes of large, highly ionized regions around luminous quasars at $z>6$ are consistent with those predicted for a quasar ``\\strom\\ sphere'' expanding into a surrounding largely neutral IGM (\\cite{mesinger04a}, \\cite{mesinger04b}, \\cite{wyithe04a}). The neutral fraction of the IGM is estimated to be as high as $>20$\\% along some lines of sight, considerably higher than the lower limit using the Gunn-Peterson optical depth. But these measurements are made uncertain because of our lack of knowledge of quasar lifetimes, bolometric corrections of quasar luminosities, the bias factor and the clumpiness of the IGM around high redshift quasars (e.g., \\cite{yu05a}, \\cite{yu05b}). \\item The observed temperature evolution of the IGM at $z\\sim 3-5$ suggests that the reionization epoch is not much higher than 8 (\\cite{hui03}). In models with reionization at much higher redshift, the IGM cools to a much lower temperature than that observed at $z\\sim 2 - 4$. \\end{itemize} Quasars and AGNs are unlikely to provide enough UV photons to reionize the universe at $z>6$ (e.g. Fan et al. 2001, \\cite{dijkstra04}, \\cite{meiksin05}). Star forming galaxies at early epochs are the most likely candidates to ionize the IGM, although current observations still have poor constraints on the total ionizing photon budget at $z\\sim 6$ (\\cite{yan04b}, \\cite{stiavelli04}). Recent measurements of the luminosity density of Lyman break galaxies using deep HST observations show a moderate decline in total UV luminosity of star forming galaxies at $z>6$ (\\cite{bunker04}, \\cite{yan04b}, \\cite{bouwens04a},b). If the reionization process was similar to a phase transition, the high-redshift quasar observations suggest that the Universe could have been mostly neutral as late as $z = 6 - 8$.  Polarization measurements of the cosmic microwave background (CMB) based on WMAP first-year data (Kogut et al. 2003) show a large optical depth due to Thompson scattering of electrons ($\\tau = 0.17 \\pm 0.04$) in the early Universe, suggesting that the IGM was largely ionized by $z\\sim 17\\pm 4$. WMAP three-year data (Spergel et al. 2006) show a smaller optical depth ($\\tau = 0.09 \\pm 0.03$) based on the new EE measurements, indicating a larely ionized IGM by $z\\sim 10 \\pm 3$. The WMAP polarization results, combined with low redshift constraints, suggest that reionization could be an extended process rather than a phase transition (Cen 2003a, b, \\cite{wyithe03a}, \\cite{haiman03}, \\cite{somerville03b}, \\cite{gnedin2004}, \\cite{choudhury05}). Even at $z\\sim 6$, several lines of evidence suggest that a uniform end of reionization with  a clean phase transition in the ionization state is too simplistic: \\begin{itemize} \\item While the optical depth no doubt has increased at $z>6$, it is not agreed whether the increase is similar to a sharp phase transition (e.g., Fan et al. 2002), or is simply a continuation of the gradual thickening of the Ly$\\alpha$ forest (e.g. Songaila 2004). Although optical depth measurements of quasar spectra are consistent among different studies, the different interpretations arise because: (1) complete Gunn-Peterson troughs give only a lower limit to the  optical depth; (2) there are complications in the interpretation of  the Ly$\\beta$ measurements; and (3) there is a lack of a clear definition of ``gradual'' vs. ``sharp'' transitions. \\item There are large line of sight variations from one quasar to another. The GP trough detected in the spectrum of SDSS J1030+0524 ($z=6.28$, Fan et al. 2001, Becker et al. 2001, White et al. 2003) shows complete absorption in the Ly$\\alpha$, $\\beta$ and $\\gamma$ transitions. However, the line of sight of SDSS J1148+5251 ($z=6.42$, Fan et al. 2002) shows clear transmission, especially in the Ly$\\beta$ and $\\gamma$ transitions (White et al. 2003, 2005, \\cite{OF2005}), indicating that the IGM along this line of sight is still more than 99\\% ionized. \\item The lack of evolution in the luminosity functions of Ly$\\alpha$ emitting galaxies from $z\\sim 5.7$ to $z\\sim 6.5$ is also consistent with the picture that the IGM is largely ionized at $z\\sim 6$; in a largely neutral IGM, an extended Gunn-Peterson absorption damping wing would attenuate the Ly$\\alpha$ emission of high-redshift galaxies (\\cite{hu04}, 2005, \\cite{malhotra04}, \\cite{stern05}, \\cite{malhotra06}). However, due to uncertainties in the clustering of Ly$\\alpha$ galaxies and the clumpiness of the IGM, a relatively high neutral fraction cannot yet be ruled out by the \\lya emitter observations (\\cite{santos04}, \\cite{haiman05}, \\cite{wyithe05}, \\cite{furlanetto05}). \\end{itemize}  Reionization is a complex process which results from the interplay of structure formation, early star formation and feedback, and radiative transfer in a clumpy IGM. Recent developments in both semi-analytic models and simulations aim at understanding these processes in the context of improved observational constraints and future observations. Models that include feedback from the formation of the earliest, metal-free populations can generally fit both the Gunn-Peterson measurements at low redshift and a high CMB polarization measurement (e.g., Cen 2003a, b, Wyithe \\& Loeb 2003a, Haiman \\& Holder 2003, \\cite{melchiorri05}). Cosmological simulations with improved resolution, box size and treatment of radiative transfer (e.g. Gnedin 2000, 2004, Razoumov et al.\\ 2002, \\cite{pn05}) are beginning to make realistic comparison with observations possible.  To make further progress observationally, spectroscopy of a large sample of luminous sources is needed to explore the line of sight variation in the GP optical depth measurements expected due to clumpiness in the IGM and clustering of ionizing sources (e.g. Wyithe \\& Loeb 2004b, \\cite{pn05}). Fan et al. (2002, Paper I) presented measurements of the ionization state of the IGM at $z\\sim 6$ and constraints on the epoch of reionization, using a sample of four quasars at $z=5.74 - 6.28$. In this follow-up paper, we describe measurements of Lyman absorption using spectroscopic observations of a  sample of nineteen quasars at $z=5.74 - 6.42$, discovered from the imaging data of the Sloan Digital Sky Survey (Fan et al. 2000, 2001, 2002, 2003, 2004, 2006). We address the following three questions: \\begin{enumerate} \\item  How fast do  {\\bf the average and the dispersion} of the IGM ionization state evolve in the redshift range $z = 5 - 6.4$? In particular, is there a well-defined break in these quantities as a function of redshift?  \\item Is the IGM consistent with a large neutral fraction at least along some lines of sight? What is the {\\em upper} limit of the neutral fraction when considering all the observational constraints?  \\item Can we define a set of statistics from the observed spectrum that are also easily extractable from cosmological simulations and allow direct comparison with theoretical models (e.g., Songaila \\& Cowie 2002, \\cite{pn05})? \\end{enumerate}  The paper is organized as follows: in \\S2, we describe our  spectroscopic sample of 19 SDSS quasars at $z=5.74 - 6.42$. In \\S3, we measure the evolution of Gunn-Peterson optical depths. We present \\lya measurements in \\S3.1, discuss in detail the optimal way of combining Ly$\\alpha$, $\\beta$ and $\\gamma$ measurements in \\S3.2, and present the measurement of the dispersion of optical depth along different lines of sight in \\S3.3. In \\S4, we derive a number of quantities characterizing the ionization state of the IGM: the evolution of the ionizing background based on a photoionization model (\\S4.1), and of the mean free path of ionizing photons and the neutral fraction of the IGM (\\S4.2), following the methods developed in Paper I.  We compare these results with simulations in \\S4.3. In \\S5, we study the presence of complete Gunn-Peterson absorption troughs in quasar spectra (\\S5.1), and discuss in detail the statistics of the dark gap distribution (\\S 5.2). In \\S5.3, we use these statistics to place an independent upper limit on the IGM neutral fraction. In \\S6, we measure the sizes of HII regions around these 19 quasars, and use their evolution to place an independent constraint on the rate of neutral fraction evolution at $z\\sim 6$. Finally, we summarize the results in \\S7. Through the paper, we use the WMAP cosmology ($\\rm H_0 = 71\\ km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\Lambda} = 0.73$, $\\Omega_{M} = 0.27$, and $\\Omega_b = 0.04$, Spergel et al. 2003) to present our results.  \\begin{figure} \\epsscale{0.9} \\plotone{Fan.zspec19.eps} \\caption{Spectra of our sample of nineteen SDSS quasars at $5.74 < z < 6.42$. Twelve of the spectra were taken with Keck/ESI, while the others were observed with the MMT/Red Channel and Kitt Peak 4-meter/MARS spectrographs. See Table 1 for detailed information.} \\end{figure} ",
        "conclusions": "The goal of this paper is to study the evolution of the IGM ionization state near the end of reionization using absorption spectra of the highest redshift quasars known. We use three types of measurements to study IGM properties. We now discuss their advantages and disadvantages, with an eye towards application to yet higher redshift.  \\noindent {\\bf 1. Evolution of Gunn-Peterson optical depth.} In \\S3, we measured the GP optical depth in the \\lya, \\lyb and \\lyc transitions. This is directly linked to the neutral hydrogen density of the IGM. Using a photoionization model, we can compute the underlying UV ionizing background, and estimate the neutral fraction and mean free path of UV photons. However, as the optical depth increases to $\\tau \\gtrsim 5$ at $z>6$, the \\lya transition becomes completely saturated. It is difficult to place stronger constraints as the measurement begins to be dominated by systematic observational errors in the red and near-IR part of the spectrum.  The \\lyb and \\lyc transitions (\\S3.2) give more stringent constraints due to their smaller oscillator strengths, but they can only improve the limit by a factor of $2-4$. In addition, one has to correct for the foreground absorption from lower order transitions. Direct measurement of the GP optical depth can only constrain the neutral fraction of the IGM to lower limits in the range $\\sim 10^{-3}$ to $10^{-2}$. Since at high redshift, all IGM transmission is through a small number of transmitting spikes with high ionization and/or low density, the {\\em effective} optical depth is actually {\\em not} sensitive to the region of the IGM where most of the neutral hydrogen lies. Furthermore, the interpretation of optical depth depends on the photoionization model and IGM density distribution models used. The exact values of derived ionization parameters, such as ionizing background and neutral fraction, are therefore model-dependent. But for a narrow redshift range, our calculation shows a robust strong evolutionary trend on the ionization parameters. Thus using the evolution of GP optical depth is most useful at the low redshift end, and it begins to lose its usefulness when the neutral fraction increases to $\\sim 10^{-3}$.  \\noindent {\\bf 2. Distribution of Gunn-Peterson troughs and dark gaps.} In \\S4, we discussed using the distribution of absorption gaps to characterize the evolution of the IGM. Qualitatively, these statistics reveal very similar trends to the GP optical depth measurements. But they also contain high order information, related to the clustering of transmission spikes and dark gaps. The evolution of dark gaps appears to be the most sensitive quantity tracing the reionization process: of the statistics we examine, it shows the most dramatic evolution at $z>6$. The distribution of dark gap length is readily computable from both observations and cosmological simulations. Simulations show that even when the average neutral fraction is high, the dark gap length is still finite due to the presence of regions that were ionized by star-forming galaxies. In \\S5.3, we used the length of dark gaps and the observed galaxy luminosity function at $z\\sim 6$ to put an independent upper limit on the  IGM neutral fraction. Of course, these statistics are not straightforward to calculate from analytic or semi-analytic models, and must  be derived from detailed simulations that include the clustering of the IGM and ionizing sources, and take into account spectral resolution and S/N of observed spectra.  \\noindent {\\bf 3. HII regions around high-redshift sources.} In \\S6, we measured the  HII region sizes around high-redshift quasars, and showed that the sizes decrease significantly with redshift, consistent with a rapid increase in the IGM neutral fraction. This method provides a powerful and independent way of measuring the IGM neutral fraction, and can in principle be applied to regions with high neutral fraction if luminous sources can be discovered. The disadvantage is that the sizes of HII regions are quite model dependent: theoretical calculations require not only knowledge of IGM properties, but also the intrinsic properties of quasars, including their lifetime and bias relative to the density field at high redshift. Further theoretical modelling will improve this constraint. As these quantities are unlikely to change over the redshift range we probe, we  used the HII region sizes as a robust measure of the {\\em fractional} change in the neutral fraction over this redshift range.  In \\S1, we posed three questions regarding the evolution of the IGM and the end of reionization. We now present our major conclusions based on our analysis of this sample of nineteen $z>5.7$ quasars.  \\begin{enumerate} \\item There is a strong evolution of the IGM ionization state at high redshift, in terms of the effective GP optical depth, the derived UV background, the neutral fraction, the mean free path of UV photons, and the extent of dark gaps. This evolution  {\\em accelerates} at $z>5.7$. The GP optical depth evolution changes from $\\tau^{\\rm eff}_{\\rm GP} \\sim (1+z)^{4.3}$ to $(1+z)^{\\gtrsim 11}$, indicating a lower limit in the  increase in the volume-averaged neutral fraction of a factor of  7 between $z=5.5$ and 6.2.  \\item Accompanying the increased average neutrality of the IGM is a rapid increase in the relative dispersion of IGM properties along different lines of sight. We find that this increased dispersion implies fluctuations by a factor of $\\gtrsim 4$ in the UV background at $z>6$. This is not surprising, as the mean free path of UV photons at $z\\sim 6$ is  comparable to the correlation length of star forming galaxies which likely provide most of the photons that ionized the Universe.  \\item The most effective measure of the IGM evolution at high redshift is the distribution of dark absorption gaps in the IGM. The average length of dark gaps shows a dramatic increase at $z\\sim 6$. This quantity shows the largest line of sight variations: dark gaps as long as 50 comoving Mpc appear by $z\\sim 5.5$, while some lines of sight at $z>6$ are still somewhat transparent.  \\item Using the evolution of HII region sizes around the quasars, we find the neutral fraction of the IGM has increased by $\\sim 14$ between $z=5.5$ and 6.4, consistent with our GP optical depth estimates. This implies a mass-averaged neutral fraction as high as $\\sim 4$\\% in regions around luminous quasars at $z\\sim 6.4$.  \\item There is no compelling evidence that the IGM at $z\\sim 6.4$ is largely neutral. On the contrary, the finite length of dark gaps in the spectra implies an upper limit on the neutral fraction less than 50\\%, likely even lower.  \\end{enumerate}  \\begin{figure} \\epsscale{1.00} \\plotone{Fan.fHm.eps} \\caption{The evolution of the mass-averaged neutral fraction using constraints from GP optical depth measurements, HII region size evolution, and distribution of dark absorption gaps.} \\end{figure}  Figure 15 summarizes our constraints on the mass-averaged neutral fraction using quasar spectra presented in this paper. The analysis here confirms the main result of Paper I based on only four quasars: the detections of deep GP troughs indicate an accelerated rate of IGM evolution at $z>5.7$, consistent with the IGM transition at the end of the overlapping stage of reionization. However, the results do not indicate that the IGM has achieved a high level of neutrality at $z\\sim 6$. The large dispersion of the IGM properties along different lines of sight  strongly suggests that the reionization process is complex and not likely a uniform phase transition over a very narrow redshift range.  If luminous sources at $z\\sim 6 - 10$ can be discovered by JWST and next generation wide-field IR surveys, statistics of dark gaps in their absorption spectra and measurements of HII region sizes surrounding these sources will provide the most useful probe of the history of reionization using high-redshift discrete sources, and complement constraints from CMB polarization measurements (e.g. PLANCK) and 21cm experiments (e.g. LOFAR, MWA, PAST and SKA).  Funding for the Sloan Digital Sky Survey (SDSS) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, the Max Planck Society, and the HEFCE. The SDSS is a joint project of The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. We thank Andrea Ferrara, Doug Finkbeiner, Steve Furlanetto, Nick Gnedin, Zoltan Haiman, Avery Meiksin, Mike Norman, Peng Oh, Stuart Wyithe,  and Idit Zehavi for helpful discussions, and the referee for valuable comments and suggestions. XF thanks KITP for hospitality during part of this project. We acknowledge support from NSF grant AST 03-07384, a Sloan Research Fellowship, a Packard Fellowship for Science and Engineering (X.F.), NSF grant AST 03-07409 (M.A.S., ) the Institute of Geophysics and Planetary Physics (operated under the auspices of the US Department of Energy by Lawerence Livermore National Laboratory under contract No. W-7045-Eng-48) (R. H. B.) and NSF grant AST 03-07582 (D.P.S.)."
    },
    "0512/astro-ph0512226_arXiv.txt": {
        "abstract": "We report phase-referencing VLBA observations of H$_{2}$O masers near the star-forming region W3(OH) to measure their parallax and absolute proper motions. The measured annual parallax is 0.489 $\\pm$ 0.017  milli-arcseconds (2.04 $\\pm$ 0.07 kpc), where the error is dominated by a systematic atmospheric contribution. This distance is consistent with photometric distances from previous observations and with the distance determined from CH$_3$OH maser astrometry presented in a related paper. We also find that the source driving the H$_{2}$O outflow, the ``TW-object'', moves with a 3-dimensional velocity of $>$ 7 km s$^{-1}$ relative to the ultracompact \\ion{H}{2} region W3(OH). ",
        "introduction": "The annual parallax is the most direct measurement of distances in astronomy. The Hipparcos satellite successfully measured the distances to numerous stars in the Solar neighborhood, typically achieving 10\\% accuracies for distances of $\\approx 100$ pc, which contributed significantly to many fields of modern astronomy (e.g. Perryman et al. 1995). However, annual parallax measurements for stars with kpc distances require sub-milliarcsecond accuracy, which has not been achieved optically.  Very Long Baseline Interferometry (VLBI) provides the highest resolution in astronomy. In phase-referencing VLBI, the position of a target source is measured relative to a nearby positional reference source (see e.g. Beasley \\& Conway 1995; Ros 2003). The feasibility of annual parallax measurements with the Very Long Baseline Array (VLBA) has been demonstrated at low frequencies by Brisken et al. (2002) who measured annual parallaxes of pulsars in the Galaxy and by van Langevelde et al. (2000) and  Vlemmings et al. (2002) who measured distances of Galactic OH masers associated with late type stars. Chatterjee et al. (2004) measured pulsar parallaxes at 5 GHz and showed that the accuracy of astrometric measurements improves with higher frequencies. Their results indicate that one can measure distances of up to a few kpc with better than 10\\% uncertainty with VLBA astrometry of maser sources. Indeed, Kurayama et al. (2005) used the VLBA to measure the annual parallax of the Mira-Type star UX Cygni with high accuracy.  Hence, VLBA measurements allow sources spread over a large part of the Milky Way to have accurate parallaxes. This enables us to probe Galactic structure and dynamics since maser sources are spread over the whole Galaxy and, especially water vapor (H$_2$O) maser sources are even found in its outer reaches (e.g. Wouterloot et al. 1993).   The 22.2 GHz transition of  H$_2$O is the most widespread and luminous known maser line. In our Galaxy it has been detected toward numerous evolved red giant stars and high- and low-mass star-forming regions (see, e.g., Valdettaro et al. 2001).   W3(OH) is a region containing several high- and intermediate-mass young stars and proto-stars of different evolutionary stages (e.g. Wilner et al. 1999 ; Wyrowski et al. 1997, 1999). In addition to strong OH and CH$_3$OH masers, which are seen projected on the archetypical ultracompact \\ion{H}{2} (UC\\ion{H}{2}) region, very strong H$_2$O maser emission is found toward the Turner-Welch (TW) Object (Turner \\&\\ Welch 1984; Reid et al. 1995; Wilner et al. 1999), a protostar projected $\\approx 10^{4}$ AU east of the UC\\ion{H}{2} region. The W3(OH) H$_2$O masers were amongst the first studied with VLBI (Moran et al. 1973). VLBI Maps of the H$_2$O maser emission have been reported by Alcolea et al (1992).  We observed W3(OH) to measure its annual parallax and to study the internal dynamics of the known bipolar H$_2$O outflow from the TW object. Moreover, our observations constitute a trial parallax and proper motion observation to explore the potential of utilizing H$_2$O masers as probes of Galactic structure. Here we report VLBA observations of the W3(OH) H$_2$O masers which yielded an extremely accurate parallax. ",
        "conclusions": "We have measured the annual parallax of the H$_{2}$O maser source in the W3(OH) region with phase-referenced VLBA observations. The distance of 2.04 $\\pm$ 0.07 kpc that we obtain is consistent with previous photometric distance estimates (but with much higher accuracy) and with the CH$_3$OH maser parallax corresponding to 1.95 $\\pm$ 0.04  kpc determined by Xu et al. (2005) in the related paper.  We also measured the proper motions of the W3(OH)-TW H$_{2}$O masers and find that the TW object is moving with a speed of $>$ 7 km s$^{-1}$ with respect to the nearby UC\\ion{H}{2} region (with its OH and CH$_3$OH masers).  Such a large speed difference between two massive objects in the same star forming region is puzzling.  Although H$_2$O masers are not perfect target sources to investigate Galactic structure and dynamics, they can still provide important information about regions in the Galaxy that are not accessible otherwise (e.g. the outer Galaxy)."
    },
    "0512/astro-ph0512010_arXiv.txt": {
        "abstract": "We show that the epoch(s) of reionization when the average ionization fraction of the universe is about half can be determined by correlating Cosmic Microwave Background (CMB) temperature maps with 21-cm line maps at degree scales ($l\\sim 100$). During reionization peculiar motion  of free electrons induces the Doppler anisotropy of the CMB, while density fluctuations of neutral hydrogen induce the 21-cm line anisotropy.  In our simplified model of inhomogeneous reionization, a positive correlation arises as the universe reionizes whereas a negative correlation arises as the universe recombines; thus, the sign of the correlation provides information on the reionization history which cannot be obtained by present means. The signal comes mainly from large scales ($k\\sim 10^{-2}~{\\rm Mpc}^{-1}$) where linear perturbation theory is still  valid and complexity due to patchy reionization is averaged out.  Since the Doppler signal comes from ionized regions and the 21-cm comes from neutral ones, the correlation has a well defined peak(s) in redshift when the average ionization fraction of the universe is about half.  Furthermore, the cross-correlation is much less sensitive to systematic errors, especially foreground emission, than the auto-correlation of 21-cm lines:  this is analogous to the temperature-polarization correlation of the CMB being more immune to systematic errors than the  polarization-polarization. Therefore, we argue that the Doppler-21cm correlation provides a robust measurement of the 21-cm anisotropy, which can also be used as a diagnostic tool for  detected signals in the 21-cm data --- detection of the cross-correlation provides the strongest confirmation that the detected signal is of cosmological origin. We show that the Square Kilometer Array can easily measure the predicted correlation signal for 1~year of survey observation. ",
        "introduction": "\\label{sec:introduction}  When and how was the universe reionized?  This question is deeply connected to the physics of formation and evolution of the first generations of ionizing sources (stars or quasars or both) and the physical conditions in the interstellar and the intergalactic media in a high redshift universe.  This field has been developed mostly theoretically \\citep{barkana/loeb:2001,bromm/larson:2004,ciardi/ferrara:2005, iliev/etal:2005,alvarez/bromm/shapiro:2005} because there are only a very limited number of  observational probes of the epoch of reionization: the Gunn--Peterson test \\citep{gunn/peterson:1965,becker/etal:2001}, polarization of the Cosmic Microwave Background (CMB)  on large angular scales \\citep{zaldarriaga:1997,kaplinghat/etal:2003,kogut/etal:2003}, mean intensity \\citep{santos/bromm/kamionkowski:2002,salvaterra/ferrara:2003,cooray/yoshida:2004,madau/silk:2005,fernandez/komatsu:2005} and fluctuations \\citep{magliocchetti/salvaterra/ferrara:2003,kashlinsky/etal:2004,cooray/etal:2004,kashlinsky/etal:2005} of the near infrared background from redshifted UV photons, Ly$\\alpha$-emitters at high redshift \\citep{malhotra/rhoads:2004,santos:2004,furlanetto/hernquist/zaldarriaga:2004,haiman/cen:2005,wyithe/loeb:2005} and fluctuations of the 21-cm line background from neutral hydrogen atoms during reionization \\citep{ciardi/madau:2003,furlanetto/sokasian/hernquist:2004,zaldarriaga/furlanetto/hernquist:2004} or even prior to reionization \\citep{scott/rees:1990,madau/meiksin/rees:1997,tozzi/etal:2000,iliev/etal:2002,shapiro/etal:2005}.  Each one of these methods probes different epochs and aspects of cosmic reionization: the Gunn--Peterson test is sensitive to a very small amount of residual neutral hydrogen present at the late stages of reionization ($z\\sim 6$), Ly$\\alpha$-emitting galaxies and the wavelengths of the near infrared background probe the intermediate stages of reionization ($7\\lesssim z \\lesssim 15$), the 21-cm background probes the earlier stages where the majority of the intergalactic medium is still neutral ($10\\lesssim z\\lesssim 30$), and the CMB polarization measures the column density of free electrons integrated over a broader redshift range ($z\\lesssim 20$, say). Since different datasets are complementary, one expects that cross-correlations between them add more information than can be obtained by each dataset alone. For example, the information content in the CMB and the 21-cm background cannot be exploited fully until the cross-correlation is studied: if we just extract the power spectrum from each dataset, we do not exhaust the information content in the whole dataset because we are ignoring the cross-correlation between the two. The cross-correlation always reveals more information than can be obtained from the datasets individually unless the two are perfectly correlated (and Gaussian) or totally uncorrelated.  Motivated by these considerations, we study the cross-correlation between the CMB temperature anisotropy and the 21-cm background on large scales. We show that the CMB anisotropy from the Doppler effect and the 21-cm line background can be {\\it anti}-correlated or correlated at degree scales ($l\\sim 100$), and both the amplitude and the sign of the correlation tell us how rapidly the universe reionized or recombined, and locations of the correlation (or the anti-correlation) peak(s) in redshift space tell us when reionization or recombination happened. This information is difficult to extract from either the CMB or the 21-cm data alone. Our work is different from recent work on a similar subject by \\citet{salvaterra/etal:2005}. While they studied a similar cross-correlation on very small scales ($\\sim$ arc-minutes), we focus on much larger scales ($\\sim$ degrees) where matter fluctuations are still linear and complexity due to patchy reionization is averaged out. \\citet{cooray:2004} studied higher-order correlations such as the bispectrum on arc-minute scales. For our case, however, fluctuations are expected to follow nearly Gaussian statistics on large scales, and thus one cannot obtain more information from higher-order statistics.  He also studied the cross-correlation power spectrum of the CMB and projected 21-cm maps, and concluded that the signal would be too small to be detectable owing to the line-of-sight cancellation of the Doppler signal in the CMB. However, we show that cancellation can be partially avoided by cross-correlating the CMB map with 21-cm maps at different redshifts (tomography). Prospects for the Square Kilometer Array (SKA) to measure the cross-correlation signal on degree scales are shown to be promising.  Throughout the paper, we use $c=1$ and the following convention for the Fourier transformation: \\begin{equation} f(\\hat{\\mathbf n},\\eta) = \\int \\frac{d^3{\\mathbf k}}{(2\\pi)^3}f_{\\mathbf k}e^{-i{\\mathbf k}\\cdot{\\hat{\\mathbf n}}(\\eta_0-\\eta)}, \\end{equation} where $\\hat{\\mathbf n}$ is the directional cosine along the line of sight pointing toward the celestial sphere, $\\eta$ is the conformal time, $\\eta(z)=\\int_0^t dt'/a(t')=\\int_z^\\infty dz'/H(z')$, and $\\eta_0$ is the conformal time at present. Note that \\begin{equation} \\eta_0-\\eta(z) = \\int_0^z \\frac{dz'}{H(z')}, \\end{equation} which equals the comoving distance, $r(z)=\\eta_0-\\eta(z)$, in flat geometry (with $c=1$). Also, using Rayleigh's formula one obtains \\begin{equation} f(\\hat{\\mathbf n},\\eta) = 4\\pi\\sum_{lm}(-i)^l \\int \\frac{d^3{\\mathbf k}}{(2\\pi)^3}f_{\\mathbf k} j_l[k(\\eta_0-\\eta)]Y_{lm}(\\hat{\\mathbf{n}})Y_{lm}^*(\\hat{\\mathbf k}). \\end{equation} The cosmological parameters are fixed at $\\Omega_m=0.3$, $\\Omega_b=0.046$, $\\Omega_\\Lambda=0.7$, $h=0.7$, and $\\sigma_8=0.85$, and we assume a scale invariant initial power spectrum for matter perturbations.  This paper is organized as follows. In \\S~\\ref{sec:21cmdop} and \\ref{sec:dop21cm} we derive the analytic formula for the Doppler--21-cm correlation power spectrum. Equations~(\\ref{exact}) or (\\ref{eq:exact_both}) are the main result. We then present a physical picture of the correlation and describe properties of the correlation in detail. We also discuss the validity of our assumptions and possible effects of more realistic reionization scenarios. In \\S~\\ref{sec:exp} we discuss detectability of the correlation signal with SKA before concluding in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have studied the cross-correlation between the CMB temperature anisotropy and the 21-cm background. The cross-correlation occurs via the peculiar velocity field of ionized baryons, which gives the Doppler anisotropy in CMB, coupled to density fluctuations of neutral hydrogen, which cause 21-cm line fluctuations. Since we are concerned with anisotropies in the cross-correlation on degree angular scales ($l\\sim 100$), which correspond to hundreds of comoving Mpc at $z\\sim 10$, we are able to treat density and velocity fluctuations in the linear regime.  This greatly simplifies the analysis, and distinguishes our work from previous work on similar subjects that dealt only with the cross-correlation on very small scales \\citep{cooray:2004,salvaterra/etal:2005}. Furthermore, because the 21-cm signal contains redshift information, the cross-correlation is not susceptible to the line of sight cancellation that is typically associated with the Doppler effect. Finally, because the systematic errors of the 21-cm and CMB observations are uncorrelated, the cross-correlation will be immune to many of the pitfalls associated with observing the high redshift universe in 21-cm emission, such as contamination by foregrounds\\footnote{ A potential source of foreground contamination is the Galactic synchrotron emission affecting both the CMB and 21-cm fluctuation maps; however, the amplitude of synchrotron emission in the CMB map at degree scales is much smaller than the Doppler anisotropy from reionization, and thus it is not likely to be a significant source of contamination. }. We argue that detection of the predicted cross-correlation signal provides the strongest confirmation that the signal detected in the 21-cm data is of cosmological origin. Without using the cross-correlation, it would be quite challenging to convincingly show that the detected signal does not come from other contaminating sources.  We find that the evolution of the cross-correlation with redshift can constrain the history of reionization in a distinctive way. In particular, we predict that a universe undergoing reionization results in a positive cross-correlation at those redshifts, whereas a recombining universe results in a negative correlation (this dependends on our simplified model of biased reionization -- a model in which reionization is homogeneous would imply a reversal of the sign of the correlation). Thus, the correlation promises to reveal whether the universe underwent a period of recombination during the reionization process (e.g., Cen 2003), and to reveal the nature of the sources of ionizing radiation responsible for reionization.  The signal we predict, on the order of $l^2C_l/(2\\pi)\\sim 500-1000~{\\rm \\mu K}^2$, should be easily detectable by correlating existing CMB maps, such as those produced by the WMAP experiment, with maps produced by upcoming observations of the 21-cm background with the Square Kilometer Array (SKA).  Our derivation of the cross-correlation rests upon linear perturbation theory and the reasonable assumption that the sizes of ionized regions are much smaller than scales corresponding to $l\\sim 100$.  However, assuming that the sizes of ionized regions are much smaller than the fluctuations responsible for the signal we predict is not equivalent to assuming that the ionized fraction is uniform.  Because our prediction depends on the correlation between ionized fraction and density, $P_{x\\delta}$, we have derived a simple approximate model for it (see Appendix).  In future work, we will use large-scale simulations of reionization to verify the accuracy of the relation we derive, and perhaps to refine our analytical predictions.  Whatever the result of more detailed future calculations, we are confident that the CMB Doppler-21--cm correlation will open a new window into the high redshift universe and shed light on the end of the cosmic dark ages."
    },
    "0512/astro-ph0512360_arXiv.txt": {
        "abstract": "{At which masses does the regime of globular clusters end and the one of dwarf galaxies begin? And what separates these two classes of hot stellar systems?} {We examine to what extend very massive ($>10^7 M_\\odot$) young star clusters are similar to their lower mass counter parts and to which degree they resemble other objects in their mass regime (dwarf--globular transition objects (DGTOs), ultra compact dwarf galaxies (UCDs), galaxy nuclei)} {The comparison is performed by placing the recently observed very massive young clusters onto known scaling relation defined by globular clusters (with typical masses $\\la10^6 M_\\odot$) and/or by hot stellar systems with sizes up to those of giant galaxies.} {The very massive ($\\ga10^{6.5-7} M_\\odot$) young clusters seem to show a mass--radius relation compatible with the one defined by hot stellar systems of galaxy mass. This, in turn, can explain their location on the other scaling relations investigated. It contrasts with the behaviour of the less massive young clusters and of globular clusters, which do not exhibit any mass-radius relation.  However,  the behaviour of the most massive clusters is similar to that of  most other objects in that mass regime ($10^6-10^8 M_\\odot$).} {We show that the properties of young massive clusters are compatible with other objects in the same mass regime such as DGTOs/UCDs. They present a possible direct avenue of formation for those objects, which does not require the transformation of a previously existing stellar system. Simulations and observations support the possibility of the formation of such very massive young clusters by early mergers of lower mass stellar clusters, which could explain the emergence of a mass--radius relation.} ",
        "introduction": " ",
        "conclusions": "The newest mass measurements of young clusters with masses greater than $10^7 M_\\odot$ show that these objects overlap in the scaling relations with DGTOs/UCDs and other objects in that mass regime. In particular, the most massive young clusters seem to follow the same mass--radius relation as DGTOs/UCDs and elliptical galaxies. {\\it This suggests that DGTOs/UCDs are compatible with having the same nature/origin as the most massive young clusters.} An open question is the ability for the evolved products of these massive young clusters to reproduce the high mass-to-light ratios observed for DGTOs (Ha\\c segan et al.~2005).  We can then ask: what is the formation process for these most massive young clusters? As mentioned already above, independent simulations (Bekki et al.~2004, Fellhauer \\& Kroupa 2005), as well as recent observations (Larsen et al.~2002, Minniti et al.~2004, Bastian et al.~2005c) point towards the idea that these objects could be products of early star cluster mergers, occurring in the first tens to hundred Myr. Our results are consistent with this hypothesis.  Note that {\\it late} mergers of star/globular clusters are not excluded as an alternative formation process for very massive star clusters.  Oh \\& Lin's (2000) simulations show that old star clusters can, through orbital decay, sink into the center of a dwarf galaxy and assemble to form a very massive star cluster. While it could potentially apply to (some?) DGTOs/UCDs, it cannot explain {\\it young} massive clusters such as discussed here.  Are the most massive young clusters a distinct class of objects with respect to globular or lower mass young clusters? At face value, the scaling relations appear to differ. In particular, globular clusters and young massive clusters with masses of less than $10^5 M_\\odot$ do not appear to follow a mass--radius relation (which in turn, when combined with the virial theorem, would explain why the $\\sigma$--mass and $\\langle \\Sigma_{\\rm h}\\rangle$--mass relations `bend over'). The study of star cluster complexes and their associated giant molecular clouds (Bastian et al.~2005c) showed that the emergence of a mass--radius relation seems to occur at scales between individual star clusters and cluster complexes. The complexes show a similar relation as their parent giant molecular clouds. This would strengthen the assumption that the most massive star clusters, and potentially some DGTOs/UCDs, are associated with star cluster complexes, i.e.~star cluster merger events.  As an alternative to the above scenario, one could speculate that all star cluster form with a primordial mass--radius relation, but only the most massive star clusters are able to retain it against processes that would erase it (cf.~Ashman \\& Zepf 2001, Bastian et al.~2005c). But given the lack of theoretical and observational support for this scenario, we currently favor the first hypothesis: objects in the mass range $10^6\\leq M \\leq 10^8 M_\\odot$ are likely to be the lowest mass structures resulting from merger of stellar systems, with the merging process being at the origin of a mass-radius relation. If the most massive star clusters indeed form by mergers and the typical globular clusters not, the two populations must overlap in some mass regime ($10^5 - 10^6 M_\\odot$?). The caveat is that the scenario cannot answer the question why globular clusters do not show a mass--radius relation, in contrast to the molecular clouds in which they are thought to be formed. But this scenario, in which young massive star clusters are the product of star cluster merger events, would explain their presence at the lowest mass end of the galaxy scaling relations extending from the dwarf galaxy regime to giant ellipticals."
    },
    "0512/astro-ph0512630_arXiv.txt": {
        "abstract": "The recent discovery of high frequency oscillations in giant flares from SGR 1806-20 and SGR 1900+14 may be the first direct detection of vibrations in a neutron star crust. If this interpretation is correct it offers a novel means of testing the neutron star equation of state, crustal breaking strain, and magnetic field configuration.  Using timing data from RHESSI, we have confirmed the detection of a 92.5 Hz Quasi-Periodic Oscillation (QPO) in the tail of the SGR 1806-20 giant flare. We also find another, stronger, QPO at higher energies, at 626.5 Hz.  Both QPOs are visible only at particular (but different) rotational phases, implying an association with  a specific area of the neutron star surface or magnetosphere.  At lower frequencies we confirm the detection of an 18 Hz QPO, at the same rotational phase as the 92.5 Hz QPO, and report the additional presence of a broad 26 Hz QPO. We are however unable to make a robust confirmation of the presence of a 30 Hz QPO, despite higher countrates.  We discuss our results in the light of neutron star vibration models. ",
        "introduction": "The Soft Gamma Repeaters (SGRs), objects that exhibit recurrent bouts of gamma-ray flare activity,  are thought to be magnetars - neutron stars with magnetic fields greater than $10^{14}$ G \\citep{dun92, tho95, woo04}. On rare occasions SGRs exhibit giant flares, hugely energetic events with peak fluxes in the range $10^{44} - 10^{46} \\mathrm{ergs~s}^{-1}$.    The giant flares consist of a short spectrally hard initial peak, followed by a softer decaying tail lasting a few hundred seconds.  Pulsations with periods of several seconds are visible in the tail, and reveal the neutron star spin period. Their presence is thought to be due to a fireball of ejected plasma, trapped near the stellar surface by the strong magnetic field \\citep{tho95}.  Three giant flares have been detected since the advent of satellite-borne high energy detectors: in 1979, from SGR 0526-66 \\citep{maz79}; in 1998, from SGR 1900+14 \\citep{hur99}; and in 2004 from SGR 1806-20 \\citep{ter05, pal05}.  The catastrophic magnetic instability that powers the giant flares is thought to be associated with large-scale fracturing of the neutron star crust \\citep{flo77, tho95, tho01, sch05}. This will almost certainly excite global seismic vibrations \\citep{dun98}:  terrestrial seismologists regularly observe such global modes after large earthquakes such as the 2004 Sumatra-Andaman event \\citep{par05}.  The 2004 flare from SGR 1806-20 was the most energetic ever recorded. Analysis of data from the {\\it Rossi X-ray Timing Explorer} (RXTE) by \\citet{isr05} revealed a transient 92.5 Hz Quasi-Periodic Oscillation (QPO) in the tail of the flare, associated with a particular rotational phase.  The presence of 18 and 30 Hz features was also suggested, although with lower significance. \\citet{isr05} suggested that the 30 Hz and 92.5 Hz QPOs might be toroidal shear vibrations of the crust \\citep{sch83, mcd88, str91, dun98, pir05}. These modes are particularly promising in terms of both ease of excitation and coupling to the magnetic field to modulate the X-ray lightcurve \\citep{bla89}.  Motivated by this work, \\citet{str05} re-analysed RXTE data for the 1998 flare from SGR 1900+14.  QPOs were found at 28, 54, 84 and 155 Hz, again associated with a particular rotational phase.  These QPOs can plausibly be identified with a series of toroidal modes.  The association with a rotational phase away from the main peak (where emission from the surface is obscured by the trapped fireball) implies that we are seeing emission from a particular area of the stellar surface or bundle of magnetic field footpoints.  The frequencies of crustal modes such as the toroidal modes depend on neutron star mass and radius, crustal rigidity, and magnetic field configuration \\citep{dun98}. For this reason the high frequency QPOs offer a powerful new diagnostic of neutron star properties provided that an accurate mode identification can be made (see \\citet{str05} for an initial attempt to constrain SGR parameters).  In this paper we present a timing analysis of the SGR 1806-20 hyperflare using data from the {\\it Ramaty High Energy Solar Spectroscopic Imager} (RHESSI), and discuss our results in the light of the crustal vibration model. ",
        "conclusions": "The strong rotational-phase dependence of the QPOs in the tail of the SGR 1806-20 hyperflare provides compelling evidence that all of the oscillations are assocated with particular regions of the stellar surface or magnetosphere.  In this section we discuss the candidate mechanisms in more detail.  \\citet{isr05} identified the 92.5 Hz QPO with the $l=7, n=0$ toroidal mode of the crust.  This intepretation was supported by the apparent detection of a QPO at 30 Hz, the expected frequency of the fundamental $l=2$ mode.  Although we cannot confirm the detection of the 30 Hz QPO, the interpretation of the 92.5 Hz QPO is probably secure. If the magnetic field of $\\sim 10^{15}$ G inferred from timing analysis \\citep{woo04} is accurate, then for the 92.5 Hz QPO to be an $l=6$ toroidal mode the stellar mass would have to be $< 1M_\\odot$ for even the softest equation of state.  To be an $l=8$ mode, the neutron star mass would have to exceed that of SGR 1900+14 by $>0.3 M_\\odot$ (see the discussion in \\citet{str05}). An alternative to the toroidal modes, with a frequency in the correct range, is the crustal interface mode, but its behavior in the presence of strong magnetic fields requires further study \\citep{pir05a, pir05}.  The detection of a QPO at 626.5 Hz is exciting, as there are several crustal modes that could have frequencies in this range.  The most likely candidate is the $n=1$ toroidal mode; the detected frequency agrees extremely well with the most recent models \\citep{pir05}.  This would be fascinating in terms of energetics; the energy required to excite $n>0$ modes is orders of magnitude larger than that required to excite an $n=0$ mode, a testament to the strength of this particular flare.  Other candidates include the crustal spheroidal modes or the crust/core interface modes \\citep{mcd88}, although the effect of a strong magnetic field on these modes remains to be calculated. If this mode is indeed a higher radial overtone, it could allow us to determine the depth at which the fracture occurred in the crust.  The nature of the 18 and 26 Hz QPOs is less clear.  The frequencies are too low for the fundamental crustal toroidal mode unless the neutron star parameters are very extreme.  A torsional mode of the core, restored by the poloidal magnetic field \\citep{tho01}, was suggested by \\citet{isr05} as a candidate mechanism for the 18 Hz QPO. Whether such modes could explain both of the QPOs is a matter for further study.  The fact that the various QPOs are strongest at different times, suggests that we may be seeing two separate fracture/reconnection events.  The first, associated with the main flare, excites the 626.5 Hz QPO.  The second, associated with the late time boost in unpulsed emission, excites the 92.5 Hz QPO.   The idea that we might be seeing separate fracture sites is given additional weight by the observation that the two QPOs are strongest at different rotational phases.  Such sequential rupturing is often observed in terrestrial seismology. It is interesting that the 18 and 26 Hz QPOs are strongest at the rotational phase where the 92.5 Hz QPO subsequently appears.  One might speculate that the magnetospheric oscillations associated with the low frequency QPOs (which are triggered by the main flare) slowly weaken the crust at this point and trigger the second fracture.  Understanding the means by which the modes couple to the magnetic field, and hence modulate the X-ray lightcurve, is now critical. The fact that the modes appear in different energy bands is intriguing, as is the fact that the 626.5 Hz QPO appears at a rotational phase when a large part of the surface should be obscured by the trapped fireball. Much theoretical effort is clearly still required, but prospects for neutron star asteroseismology are bright."
    },
    "0512/astro-ph0512406_arXiv.txt": {
        "abstract": "We present a model independent method of reconstructing the Lagrangian for the k-essence field driving the present acceleration of the universe. We consider the simplest k-essence model for which the potential is constant. Later we use three parametrizations for the Hubble parameter $H(z)$, consistent with recent the SN1a data, to yield the Lagrangian $F$. Our reconstruction program does not generate any physically realistic Lagrangian for models that allow phantom crossing, whereas models without phantom crossing, yield well behaved Lagrangians. ",
        "introduction": "One of the most interesting observational discoveries in the past decade has been the evidence that the expansion of the universe is speeding up rather than slowing down. A currently accelerating universe is strongly favored by type Ia supernova data \\cite{super}. Recent observations of cosmic microwave background radiation (CMBR) by the Wilkinson Microwave Anisotropy Probe (WMAP) experiment \\cite{wmap} and of the  large scale structure by redshift surveys e.g the Sloan Digital Sky Survey (SDSS) \\cite{sdss} and the 2-Degree Field (2df) redshift survey \\cite{2df} have further strengthened this result. This result together with the observational evidence that we are living in a low matter density (around $30\\%$ of the critical density) flat ($\\Omega_{total} =1$) universe, suggests the existence of a relatively smooth component called ``dark energy'' which dominates the energy density of the universe and also has a large negative pressure.  Although a simple cosmological constant $\\Lambda$ ($\\omega = -1$) can serve the purpose of this dark energy, it faces a  serious problem of fine tuning (for details, see recent review by Padamnabhan \\cite{paddy}). The alternative possibility called quintessence \\cite{quint}, is a varying  $\\Lambda$ model, where dark energy arises from a minimally coupled scalar field rolling over its potential. For a sufficiently flat potential, the potential energy term dominates and the scalar field can mimic $\\Lambda$.  Recently, an alternative possibility, that of an effective scalar field theory governed by a Lagrangian containing a non-canonical kinetic term (${\\cal{L}} = - V(\\phi)F(X)$, where $ X = (1/2) \\partial_{\\mu}\\phi \\partial^{\\mu}\\phi$) has been proposed. Such a model can lead to a late time acceleration in the present universe and is named as ``k-essence'' \\cite{kes}. It is worth noting that the Generalized Chaplygin Gas model \\cite{gcg} and the tachyon dark energy models \\cite{tirtha} may be seen as  special cases of k-essence (e.g k-essence with a constant potential can mimic simple Chaplygin gas (CG) model \\cite{frolov}).  From the theoretical side, it is not only important to check whether dark energy is a constant or dynamical, but also if dynamical, it is equally important to know the Lagrangian for the underlying field. For a minimally coupled scalar field model, a number of studies have been done for reconstructing the potential for this scalar field in a model independent way by using observational data \\cite{reconsq}. Recently Tsujikawa has studied the reconstruction of the potential function $V(\\phi)$ in the k-essence Lagrangian once one assumes a certain form for the kinetic energy function $F(X)$ \\cite{tsuji}. But until now, there is no attempt in the literature to reconstruct in a model independent way the function $F(X)$ in the k-essence Lagrangian.  In this work, we propose a algorithm (similar to that proposed by Simon et al. \\cite{simon} for a minimally coupled scalar field model) to reconstruct $F(X)$ in a model independent way. For present purposes, we restrict ourselves to the case where the potential $V(\\phi)$ is a constant. Using this algorithm and three specific ansatzes for the Hubble parameter $H(z)$ (consistent with the supernova observations), we reconstruct  $F(X)$ in the redshift range between $z=0$ and $z=1.7$ which is the current range for the supernova data. ",
        "conclusions": "In this work, we have proposed a method for reconstructing the k-essence Lagrangian for constant potential. In doing so, we have defined the parameters $\\epsilon_{n}$, which are similar to the horizon-flow parameters for inflation. Using these parameters, we have related $F(z)$ and $F'(z)$ to the directly observable quantities like $\\Omega_{m}$, $q(z)$, $H(z)$. Equations (11) and (19) are the core of our reconstruction program. By knowing $H(z)$ from distance-redshift measurements like SNIa observations, one can reconstruct $F$ as a function of redshift $z$ and $X$, the kinetic energy. Later, we have used three well known ansatz for $H(z)$, in order to reconstruct $F$ for $z=0$ to $z=1.7$ which is the current redshift range for SNIa data. We should emphasize that our reconstruction program only involves that portion of $F(X)$ over which the field evolves to give the required $H(z)$ and is insensitive to certain ranges in $X$.   Our study shows that in order to have a physically realistic k-essence Lagrangian $F(X)$, the crossing of the phantom boundary $\\omega = -1$ is forbidden. For ansatzes 1 and 2, which allow the phantom crossing, the reconstructed $F$'s are physically unrealistic double-valued functions of $X$. This result is consistent with earlier investigations with similar conclusion \\cite{vik, doran, neto}. For ansatz 3, which is also consistent with SN1a data but does not allow phantom crossing, the reconstructed $F(X)$ is a well-behaved, single valued function. The result also shows that one has to be careful while choosing different ansatzes for $H(z)$ or $\\omega(z)$ for dark energy to fit different observational data. While some of them are perfectly consistent with observation, it may not be possible to relate them to any physically realistic theory."
    },
    "0512/astro-ph0512295_arXiv.txt": {
        "abstract": "We introduce and discuss the properties of a theoretical $(B-K)$-$(J-K)$ integrated colour diagram for single-age, single-metallicity stellar populations. We show how this combination of integrated colours is able to largely disentangle the well known age-metallicity degeneracy when the age of the population is greater than $\\sim$ 300~Myr, and thus provides valuable estimates of both age and metallicity of unresolved stellar systems. We discuss in detail the effect on this colour-colour diagram of $\\alpha$-enhanced metal abundance ratios (typical of the oldest populations in the Galaxy), the presence of blue horizontal branch stars unaccounted for in the theoretical calibration, and of statistical colour fluctuations in low mass stellar systems. In the case of populations with multiple stellar generations, the luminosity-weighted mean age obtained from this diagram is shown to be heavily biased towards the youngest stellar components.  We then apply this method to several datasets for which optical and near-IR photometry are available in the literature.  We find that LMC and M31 clusters have colours which are consistent with the predictions of the models, but these do not provide a sensitive test due to the fluctuations which are predicted by our modelling of the Poisson statistics in such low-mass systems. For the two Local Group dwarf galaxies NGC185 and NGC6822, the mean ages derived from the integrated colours are consistent with the star formation histories inferred independently from photometric observations of their resolved stellar populations.  The methods developed here are applied to samples of nearby early-type galaxies with high quality aperture photometry in the literature. A sample of bright field and Virgo cluster elliptical galaxies is found to exhibit a range of luminosity-weighted mean ages  from 3 to 14~Gyr, with a mean of $\\sim$8~Gyr, independent of environment, and mean metallicities at or just above the solar value. Colour gradients are found in all of the galaxies studied, in the sense that central regions are redder.  Apart from two radio galaxies, where the extreme central colours are clearly driven by the AGN, and one galaxy which also shows a radial age gradient, these colour changes appear consistent with metallicity changes at a constant mean age.  Finally, aperture data for five Virgo early-type dwarf galaxies shows that these galaxies appear to be shifted to lower mean metallicities and lower mean ages (range 1 to 6 Gyr) than their higher luminosity counterparts. ",
        "introduction": "Within the currently favoured $\\Lambda$CDM models of galaxy and large scale structure formation, dwarf galaxies play a crucial role. The first stars are predicted to form in small dark matter halos of mass $\\approx 10^{6}$ $M_{\\odot}$ at redshifts of $z \\approx 20-30$ (see Bromm and Larson 2004). These small dark matter halos subsequently form the building blocks from which the larger galaxies we see around us today are assembled (Kauffman et al. 1993; Cole et al. 2000; Mathis et al. 2002). As with every building site many of the smaller component parts are left behind once the larger structures are formed and so the $\\Lambda$CDM model predicts that there should be many small dark matter halos around us today.  If this debris of small dark matter halos can be associated with dwarf galaxies then there is a large discrepancy between what is observed and what is predicted by theory, this is known as the sub-structure problem (Moore et al. 1999; Klypin et al. 1999). The standard $\\Lambda$CDM model, without dwarf galaxy formation suppression processes, predicts many more small dark matter halos (with a low mass slope of the mass function of $\\alpha \\approx -2$) compared to observations of the faint end of the global luminosity function of galaxies ($\\alpha \\approx -1.2$; Blanton et al. 2001, Norberg et al., 2002). There have been numerous suggestions for a solution to this problem. These typically involve the prevention of gas falling into small dark matter halos (Efstathiou 1992), the removal of gas via supernova winds from the first generation of stars (Dekel and Silk 1986) or the `squelching' of dwarf galaxy formation by ionising photons (Tully et al. 2002). These mechanisms prevent or severely inhibit star formation in small halos; the halos still exist, but they remain hidden from us. Interest in the possiblity that star formation in dwarf galaxies has in some way been delayed compared to larger galaxies has come from the extensive studies made using Sloan Digital Sky Survey data. This seems to indicate that it is the low luminosity galaxies that contain the youngest stars (Kauffmann et al., 2003). It is difficult to see how this `downsizing' of the galaxy population fits in with the standard $\\Lambda$CDM model (Kodama et al. 2004), but it is an observation we can test with the methods we discuss in this paper. Recently other solutions have been suggested, for example by adjusting the initial dark matter fluctuation spectrum, something suggested by combining the WMAP, 2dF and Lyman-$\\alpha$ forest observations (Spergel et al. 2003), small dark matter halos themselves become rare.  By studying in detail the sub-structure problem in different galaxy environments we can hope to gain an insight into the mechanisms that prevent dark matter halos revealing themselves to us as dwarf galaxies or comment on how universally rare small halos might be. We have previously quantified the relative numbers of dwarf galaxies in different environments and find many more dwarf galaxies in the Virgo and Fornax clusters than can be found in the general field (Phillipps et al. 1998a; Sabatini et al. 2003, 2004; Roberts et al. 2004, 2005; Davies et al. 1988; Kambas et al. 2000). The general field result is also consistent with observations of the dwarf galaxy population of the Local Group (Mateo 1998; Pritchet \\& van den Bergh 1999). It appears that global star formation suppression mechanisms do not work. The environment plays a major role in either producing a large number of additional dwarf galaxies or in revealing more of the existing dark matter halos to us.  The expected properties of the first galaxies to form within the $\\Lambda$CDM model seem to be more closely matched to the properties of globular clusters (a star formation event followed by almost total gas expulsion) than it does to the more complicated star formation histories of dwarf galaxies. Although dwarf galaxies appear to be relatively simple stellar systems (there is little evidence for spiral stellar density waves, bars or jets, but see Barazza et al. 2002) they can be more susceptible to exterior influences. In the main dwarf galaxies show a morphology density relation with the gas poor dSph galaxies being closer to the bright galaxies in the Local Group and closer to the centres of galaxy clusters than are the dIrr galaxies (Ferguson \\& Sandage 1989; Mateo 1998; Boyce et al. 2001; Sabatini et al. 2003; Roberts et al. 2004). Within a cluster environment, or when they are close to a more massive galaxy, dwarf galaxies may be subject to ram pressure and/or tidal stripping, tidal destruction or `harassment' processes (for a review see Sabatini et al. 2004).  Currently we essentially only have detailed star formation histories for Local Group dwarf galaxies because these are the only ones for which we have resolved observations of their stellar populations. These have been interpreted using stellar colour magnitude diagrams and models that implement different stellar ages and metallicities (Grebel \\& Stetson 1999; Smecker-Hane et al. 1996; Gallart et al. 1999; Mighell and Rich 1996; Skillman et al. 2003). Each of the Local Group dwarf galaxies has a quite different star formation history when studied in detail, with stellar populations covering a wide range of ages and metallicities. In general the dSph galaxies have no young stars; the stellar population is dominated by stars that are older than 10 Gyr, but there are some intermediate age stars (1-10 Gyr) indicating that star formation continued for many Gyr after their formation (Grebel 2002). The current, most widely used, $\\Lambda$CDM solution to the substructure problem (that small galaxies in the early Universe should experience a single catastrophic episode of star formation) is not supported by observations of the star formation histories of Local Group dwarf galaxies or by the large numbers of dwarf galaxies found in clusters.  Although subject to similar physical processes dwarf galaxies in clusters may have very different origins and histories to those in the Local Group. Within the cluster environment there are many more dwarf galaxies per giant (Phillipps et al. 1998b, Sabatini et al. 2003), so the clusters are not being assembled from groups like the Local Group. Also unlike the Local Group there are many early type (dE) dwarf galaxies that are not companions of the giant galaxies (Sabatini et al. 2003, 2004; Roberts et al. 2004). Ideally we would like to study the star formation histories of these galaxies in the same detail as has been done for the Local Group galaxies to see if they are consistent with each other and with the properties predicted by $\\Lambda$CDM - predominantly old, gas poor, low metallicity galaxies. Currently this is not possible for dwarf galaxies in the two nearest clusters to us (Virgo and Fornax) for two reasons. Firstly, because the bulk of the stellar population cannot be resolved at these distances (although the resolution of galaxies at these distances is a key science driver for the proposed Extremely Large Telescopes). Secondly, many of the dwarf galaxies we are particularly interested in are of such low surface brightness (LSB) that it is impossible to obtain spectra and derive line indices (Worthey 1994, Vazdekis and Arimoto 1999, Kuntschner 2000). Thus, at present we are forced to use population indicators that can be applied to completely unresolved LSB stellar populations.  In the present study we explore the use of optical/near-IR colours as diagnostics of luminosity-weighted mean ages and metallicities of stellar populations.  We demonstrate that such colours have similar diagnostic power to spectral line indices, while being much less demanding in terms of observing time, data reduction and analysis. The development of such indicators is particularly timely with the recent release of the full Two-Micron All-Sky Survey (Jarrett et al. 2000; henceforth 2MASS), and the prospect of deeper surveys such as the United Kingdom Infrared Deep Sky Survey (UKIDSS) recently commenced with the United Kingdom Infrared Telescope (UKIRT).  Model predictions of optical/near-IR colours have previously been presented by several groups (e.g. Fioc and Rocca-Volmerange 1997; Smail et al. 2001, Girardi et al.  2002; Bruzual and Charlot 2003) and observational studies have already highlighted the potential of such measurements for breaking the age-metallicity degeneracy, which is the first requirement of any successful population diagnostic method. Peletier, Valentijn \\& Jameson (1990) used $(U-V)$, $(B-V)$ and $(V-K)$ colours to constrain the ages and metallicities of 12 elliptical galaxies, while Peletier \\& Balcells (1996) demonstrated that $(U-R) - (R-K)$ colour-colour plots gave useful information on the mean ages of galaxy bulges and disks.  Bell et al. (2000) and Bell \\& de Jong (2000) used $(B-R)$ and $(R-K)$ colours in the analysis of the star formation histories of low surface brightness and spiral galaxies respectively. Puzia et al. (2002) and Hempel \\& Kissler-Patig (2004) have recently demonstrated the power of $(V-I) - (V-K)$ colour-colour diagrams applied to analysis of the star-formation histories of unresolved extragalactic globular clusters.  In this paper we first develop our own predictions of the optical/near-IR colours of stellar populations as a function of age and metallicity, using models developed by two of the authors (MS and SC).  We identify the $(B-K)$-$(J-K)$ colour-colour plane as a particularly sensitive one for breaking age-metallicity degeneracy, and use photometry to test these predictions for systems (globular clusters and two dwarf galaxies in the Local Group) with existing independent measurements of age and metallicity.  We also apply these methods to a small number of Virgo cluster elliptical and dwarf elliptical galaxies which have sufficiently high quality photometry in the 2MASS catalogues and $B$-band photometry in matched apertures, in the first stage in a study of the star formation history of cluster galaxies. ",
        "conclusions": "We will summarise briefly our conclusions on the nature of the stellar populations of the various galaxy samples as revealed by their optical/near-infrared colours.  It is useful to remember at this point that the ages and metallicities derived are luminosity-weighted mean values, and we have shown that, as for spectroscopic indices, the results can be heavily skewed by the presence of even a small (a few per cent by mass) young stellar population in an otherwise old galaxy.  For the bright elliptical galaxy sample of Michard (2005), we find that the integrated colours within the effective radius for these galaxies reveal stellar populations with mean ages averaging about 8~Gyr, but with a real scatter, larger than the measurement errors, covering the range from 3--14~Gyr.  Mean [Fe/H] values are just above solar, but again there is a real scatter, from just below solar to just above [Fe/H]$=$0.4.  The Virgo cluster galaxies contained within Michard's sample show no statistically significant differences in their integrated parameters than the sample as a whole, so there is no evidence for strong environmental effects on either mean stellar ages or metallicities.  A more detailed study of population gradients within 8 Virgo cluster giant galaxies and 5 Virgo cluster dwarf galaxies reveals a strong trend for central regions of all galaxies to have higher metallicities than the galaxies as a whole.  Evidence for age gradients is much more marginal, but one giant galaxy shows evidence for a younger central population, while the dwarfs show weak evidence for age gradients in both directions.  Considering the 5 Virgo dwarf galaxies together, we find that they have younger mean stellar ages than their more luminous counterparts. Clearly, this is too small a sample to be drawing general results about the dwarf galaxy population as a whole, but it should be noted that such a result would be consistent with recent observational results by Caldwell et al. (2003), using spectral indices of low-luminosity galaxies in Virgo and lower density environments, and Kauffmann et al. (2003), who perform a statistical study of the star formation histories of galaxies from the Sloan Digital Sky Survey.  There are two conclusions we can make from this preliminary study. Firstly, if we assume that all these galaxies contain a simple stellar population (SSP), then quite contrary to the hierarchical model of galaxy formation the smallest galaxies have not formed first at redshifts of 6-10 (ages greater than 12~Gyr for standard $\\Lambda$CDM) but, this does fit in with the `downsizing' scenario discussed in the introduction. Secondly, if these galaxies do not contain a SSP but do have a small fraction of young stars then the feedback mechanisms discussed in the introduction are inefficient - these galaxies have retained at least some gas to enable them to continue to form stars over long periods of time (see also the model by Kobayashi, 2005, which has this property).  This second interpretation fits in very well with observations of the resolved stellar populations found in Local Group dwarf galaxies (Grebel \\& Stetson 1999; Smecker-Hane et al. 1996; Gallart et al.  1999; Mighell \\& Rich 1996) and discussed in the introduction.   Unfortunately, it is only possible to study the brightest of the Virgo dwarf elliptical galaxies using 2MASS photometry.  However, with new data becoming available, e.g.  from the UKIDSS project, it should be possible to derive the required colours for some hundreds of Virgo cluster dwarfs, pushing significantly further down the luminosity function."
    },
    "0512/astro-ph0512589_arXiv.txt": {
        "abstract": "We present the initial results of a 3-mm spectral line survey towards 83 methanol maser selected massive star-forming regions. Here we report observations of the J\\,=\\,5\\,--\\,4 and 6\\,--\\,5 rotational transitions of methyl cyanide (\\chthreecn) and the J\\,=\\,1\\,--\\,0 transition of \\hcop~and \\hthirteencop.  \\chthreecn~emission is detected in 58 sources (70\\,\\% of our sample). We estimate the temperature and column density for 37 of these using the rotational diagram method. The temperatures we derive range from 28\\,--\\,166\\,K, and are lower than previously reported temperatures, derived from higher J transitions. We find that \\chthreecn~is brighter and more commonly detected towards ultra-compact H{\\scriptsize II} (UCH{\\scriptsize II}) regions than towards isolated maser sources. Detection of \\chthreecn~towards  isolated maser sources strongly suggests that these objects are internally heated and that \\chthreecn~is excited prior to the UCH{\\scriptsize II} phase of massive star-formation.  \\hcop~is detected towards 82 sources (99\\,\\% of our sample), many of which exhibit asymmetric line profiles compared to \\hthirteencop. Skewed profiles are indicative of inward or outward motions, however, we find approximately equal numbers of red and blue-skewed profiles among all classes. Column densities are derived from an analysis of the \\hcop~and \\hthirteencop~line profiles.  80 sources have mid-infrared counterparts: 68 seen in emission and 12 seen in absorption as `dark clouds'. Seven of the twelve dark clouds exhibit  asymmetric \\hcop~profiles, six of which are skewed to the blue, indicating infalling motions. \\chthreecn~is also common in dark clouds, where it has a 90\\,\\% detection rate. ",
        "introduction": "Identifying young massive stars ($>$\\,8\\,M$_{\\odot}$) in their early evolutionary phases is an important step in attempting to understand massive star formation and its effect on the Galactic ecology.  Massive stars begin their lives at the centre of dense contracting cores, embedded in giant molecular clouds (GMCs). At the earliest stages of their evolution they heat the surrounding dust, and are visible as luminous sub-millimetre (mm) and far-infrared sources \\citep{Kurtz2000}. High densities existing prior to the final collapse cause simple organic molecules present in the interstellar medium to freeze onto the dust  mantles. Grain surface chemistry leads to the production of heavy ices such as methanol (\\chthreeoh). As the protostar evolves, heat and radiation evaporate these species, expelling them into gas phase where they fuel a rich `hot core' chemistry \\citep{VanDishoeck1998}. Eventually, the embedded protostar ionises the surrounding environment, forming a ultra-compact H{\\scriptsize II} (UCH{\\scriptsize II}) region and complex molecules are destroyed owing to continued exposure to heat and radiation.  Targeted surveys towards known star forming regions have suggested a close relationship between class II methanol masers, hot cores and UCH{\\scriptsize II} regions \\citep{Menten1991,Minier2001}. Parkes 64-m and Australia Telescope Compact Array (ATCA) observations by \\citet{Walsh1998,Walsh1999} have shown that only 25\\,\\% of methanol masers are directly associated with radio continuum emission, tracing UCH{\\scriptsize II} regions. Subsequent follow-up observations towards the {\\it isolated} masers confirm that all maser sites are associated with luminous sub-mm continuum emission \\citep{Walsh2003}.  In the present paper we investigate the link between methanol masers, UCH{\\scriptsize II} regions and the chemically active `hot molecular core' (HMC) phase of massive star formation. We present the initial results of a single-dish molecular line survey of 83 methanol maser selected massive star forming regions.   \\subsection{\\chthreecn~Spectroscopy} The primary chemical tracer in this survey is methyl cyanide (\\chthreecn). \\chthreecn~is a good tracer of the conditions found in HMCs owing to its favourable abundance and excitation in warm ($>$\\,100\\,K) and dense ($>$\\,10$^5$\\,\\cmmthree) regions. It is thought to form through reactions between species evaporated from dust grain mantles (e.g. \\citealt*{Millar1997b}) or through gas phase chemistry in the envelope around massive young stars (see \\citealt*{Mackay1999}). Grain-surface chemistry has also been invoked to explain the enhanced abundances observed in HMCs, however, in all formation scenarios \\chthreecn~is evident in observations after the core temperature rises past $\\sim$\\,90\\,K.  \\chthreecn~is a member of the ${\\rm c_{3v}}$ group of symmetric tops, whose rotational energy levels may be described by two quantum numbers: J, the total angular momentum and K, the projection of J along the axis of symmetry. Individual J\\,$\\rarr$\\,(J\\,--\\,1) transitions are grouped into `rotational ladders' labelled by their K values (see \\citet*{LorenMundy1984} for a detailed description). For each J\\,$\\rarr$\\,(J\\,--\\,1) transition, selection rules prohibit radiative transitions between the K ladders and their relative populations are determined exclusively by collisional excitation. Assuming local thermal equilibrium (LTE) and optically thin lines, the relative intensities of the K components yield a direct measure of the kinetic temperature and column density.  The energy spacings between individual J levels are almost independent of K ladder, however, increasing centrifugal distortion causes successive K components to shift to progressively lower frequencies. The offset in frequency is slight, and the K components of a particular J\\,$\\rarr$\\,(J\\,--\\,1) transition may be observed simultaneously in a single bandpass, minimising errors in their relative calibration.  Spin statistics of the hydrogen nuclei divide \\chthreecn~into two spin states, dubbed A and E. Energy levels with ${\\rm K=3n,~n=0,1,2\\,\\ldots}$ belong to the A state, while those with ${\\rm K\\neq 3n,~n=0,1,2\\,\\ldots}$ belong to the E state. The A states have twice the statistical weight of the E states. Neither radiation nor collisions convert between states; if formed in equilibrium conditions the A\\,/\\,E abundance ratio is expected to be $\\sim$\\,1 (\\citealt{Minh1993}).   \\subsection{\\hcop~Spectroscopy} \\hcop~and its isotopomer, \\hthirteencop, were also observed in order to probe the kinematics of the extended envelopes and as a diagnostic of the optical depth towards the cores. A comparison between the generally optically thick \\hcop~line and the optically thin \\hthirteencop~line yields information on the bulk motions of gas in the regions.  \\hcop~has a large dipole moment and is a highly abundant molecule, with abundance especially enhanced around regions of higher fractional ionisation. \\hcop~is also enhanced by the presence of outflows where shock generated radiation fields are present (\\citet*{Rawlings2000}; \\citet{Rawlings2004}). It shows saturated and self absorbed line profiles around massive star-forming regions and is a good tracer of the dynamics in the vicinity of young protostellar objects (e.g. \\citealt*{DeVries2005}). The emission from its isotopomer, \\hthirteencop, is generally optically thin and traces similar gas densities to \\hcop~at $n\\sim 10^5$\\,cm$^{-3}$.  Table~\\ref{tab:transitions} presents the details of the observed transitions. All molecular constants are in SI units and come from the catalog of \\citet{Pickett1998}.   We justify the source selection criteria and derive the distances and luminosities in the following section. In \\S\\,3 we describe the observations and data reduction methodology. In \\S\\,4, we present the initial results of the line survey and describe the data. We outline the analysis techniques used and present derived physical parameters in \\S\\,5, and in \\S\\,6 we analyse and discuss general trends in the data and test the validity of our results. Finally, in \\S\\,7 we conclude with a summary of our investigations and suggest future avenues of investigation. The full set of spectra, MSX images and related plots are presented as an online addition to this paper. ",
        "conclusions": "We detect \\chthreecn~towards all classes of sources: UCH{\\scriptsize} II regions, isolated masers and maserless cores. This result confirms a direct link between each of these objects and massive star formation. In the following analysis we will discuss the implications for a possible evolutionary sequence.   \\subsection{The presence of an UCH{\\scriptsize II} region} The \\chthreecn~detection rate towards maser-sources associated with an UCH{\\scriptsize II} region (radio-loud), compared to isolated maser sources (radio-quiet) is strikingly different. \\begin{figure*} \\begin{center} \\includegraphics[height=6.9cm, trim=0 -0 -0 0]{figs/fig_5.epsi} \\caption{~{\\bf  A}) Distributions of integrated \\chthreecn\\,(5\\,--\\,4) flux, summed over all detected K-components, for sources with UCH{\\scriptsize II} regions within the Mopra beam (bottom) and those without (top). The shaded bar represents the number of \\chthreecn~non-detections for each population. A KS-test yields a maximum difference of 0.48 between the distributions, with a probability of 0.13\\,\\% of being drawn from the same parent distribution. {\\bf B}) Distributions of integrated \\hthirteencop\\,(1\\,--\\,0) flux for the same populations. Lines associated with UCH{\\scriptsize II} regions are markedly brighter than those without. A KS-test yields a difference of 0.55 between the distributions, with an associated probability of 0.02\\,\\% of being drawn from the same parent distribution. A similar result has been found for the \\hcop~lines.} \\label{fig:ch3cna_uchii} \\end{center} \\end{figure*}  Figure~\\ref{fig:ch3cna_uchii}-A shows the distributions of integrated \\chthreecn\\,(5\\,--\\,4) intensity for the two  classes of source, including non-detections, which are summed in the shaded bar. \\chthreecn~was detected towards 18 out of the 19 radio-loud sources (95\\,\\%) compared to only 40 out of 64 radio-quiet sources (63\\,\\%). Furthermore, the average integrated \\chthreecn\\,(5\\,--\\,4) intensity measured towards the radio-loud sources is 7.4\\,K\\,\\kms, stronger than the average measured towards the radio-quiet sources, which is 3.5 K\\,\\kms. We use a Kolomogorov-Smirnov (KS) test to measure the difference between the distributions of the two classes. Non-detections are included as upper limits, measured from the RMS noise in the spectra, and by assuming K\\,=\\,0 and K\\,=\\,1 \\chthreecn~components of linewidth 4\\,\\kms~are present at a 3\\,$\\sigma$ level. Using this conservative estimate a KS-test indicates there is only a 0.13\\,\\% probability of the samples being drawn from the same population.  Probabilities below $\\sim$\\,1\\,\\% indicate that the distributions are significantly different.  A similar result applies for \\hthirteencop. Figure~\\ref{fig:ch3cna_uchii}-B shows the distributions of \\hthirteencop~intensity for the two classes of objects. The mean intensity for the radio-loud sources is 4.3 K\\,\\kms~compared to 1.9 K\\,\\kms~for the radio-quiet sources. A KS-test returns a probability of 0.02\\,\\% that the distributions are similar.  We find the \\hcop~lines are also brighter towards radio-loud sources, however we do not consider them in this analysis as they show evidence of optical depth effects which may confuse their interpretation.  \\subsubsection{The effect of distance on the results} A higher \\chthreecn~detection rate towards UCH{\\scriptsize II} regions is consistent with a fall-off in flux with distance in the original radio survey. Our classification of sources into radio-loud and radio-quiet regions may be artificial, and undetected UCH{\\scriptsize II} regions may simply be a result of the sources being further away. Closer examination shows this not to be the case. The distribution of distances to the two populations is shown in Figure~\\ref{fig:hist_dist_uchii}. No UCH{\\scriptsize II} regions have been detected further than 9\\,kpc away, however only 9 out of 83 sources lie at greater distances. A comparison yields an 53\\,\\% probability that the classes are consistent with being drawn from the same population and thus probe the same distance range. \\begin{figure} \\begin{center} \\includegraphics[height=6.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_6.epsi} \\caption{~Distributions of distances in kiloparsecs for sources with and without associated UCH{\\scriptsize II} regions. A KS-test yields a probability of 53\\,\\% that samples are consistent with being drawn from the same population and thus probe the same distances.} \\label{fig:hist_dist_uchii} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[height=7.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_7.epsi} \\caption{~A plot of distance versus bolometric luminosity for the sample. The solid line shows the 3\\,$\\sigma$ detection limit for 8.67\\,GHz radio continuum emission, assuming all the Lyman-continuum flux contributes to creating an UCH{\\scriptsize II} region. The dashed line shows how the detection limit changes if 90\\,\\% of the ionising photons are absorbed by dust. UCH{\\scriptsize II} regions are represented by squares, and isolated maser sources by crosses.} \\label{fig:xy_dist_lum} \\end{center} \\end{figure} We further examine the detection limits on the \\citet{Walsh1998} 8.64 GHz radio survey from which our sample was drawn. The quoted 1\\,$\\sigma$ sensitivity is $\\sim$\\,1\\,mJy. Assuming a 3\\,$\\sigma$ detection, the limiting distance, D, in kpc at which free-free radio emission from a UCH{\\scriptsize II} region is detectable, may be calculated from \\citet*{KurtzChurchwellWood1994}: \\begin{equation} {\\rm \\left(\\frac{S_{\\nu}}{Jy}\\right)= 1.32 \\times 10^{-49} N_L \\left(\\frac{D}{kpc}\\right)^{-2} a(\\nu,T_{e})\\left(\\frac{\\nu}{GHz}\\right)^{-0.1} \\left(\\frac{T_{e}}{K}\\right)} \\end{equation} where ${\\rm \\nu}$ = 8.64 GHz is the frequency, ${\\rm T_e \\approx 10^4}$\\,K is the electron temperature, ${\\rm a(\\nu,T_{e})}$ is a factor of order unity tabulated by \\citet*{Mezger1967} and ${\\rm N_L}$ is the number of Lyman-continuum photons. Values of ${\\rm N_L}$ for early-type stars are tabulated by \\citet{Panagia1973}, however, in practise some fraction of the emitted ionising photons will be absorbed by dust before ionising the surrounding medium. In a study of UCH{\\scriptsize II} regions, \\citet{KurtzChurchwellWood1994} found that between 50\\,\\% and 90\\,\\% of the ionising photons are absorbed for the majority of their sample. Figure~\\ref{fig:xy_dist_lum} is a plot of luminosity against distance, in kpc, for all the sources whose luminosity could be determined. The solid line marks the 3\\,$\\sigma$ sensitivity limit to 8.64 GHz radio continuum, assuming no Lyman-continuum photons are absorbed, while the dotted line shows the sensitivity limit assuming 90\\,\\% of all ionizing photons are absorbed. The luminosities of our sources were determined by fitting a 2-component greybody to available MSX, SCUBA, and IRAS fluxes as discussed earlier and are accurate to a factor of $\\sim$\\,2. For those sources above the 90\\,\\% absorption limit, UCH{\\scriptsize II} regions would have been detected if they existed, and it is most unlikely that the comparison between \\chthreecn~detected for UCH{\\scriptsize II} regions and isolated maser sources is biased. We plot Figure~\\ref{fig:ch3cna_uchii_limit}-A as per Figure~\\ref{fig:ch3cna_uchii}, but include only sources luminous enough to have detectable radio-continuum emission, assuming 90\\,\\% attenuation of ionising photons. We find that the \\chthreecn~distributions are still significantly different, with only a 0.9\\,\\% probability of being drawn from the same population. \\begin{figure*} \\begin{center} \\includegraphics[height=6.9cm, trim=0 -0 -0 0]{figs/fig_8.epsi} \\caption{~As for Figure \\ref{fig:ch3cna_uchii} but limited to sources which are luminous enough to have detectable UCH{\\scriptsize II} regions, assuming no less than 90\\,\\% of the ionising photons are attenuated. {\\bf A}) A KS-test yields a maximum difference of 0.50 between the \\chthreecn~distributions, with an associated probability of 0.93\\,\\% of being drawn from the same parent distribution. {\\bf B}) Similarly the maximum difference between the \\hthirteencop~populations is 0.56, with an associated probability of 0.12\\,\\%.} \\label{fig:ch3cna_uchii_limit} \\end{center} \\end{figure*} The \\hthirteencop~intensity distribution for the two classes, filtered to the same luminosity limit, is shown in Figure~\\ref{fig:ch3cna_uchii_limit}-B. A KS-test returns a 0.1\\,\\% probability of them being drawn from the same population.  The apparent enhancement in {\\it line brightness} towards UCH{\\scriptsize II} regions could be mimicked by the effect of beam dilution. A previous high resolution survey for \\chthreecn~has found the typical size of the emitting region to be $<$\\,10$''$ \\citep{Remijan2004}. Assuming a constant spatial size, the effect of beam dilution on the \\chthreecn~brightness temperature will depend on the angular size and hence distance to the source. We have shown that the radio-loud and radio-quiet samples probe the same distance range and thus neither class is biased towards being nearer.  Figure~\\ref{fig:xy_ch3cn_dist_uchii} plots measured \\chthreecn~(5\\,--\\,4) flux against kinematic distance and we do  not find any correlation. As an additional test we can recover the line-luminosity by multiplying by the square of the distance. This has the effect of increasing the spread in the brightness distribution for both samples, however a KS-test returns a 1.5\\,\\% probability in the case of \\chthreecn~and a 1.0\\,\\% probability in the case of \\hthirteencop~that the  radio-quiet and radio loud line-luminosity distributions are drawn from the same population. The average line-luminosity remain higher for the radio-loud class of objects. \\begin{figure} \\begin{center} \\includegraphics[height=7.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_9.epsi} \\caption{~Plot of integrated \\chthreecn~(5\\,--\\,4) flux vs kinematic distance, in kpc. UCH{\\scriptsize II} regions are marked with squares and isolated maser sources are marked with crosses. No correlation is evident.} \\label{fig:xy_ch3cn_dist_uchii} \\end{center} \\end{figure}  \\subsubsection{Interpretation.} Previous work on hot cores suggests \\chthreecn~is the tracer of choice for massive protostars in the hot core phase (e.g. \\citealt{Kurtz2000}), however we find \\chthreecn~is preferentially associated with UCH{\\scriptsize II} regions. \\citet{Pankonin2001} conducted a \\chthreecn\\,(12\\,--\\,11) survey for hot cores and also found a correlation between the presence of \\chthreecn~and UCH{\\scriptsize II} regions. In one of the few high resolution studies \\citet{Remijan2004} used the BIMA array to image the hot core regions W51e1 and W51e2 and found that the \\chthreecn~emission was centred on known UCH{\\scriptsize II} regions.  That we detect \\chthreecn~towards isolated methanol masers, for which there are no other tracers indicative of star-formation, clearly suggests that these objects are internally heated. The lack of \\chthreecn~towards some isolated maser sources is consistent with them being at a less advanced stage of evolution compared to UCH{\\scriptsize II} regions. In young hot cores we would expect any emission to be concentrated in the central regions, where the protostar has heated the dust sufficiently for icy mantles to evaporate and for \\chthreecn~to form. As the temperature rises icy mantles continue to evaporate and the emitting region expands outwards and becomes easier to detect. UCH{\\scriptsize II} regions represent the most advanced evolutionary stage before the young star emerges from its natal cocoon. It is reasonable to assume a relatively extended envelope of \\chthreecn~may exist around the UCH{\\scriptsize II} region, which is less beam diluted and easier for a single-dish survey to detect. Such a scenario is lent weight by high resolutions ($\\sim$\\,1\\arcsec) observations of the hot core in G29.96$-$0.02. \\citet{Olmi2003} derive a kinetic temperature of $\\sim$\\,150\\,K from vibrationally excited \\chthreecn, significantly higher than the $\\sim$\\,90\\,K found by \\citet{Pratap1999}, using ground state lines. \\citet{Olmi2003} also find evidence for a temperature gradient, as the emission from higher energy transitions is confined to increasingly compact regions.  Alternatively, the dominant \\chthreecn~emission may emanate from an unresolved hot core within the same beam. G29.96$-$0.02 is the classic example of a hot core on the edge of a cometary UCH{\\scriptsize II} region (c.f. \\citealt{Cesaroni1998}). We note that in the present survey the \\chthreecn~emission is classified as being associated with the UCH{\\scriptsize II} region, as at 3-mm wavelengths, Mopra does not have the resolution to disentangle the hot core from the nearby ($\\sim$\\,5\\,\\arcsec) radio emission. Conclusions drawn from this work refer to the properties of ``clumps'', which are may contain more than one core. We speculate that clumps containing UCH{\\scriptsize II} regions are more likely to contain an evolved hot core and thus have greater abundances of daughter species, like \\chthreecn.  Further enhancements come from the work of \\citet*{Mackay1999} who developed a chemical model of the hot core G34.3$+$0.15. The model predicts an enhanced abundance of \\chthreecn~in the presence of a far-ultraviolet radiation field. MacKay assumed a spherically symmetric core surrounded by a photon dissociation region (PDR), created by the ultraviolet radiation from a nearby OB association. It is more likely that the cores are inhomogeneous and have clumpy density distributions, in which case the incident radiation will penetrate deeper and the \\chthreecn~abundance will be further enhanced. Seven out of the nineteen UCH{\\scriptsize II} regions in our sample are exactly coincident with methanol maser sites, however the remaining twelve are offset by less than a beam and it is likely that these sources have a complex clumpy structure.   \\subsection{V$_{\\rm LSR}$ and Linewidths} The difference between the \\chthreecn~and \\hthirteencop~V$_{\\rm LSR}$ is plotted in Figure~\\ref{fig:veldiff}. The velocity offsets are all less than 2\\,\\kms, indicating that the molecules are probably within the same star forming region. 92\\,\\% of the sources have velocity offsets less than the fitting errors ($\\pm$\\,0.4\\,\\kms), suggesting that they trace the same clump or core within the beam. \\begin{figure} \\begin{center} \\includegraphics[height=6.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_10.epsi} \\caption{~V${\\rm _{LSR}}$ (\\chthreecn~--~\\hthirteencop). The difference in velocity between the \\chthreecn~and \\hthirteencop~lines is within the errors, making it likely that the emission arises from the same source. Velocity errors are derived from Gaussian fits and are typically $\\pm$\\,0.4\\,\\kms.} \\label{fig:veldiff} \\end{center} \\end{figure}  Figure~\\ref{fig:vlsr_ch3cn_h13cop} displays the distributions of linewidth from the Gaussian fits to the \\chthreecn~and \\hthirteencop~lines. Both distributions are roughly symmetrical about means of 4.9 and 3.5\\,\\kms respectively, higher than typical linewidths of $<$\\,2\\,\\kms~measured in low-mass star-forming regions. At 100\\,K the thermal linewidths of \\chthreecn\\,(5\\,--\\,4) and \\hthirteencop\\,(1\\,--\\,0) are 0.20\\,\\kms~and 0.24\\,\\kms~respectively, with turbulence or bulk gas motions accounting for most of the broadening. The \\chthreecn~linewidth is broader than the \\hthirteencop~linewidth, which suggests that the two species trace different spatial scales. This is consistent with \\chthreecn~being concentrated around a dynamic core and the \\hthirteencop~tracing a more extended quiescent envelope. \\begin{figure} \\begin{center} \\includegraphics[height=6.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_11.epsi} \\caption{~\\chthreecn\\,(5\\,--\\,4) (top panel) and \\hthirteencop\\,(1\\,--\\,0) (bottom panel) linewidth distributions. Both distributions are roughly symmetrical about means of 4.9 and 3.5\\,\\kms~respectively.} \\label{fig:vlsr_ch3cn_h13cop} \\end{center} \\end{figure}  We also note that sources with associated radio continuum emission have larger linewidths than sources without (i.e. isolated maser sources). The mean \\chthreecn\\,(5\\,--\\,4) linewidths are 6.2\\,\\kms~for UCH{\\scriptsize II} regions versus 4.7\\,\\kms~for isolated masers, while the mean \\hthirteencop\\,(1\\,--\\,0) linewidths are 4.0\\,\\kms~for UCH{\\scriptsize II} regions versus 3.4\\,\\kms~for isolated masers. Linewidths are a rough indicator of star-formation activity: we expect to see greater linewidths towards more dynamic regions. The greater linewidths reported for UCH{\\scriptsize II} regions are certainly consistent with a scenario where the isolated maser sources are precursors to the UCH{\\scriptsize II} phase.   \\subsection{\\hcop~Line Profiles} \\begin{figure} \\begin{center} \\includegraphics[width=7.9cm, trim=0 -0 -0 0]{figs/fig_12.epsi} \\caption{~Example \\hcop~line profiles. {\\bf a}) Blue double-peaked profile, {\\bf b}) blue skewed profile, {\\bf c}) red skewed profile, {\\bf d}) red double-peaked profile.} \\label{fig:profile_examples} \\end{center} \\end{figure} Almost all \\hcop~lines exhibit asymmetric line profiles or high-velocity line wings. Complex profiles may be interpreted as either multiple emitting regions along the same line of sight, or as a single emitting region with cold absorbing gas intervening. By examining the line profiles of optically thin \\hthirteencop~we attempt to distinguish between these two causes, and find that most sources are composed of a single broad line with a self absorption dip. Depending on the shape of the profile, we can infer inward or outward motions. Blue-skewed profiles are predicted by collapse models of star formation, however, rotation or outflow-blobs can also produce similar line shapes. The presence of a statistical excess of blue profiles in a survey can indicate that inflow is a likely explanation \\citep*{WuEvans2003}. Two methods of characterising line profiles appear in the literature. Where the opacity is high and the line takes on a double-peaked form, \\citet{WuEvans2003} measure the ratio of the blue to the red peak ${\\rm [T_{MB}(B)/T_{MB}(R)]}$. A blue profile fulfils the criterion: ${\\rm [T_{MB}(B)/T_{MB}(R)]>1}$ by a statistically significant amount. At lower optical depths the absorption will be less severe and the line will appear as a skewed peak with a red or blue shoulder. \\citet{Mardones1997}  suggest an alternative measurement: the line may be classed as blue  if the difference between the peak velocity of the optically-thick  and optically-thin lines is greater than 1/4 the line width of the optically thin line: ${\\rm \\delta v=(v_{thick}-v_{thin})/v_{thin}<-0.25}$. Similarly for a red profile:  ${\\rm \\delta v=(v_{thick}-v_{thin})/v_{thin}>0.25}$. Figure~\\ref{fig:profile_examples} displays a range of \\hcop~profiles from blue to red. We have classed all detected \\hcop~lines using one of the two schemes above. The results are displayed in Figure~\\ref{fig:profile_distrib}. We identify 9 blue and 8 red profiles using the ${\\rm \\delta v}$ measurement, and 12 blue and 10 red profiles using the brightness method. To effect a comparison with other surveys we adopt the concept of ``excess'' introduced by \\citet{Mardones1997} and later used by \\citet{WuEvans2003}: ${\\rm E=(N_{blue}-N_{red})/N_{total}}$, where ${\\rm N_{total}}$ is the number of sources in the sample. The excesses using the ${\\rm \\delta v}$ and brightness methods are E = 0.02 and 0.08, respectively. In comparison \\citet{WuEvans2003} measure excesses of 0.29 and 0.21 in a sample of 28 low-mass star-forming regions. \\begin{figure} \\begin{center} \\includegraphics[width=6.9cm, trim=0 -0 -0 0]{figs/fig_13.epsi} \\caption{Distributions of red and blue profiles measured by ${\\rm (v_{thick}-v_{thin})/v_{thin}}$ (upper panel) and ${\\rm[T_{MB}(B)/T_{MB}(R)]}$ (lower panel). In the top panel the vertical dashed lines mark the absolute velocity difference beyond which profiles are considered to be blue or red skewed. In the lower panel the line divides red and blue profiles. One red and one blue source were discarded as the difference in peak height was within the 1\\,$\\sigma$ noise on the spectrum.} \\label{fig:profile_distrib} \\end{center} \\end{figure}  We also split our sample into radio-strong and radio-weak population, searching for differences in the incidence of blue or red profiles; however we do not find any convincing difference between the populations.   \\subsection{MSX colours} The MSX colours of our sample are primarily determined by the temperature and optical depth of the dust in which they are embedded. As a first attempt to distinguish between classes of source we have derived the colour temperature from the 21/14\\,\\micron~ratio, assuming blackbody emission. Larger ratios are equivalent to lower temperatures and increasingly reddened colours. We have avoided using the 21/8\\,\\micron~or 21/12\\,\\micron~ratios because of possible contamination from PAHs emission at 7.7, 8.6 and 13.3\\,\\micron, and silicate features at 9.7 and 11.3\\,\\micron. Distributions of derived temperatures for sources split into classes with and without UCH{\\scriptsize II} regions and \\chthreecn~emission are displayed in Figure~\\ref{fig:hist_21_14}. A KS-test fails to find any significant difference between the classes. \\begin{figure*} \\begin{center} \\includegraphics[height=6.9cm, trim=0 -0 -0 0]{figs/fig_14.epsi} \\caption{~Colour temperatures derived from the 21\\,/\\,14\\,\\micron~ratio, for sources with and without associated UCH{\\scriptsize II} regions, (A), and \\chthreecn~emission, (B). A KS-test cannot distinguish between the distributions with confidence.} \\label{fig:hist_21_14} \\end{center} \\end{figure*}  Figure \\ref{fig:xy_21_14_trot_uchii} shows a plot of colour temperature against ${\\rm T_{rot}}$, however no correlation is evident, most likely because the thermal infrared and \\chthreecn~emission arise in different regions. Although optical depth effects render colour ratios a poor indicator of kinetic temperature, we note that the rotational temperatures are on average three times lower than the colour temperatures. \\begin{figure} \\begin{center} \\includegraphics[height=7.9cm, angle=-90, trim=0 -0 -0 0]{figs/fig_15.epsi} \\caption{~21\\,/\\,14 \\micron~ratio vs ${\\rm T_{rot}}$. Squares represent sources containing UCH{\\scriptsize II} regions. The 21\\,/\\,14\\,\\micron~ratio does not appear to be correlated with the temperature obtained from the rotational analysis of the \\chthreecn. Note also that the distribution of T$_{\\rm rot}$ is identical for the UCH{\\scriptsize II} regions and isolated masers.} \\label{fig:xy_21_14_trot_uchii} \\end{center} \\end{figure}  \\citet{Lumsden2002} have made a study of sources in the MSX catalogue in an effort to perform a census of massive protostars in the Galactic plane. They have arrived at a set of mid-infrared colour criteria forming the first step in a selection process designed to find the majority of massive protostars present in the Galaxy. As most massive protostars have a featureless red continuum rising between 1 and 100\\,\\micron, they require that ${\\rm F_8<F_{14}<F_{21}}$. Known massive protostars, when plotted on a mid-infrared colour-colour diagram, have an infrared excess such that ${\\rm F_{21}/F_{8}>2}$. In the absence of other selection criteria, sources which satisfy these colour-cuts will include massive protostars, evolved stars, planetary nebula and UCH{\\scriptsize II} regions. Applying these criteria to our data-set we initially discard  28 sources which have incomplete data. Figure~\\ref{fig:colour-colour} shows the mid-infrared colour-colour diagrams for the remaining 55 sources. Associations with UCH{\\scriptsize II} regions and \\chthreecn~are noted by way of different symbols and the selection limits are marked as horizontal and vertical dashed lines. Five sources fail the selection criteria: G23.26$-$0.24 fails both colour-cuts, G30.90$+$0.16 fails because ${\\rm F_{21}/F_{8}<2}$ and the remaining three (G6.54$-$0.11, G12.72$-$0.22, G16.59$-$0.05) fail because ${\\rm F_{14}<F_{8}}$. All sources are radio-quiet and three have detected \\chthreecn~emission, but are unremarkable in other ways. No clustering is seen on the colour-colour plots that would distinguish the classes of source, except perhaps a slight enhancement of the 14/12\\,\\micron~ratio for UCH{\\scriptsize II} regions. UCH{\\scriptsize II} regions are on average also a factor of two brighter in all bands than radio-quiet sources. To further separate young massive stars from planetary nebula and evolved stars \\citet{Lumsden2002} resorted to near-infrared JHK colour selection criteria. 6.7\\,GHz methanol-masers have not, to date, been found towards low-mass star forming regions and we believe that all of our maser sources represent massive stars \\citep{Minier2005}.  Potentially the most interesting sources are `MSX-dark clouds', seen in absorption against the Galactic plane, and the `MSX-red' sources, detected only at 21\\,\\micron. It has been suggested that these sources harbour massive protostars at earlier, more obscured phases of massive star-formation, characterised by low temperatures (T$\\sim$\\,30\\,K) and spectral energy distributions which peak in the sub-millimetre and far-infrared. Table~\\ref{tab:msx_red_dark} summarises the properties of the sources which fall into these categories. The `MSX-dark' category contains thirteen sources, one of which (G30.71$-$0.06) is within 30$''$ of a UCH{\\scriptsize II} region. The `MSX-red' category contains seven sources, one of which (G30.79$+$0.20) is also seen in absorption at 8\\,\\micron~but is detected in emission at 21\\,\\micron.  We find that the `MSX-red' sources are indistinguishable from the bulk of the sample, having approximately the same molecular detection rates and brightness. Four out of eight sources exhibit asymmetric \\hcop~profiles, three of which are skewed to the red, suggestive of outflowing motions. \\begin{figure*} \\begin{center} \\includegraphics[height=6.9cm, trim=0 -0 -0 0]{figs/fig_16.epsi} \\caption{~Mid-infrared colour-colour diagrams for sources detected in all bands. The following symbols mark classes of source: Square = UCH{\\scriptsize II} region with \\chthreecn, Triangle = UCH{\\scriptsize II} without \\chthreecn, Star = Non-UCH{\\scriptsize II} region with \\chthreecn, Cross = Non-UCH{\\scriptsize II} region without \\chthreecn. The dashed lines mark the selection criteria proposed by \\citet{Lumsden2002}: ${\\rm F_8<F_{14}<F_{21}}$ and ${\\rm F_{21}/F_{8}<2}$. Sources which fulfil these criteria include MYSOs, UCH{\\scriptsize II} regions, evolved stars and planetary nebula.} \\label{fig:colour-colour} \\end{center} \\end{figure*}  `MSX-dark' sources exhibit distinct differences from the bulk of the sample. 91\\,\\% were detected in \\chthreecn, and the \\chthreecn~and \\hthirteencop~brightness temperatures tend to be higher. However, by their nature dark-clouds are located nearby making it easier to detect weak lines. Seven of twelve sources exhibit asymmetric \\hcop~profiles, six of which are skewed to the blue, indicative of in-falling motions, and consistent with a scenario of cloud collapse. Further analysis is required to judge if these cores will collapse to form massive stars.   \\subsection{Maserless Sources} Six maserless sources, identified by thermal emission at 1.2-mm from the work of \\citet{Hill2005}, were included in the source list as potential precursors to the hot core phase. One of these, G5.89$-$0.39, was subsequently identified as an UCH{\\scriptsize II} region. The remaining five are spread among the classes of source identified in previous section: G0.26$+$0.01 and G14.99$-$0.70 are MSX-dark clouds, G5.90$-$0.44 and G12.72$-$0.22 are bright and partially resolved MSX sources, and G15.03$-$0.71 is embedded in confused region of extended emission near the M17 molecular cloud. All lie at distances nearer than 4\\,kpc, except for G0.26+0.01 which we place at the distance of the Galactic centre. \\chthreecn~is moderately detected towards all sources, except G15.03-0.71. \\hcop~and \\hthirteencop~are found towards all sources while the \\hthirteencop~linewidths vary from 2\\,--\\,4.5\\,\\kms~without any obvious relation to MSX brightness or class. The absence of a further link between these maserless sources and the spread in observed properties suggests a range of evolutionary phases; i.e. the absence of maser emission towards these sources does not necessarily indicate an early phase star-formation, even with the presence of \\chthreecn."
    },
    "0512/astro-ph0512540_arXiv.txt": {
        "abstract": "We present a model for the accretion flow around the supermassive black hole in the LINER nucleus of \\ngc\\ which fits the optical to X-ray spectral energy distribution (SED). The X-ray segment of the SED is based on observations with the {\\it Chandra X-Ray Observatory}, which are reported here for the first time.  The inner part of the flow is modeled as a radiatively inefficient accretion flow (RIAF) and the outer part as a standard thin disk.  The value of the transition radius ($r_{tr} \\approx 225\\,R_S$, where $R_S=2GM/c^2$) between the RIAF and outer thin disk was obtained from our previous fitting of the double-peaked Balmer emission line profile, which originates in the thin disk. The black hole mass was inferred from measurements of the stellar velocity dispersion in the host galaxy.  When these parameters are used in the accretion flow model, the SED can be successfully reproduced, which shows that the line profile model and the accretion flow model are consistent with each other.  A small remaining excess in the near-UV is accounted by the contribution of an obscured starburst located within 9~pc from the nucleus, as we reported in an earlier paper.  The radio flux is consistent with synchrotron emission of a relativistic jet modeled by means of the internal shock scenario.  In an appendix we also analyze the {\\it Chandra} X-ray observations of the $\\sim$ 1 kpc circumnuclear star-forming ring and of an ultraluminous compact X-ray source located outside the ring. ",
        "introduction": "\\label{sec:intro}  Active galactic nuclei (hereafter AGNs) are thought to be powered by accretion onto a supermassive black hole. The accreting matter is thought to form a thin, equatorial accretion disk, which may reveal itself observationally via its continuum emission (the \\qm{big blue bump}, produced by the sum of black body spectra from annuli in the disk at different temperatures; e.g., \\citealt{frank02,koratkar99}). Another possible signature of a thin accretion disk is the profile of the Fe~K$\\alpha$ X-ray line, which is emitted from the inner parts of the disk and distorted by strong relativistic effects \\citep[see, for example,][]{fabian89,tanaka95}. An alternative signature, broad, double-peaked Balmer emission lines from the outer accretion disk (analogous to the lines of cataclysmic variables) has been observed in the optical spectra of some AGNs \\citep[see review by][]{eracleous98}.  Many AGNs with double-peaked Balmer lines (hereafter ``double-peakers'') were discovered in the spectroscopic survey of radio-loud AGNs \\citep{eh94,eh03} and by the Sloan Digital Sky Survey \\citep{strateva03,wang05}.  However, \\ngc\\ is an important object because it was the first (and the best studied yet) case of a low-luminosity LINER nucleus to display broad, double-peaked H$\\alpha$ and H$\\beta$ emission lines \\citep{sb93}, with a full width at half maximum of $\\approx 7500$~km~s$^{-1}$. This discovery was followed by observations of such lines in other low-luminosity AGNs (LLAGNs), mostly with the \\textit{Hubble Space Telescope} (HST, \\citealt{bower96,shields00,ho00,barth01}). The appearance of the disklike emission lines of \\ngc\\ is thought to be a transient event because they were not seen in previous observations. This transient event was possibly a tidal disruption of a star by a nuclear supermassive black hole, where the debris of the disrupted star feeds the accretion flow, although other possibilities were also discussed by \\citet{sb97}. Since \\ngc\\ is a LLAGN, this transient phenomenon opened up the possibility that we could witness variations in the disklike emission lines on a timescale of years or even months, and the variability of the broad H$\\alpha$ lines was indeed monitored (\\citealt{sb95,sb03}, hereafter SB03). The observed H$\\alpha$ profiles of \\ngc\\ were reproduced by kinematical models of non-axisymmetric thin accretion disks by \\citet{sb97}, who inferred that the line-emitting disk is truncated at an inner radius of $\\approx 225 R_S$, where $R_S=2 G M / c^2$ is the Schwarzschild radius and $M$ is the black hole mass, the black hole mass was measured by \\citet{lewis05} as $(1.2 \\pm 0.2) \\times 10^8 M_\\odot$ from the stellar velocity dispersion using the Ca~II near-IR triplet lines.  On the assumption that the broad double-peaked emission lines observed in AGNs originate in a thin accretion disk, an external source of illumination is needed to power them. Some compelling arguments for this have been put forward by different authors, involving the energy budget of the line-emitting region (e.g., \\citealt{chen89a}) and the low temperatures achieved by the disk under local viscous dissipation \\citep{collin87}.  For instance, considering the energy budget of the thin disk in a small sample of {\\it radio-loud} double-peakers \\citep{eh94,eh03}, the disk can power its own lines only if virtually all the energy available locally is converted into recombination lines, which is unlikely (\\citealt{chen89a}).  The presence of an inner, inflated, radiatively inefficient accretion flow (RIAF, see reviews by \\citealt{quataert01,narayan98r}, hereafter NMQ98) emitting an X-ray continuum which illuminates the outer thin disk solves all these energetic problems, as was proposed initially by \\citet{chen89b} in the form of the ion torus \\citep[an early version of current RIAF models; see][]{rees82}.  While spectral models of RIAFs have successfully accounted for the spectral energy distributions (hereafter SEDs) of a handful of LINERs (\\citealt{lasota96,quataert99,ptak04}, but see \\citealt{yuan02b,herrnstein05} in the case of NGC 4258), RIAFs have yet to be studied in a double-peaker LLAGN. The reason for this is that there was no object until now which had both good observations of the nuclear SED and displayed disklike broad emission lines. We have now such observations for \\ngc, which have allowed us to investigate the structure of the accretion flow in the nucleus by comparing the observed broadband SED with predictions from theoretical models in which the plasma at small radii is in the form of a RIAF and the outer parts of the flow are a standard thin disk. We explore the conditions under which the parameters of the thin disk inferred from the H$\\alpha$ profile model presented by SB03 are consistent with the modelling of the continuum. In \\textsection \\, \\ref{sec:obs} we describe the multiwavelength data comprising the SED, including new X-ray observations as well as previously published data in the optical/ultraviolet, infrared and radio wavebands. In \\textsection \\, \\ref{sec:sed} we derive several properties of the SED, such as the degree of radio-loudness and the bolometric and X-ray luminosities, and we compare the SED of \\ngc\\ with the typical SEDs of quasars and LLAGNs. In \\textsection \\, \\ref{sec:riaf} we discuss the accretion flow model and compare its predictions with the observations. We discuss our results in \\textsection \\, \\ref{sec:discuss} and present our conclusions in \\textsection \\, \\ref{sec:conc}. Throughout this paper we assume a distance\\footnote{We note that \\citet{sb05} assumed a distance to \\ngc\\ of 17 Mpc. We choose the distance 14.5 Mpc in order to compare directly our results for the SED with those of \\citet{ho00}, who adopt the latter distance to \\ngc. Furthermore, this difference in the assumed distance does not change our conclusions and its effect is within flux calibration uncertainties.} to \\ngc\\ of 14.5~Mpc, thus $1\\arcsec = 70$~pc. In an appendix, we present a spectral analysis of the $\\sim$ 1 kpc starforming ring and of an ultraluminous compact X-ray source located outside the ring. ",
        "conclusions": "\\label{sec:conc}  We have reproduced the nuclear optical to X-ray SED of \\ngc, composed of HST and \\emph{Chandra} observations, using a model in which the accretion flow consists of an optically thin, geometrically thick RIAF for radii smaller than the transition radius $r_{\\rm tr} = 225 R_S$ and an optically thick, geometrically thin disk for larger radii. The value of $r_{\\rm tr}$ is the same as the inner disk radius obtained from fitting the double-peaked H$\\alpha$ line profile (SB03) and the value of the black hole mass is that inferred from the stellar velocity dispersion of the host galaxy. Thus, an important result of this work is the consistency between the result of fitting the SED with the accretion flow model and the double-peaked Balmer line with the thin disk model, i.e. we obtain acceptable fits of these two different data sets with the same transition radius.  The accretion flow model underpredicts slightly the UV continuum, but the added contribution of a young and obscured starburst of mass $\\sim 10^6 \\; M_\\odot$ (SB05) can account for the excess UV emission.  The accretion model also underpredicts the radio emission, suggesting the need for a jet component. We have been able to reproduce the radio data with a jet model employing the internal shock scenario (YCN05), although the jet parameters are not as well constrained as those of the accretion flow models, because we have few constraints on the jet emission.  The results of this work support the proposal that the basic components of the nuclei of LLAGNs are a RIAF, a truncated thin disk and a jet (e.g., \\citealt{ho05}), with the possibly important role of starbursts presumably obscured by the circumnuclear torus predicted in the unified model of AGNs."
    },
    "0512/astro-ph0512248_arXiv.txt": {
        "abstract": "The development of fast numerical methods for multilevel radiative transfer (RT) applications often leads to important breakthroughs in astrophysics, because they allow the investigation of problems that could not be properly tackled using the methods previously available. Probably, the most familiar example is the so-called Multilevel Accelerated $\\Lambda$-Iteration (MALI) technique of Rybicki \\& Hummer for the case of a local approximate operator, which is based on Jacobi iteration. However, there are superior operator-splitting methods, based on Gauss-Seidel (GS) and Successive Overrelaxation (SOR) iteration, which provide a dramatic increase in the speed with which non-LTE multilevel transfer problems can be solved in one, two and three-dimensional geometries. Such RT methods, which were introduced  by Trujillo Bueno \\& Fabiani Bendicho ten years ago, are the main subject of the first part of this paper. We show in some detail how they can be applied for solving multilevel RT problems in spherical geometry, for both atomic and molecular line transitions. The second part of the article addresses the issue of the calculation of the molecular number densities when the approximation of instantaneous chemical equilibrium turns out to be inadequate, which happens to be the case whenever the dynamical time scales of the astrophysical plasma under consideration are much shorter than the time needed by the molecules to form. ",
        "introduction": "Astrophysical plasma spectroscopy is a fascinating but very complex research field because non-equilibrium physics is required. First, for given atomic or molecular number densities at each point within the astrophysical plasma under consideration, we have the problem of how to calculate the population numbers of the levels included in the atomic or molecular model. Second, it is also important to remember that the atomic and/or molecular number densities themselves are not always given by the instantaneous values of the physical conditions of the plasma (e.g., kinetic temperature, density), simply because such conditions may be changing rapidly with time as as result of magneto-hydrodynamical processes. Therefore, for instance, it may be ``dangerous'' to infer chemical abundances in stellar atmospheres by computing the molecular number densities via the approximation of instantaneous chemical equilibrium.  The present keynote article addresses these two aspects of astrophysical plasma spectroscopy, with emphasis on the increasingly attractive topic of non-LTE radiative transfer in molecular lines (e.g., note that in the very near future the Herschell Space Observatory of ESA is expected to be fully operative). ",
        "conclusions": "Thirty years ago, the total computational work ($W$) of the best available multilevel transfer codes scaled approximately as NP$^3$, with NP the number of spatial grid points in a computational domain of fixed dimensions (e.g., Mihalas 1978). Twenty years ago, operator splitting methods based on Jacobi iteration were developed that yielded $W{\\sim}{\\rm NP}^2$ (e.g., Olson, Auer \\& Buchler 1986; Rybicki \\& Hummer 1991, 1992; Auer, Fabiani Bendicho \\& Trujillo Bueno 1994). Ten years ago, radiative transfer methods based on Gauss-Seidel and SOR iteration were developed for which $W{\\sim}{\\rm NP}{\\,}\\sqrt{{\\rm NP}}$, which imply an-order-of-magnitude of improvement with respect to the MALI method (Trujillo Bueno \\& Fabiani Bendicho 1995; Trujillo Bueno \\& Manso Sainz 1999; Trujillo Bueno 2003). Such methods are also suitable for 2D and 3D radiative transfer applications because they do not require neither the construction nor the inversion of any non-local approximate $\\Lambda$-operator (see the MUGA-3D code of Fabiani Bendicho \\& Trujillo Bueno 1999). Moreover, MUGA is the iterative method of choice for the development of the only multilevel RT method for which $W{\\sim}{\\rm NP}$ -that is, the non-linear multigrid method described by Fabiani Bendicho, Trujillo Bueno \\& Auer (1997).  In this article we have shown how the Multilevel Gauss-Seidel (MUGA) and Multilevel Successive Overrelaxation (MUSOR) methods can be applied for solving multilevel radiative transfer problems in spherical geometry, for both atomic and molecular lines. Our multilevel codes based on the MUGA and MUSOR methods offer a very powerful tool for the fast solution of realistic RT problems."
    },
    "0512/astro-ph0512281_arXiv.txt": {
        "abstract": "We investigate the impact finite simulation box size has on the structural and kinematic properties of Cold Dark Matter haloes forming in cosmological simulations. Our approach involves generating a single realisation of the initial power spectrum of density perturbations and studying how truncation of this power spectrum on scales larger than $L_{\\rm cut} = 2\\pi/k_{\\rm cut}$ affects the structure of dark matter haloes at $z=0$. In particular, we have examined the cases of $L_{\\rm cut} = f_{\\rm cut} L_{\\rm box}$ with $f_{\\rm cut}=1$ (i.e. no truncation)$, 1/2, 1/3$ and $1/4$. In common with previous studies, we find that the suppression of long wavelength perturbations reduces the strength of clustering, as measured by a suppression of the 2-point correlation function $\\xi(r)$, and reduces the numbers of the most massive haloes, as reflected in the depletion of the high mass end of the mass function $n(M)$. Interestingly, we find that truncation has little impact on the internal properties of haloes. The masses of high mass haloes decrease in a systematic manner as $L_{\\rm cut}$ is reduced, but the distribution of concentrations is unaffected. On the other hand, the median spin parameter is $\\sim 50\\%$ lower in runs with $f_{\\rm cut}<1$. We argue that this is an imprint of the linear growth phase of the halo's angular momentum by tidal torquing, and that the absence of any measurable trend in concentration and the weak trend observed in halo shape reflect the importance of virialisation and complex mass accretion histories for these quantities. These results are of interest for studies that require high mass resolution and statistical samples of simulated haloes, such as simulations of the population of first stars. Our analysis shows that large-scale tidal fields have relatively little effect on the internal properties of Cold Dark Matter haloes and hence may be ignored in such studies. ",
        "introduction": "\\label{sec:intro}  Over the last twenty years, the cosmological N-Body simulation has come to be firmly established as the preeminent tool for theoretical studies of the formation and evolution of structure in the Universe (see \\citet{bert98} for a relatively recent review). Whereas early studies explored the gross features of different cosmological models and led to the emergence of the Cold Dark Matter (CDM) paradigm \\citep{frenk83,white83,davis85}, the most recent generation of simulations are being used to provide detailed predictions of the CDM model and in particular its $\\Lambda$CDM variant. These include the behaviour of CDM halo mass profiles on scales that can be directly compared with observations \\citep[e.g.][]{hayashi04,navarro04,diemand04} and sophisticated semi-analytical models of galaxy formation built upon merger trees harvested from high resolution simulations \\citep[c.f. the Millennium Simulation described in][]{springel05}. The N-Body approach is not without its limitations, however, and given its significance, it is of fundamental importance to understand what these limitations are and how they affect the interpretation of the results of simulations.  Such considerations generally lead to ``convergence'' studies that seek to quantify the degree to which the results of a simulation are affected by choices made in its setting up and running. For example, both \\citet{nfw96} and \\citet{moore98} demonstrated that their results on the form of CDM halo mass profiles were not compromised by force resolution (or ``softening'') as well as starting redshift \\citep[][]{nfw96} and particle number \\citep[][]{moore98}. The recent study of \\citet{power03} established a set of convergence criteria \\citep[reaffirmed by][]{hayashi04} that determine the radial range over which the circular velocity profile of a simulated CDM halo is reliable to better than $10\\%$ accuracy, for a number of simulation parameters including softening, time integration accuracy, and particle number \\citep[see also][]{diemand04,binney04}. More generally, \\citet{alex} investigated how properties such as the correlation function, mass function and central densities of dark matter haloes varied with softening and time integration accuracy at fixed particle mass in simulations run with different N-Body codes.\\\\  Recently, \\citet[][hereafter BR05]{bagla05} have demonstrated that size of simulation box $L_{\\rm box}$ can also affect certain properties of the dark matter halo population, in particular the 2-point correlation function $\\xi(r)$ and the mass function $n(>M)$. $L_{\\rm box}$ sets the longest wavelength perturbation that can be resolved in a simulation and it follows that reducing $L_{\\rm box}$ suppresses the contribution of large scale perturbations to the power spectrum. The effect of this suppression was highlighted by \\citet[][hereafter GB94]{gelb94b}, who examined the importance of long wavelength perturbations in the initial conditions for $\\xi(r)$ and the correlation length $r_0$ in the context of the SCDM model. They investigated how $\\sigma_8$ (the linear mass variance in a sphere of radius $8 h^{-1}\\,{\\rm Mpc}$) and consequently $\\xi(r)$ varied with $L_{\\rm box}$ (see their figures 1 and 2), and found that $L_{\\rm box} \\gtrsim 100 {\\rm Mpc}$ (or $50 h^{-1} {\\rm Mpc}$ for $h=0.5$) was required to correctly recover both $\\sigma_8$ and the amplitude of $\\xi(r)$. In other words, studies that wish to accurately characterise the clustering properties of dark matter haloes require large simulation boxes.  How large the simulation box must be to accurately recover $\\xi(r)$ was the question addressed by BR05. They noted that finite box size can affect not only $\\xi(r)$ but also the high mass end of the mass function $n(>M)$; as $L_{\\rm box}$ is reduced below some threshold value, the numbers of the most massive haloes decrease in a systematic fashion. The mass function $n(>M)$ can be characterised by the \\citet{ps74} form which is a relatively simple function of the linear mass variance $\\sigma(M)$. For a given mass $M = {4\\pi/3} \\bar{\\rho} R^3$ with $\\bar{\\rho}$ mean density, \\begin{equation} \\label{eq:sigma} {\\sigma^2(M) = \\int^\\infty_{2\\pi/{L_{\\rm box}}} P(k) W^2(kR) d^3k,} \\end{equation} where $P(k)$ is the linear power spectrum, $W^2(kR)$ is the top-hat filter, and the lower limit of the integral $2\\pi/L_{\\rm box}$ corresponds to the fundamental mode in the simulation box. BR05 investigated how varying $L_{\\rm box}$ impacted on the form of $n(>M)$ and used this as a simple criterion for determining how large $L_{\\rm box}$ must be to recover $n(>M)$ (and consequently $\\xi(r)$) to a given accuracy.\\\\  The studies of GB94 and BR05 clearly indicate that large simulation boxes are required if we wish to recover accurate mass and two point correlation functions of the dark matter halo population, but does the choice of $L_{\\rm box}$ also affect the internal properties of the haloes? This question is of interest for studies that require the high spatial resolution offered by simulations of small volumes but that may not require accurate clustering information, such as studies of the first objects at high redshift \\citep[e.g.][]{abel02}, and it is one that we shall address in the present study.  Why might we expect finite $L_{\\rm box}$ to be of importance for the internal structure of haloes? The suppression of long wavelength perturbations will modify the ``global'' tidal field that a dark matter halo is subject to, which may have an impact on the halo's shape, its angular momentum content \\citep[e.g.][]{white84} and the infall pattern of substructures \\citep[e.g.][]{aubert04,benson05}. Similarly, the linear mass variance controls the formation time of dark matter haloes \\citep[][]{nfw96,bullock01,ens01} and so we might expect haloes to form at systematically later times in smaller boxes with lower ``effective'' values of $\\sigma_8$; this may then affect both the mass of the halo and its central density (or concentration) measured at $z=0$. We also note that small- and large-scale modes are coupled during the non-linear clustering phase and there is a transfer of power from large scales down to small scales \\citep[\\eg][]{bagla97}; neglecting this power may leave an imprint on the formation and evolution of gravitationally bound objects. On the other hand, dark matter haloes are, by definition, virialised systems and much of the information that was present during the initial stages of their collapse will be erased during virialisation \\citep[c.f. the universal mass profile of][]{nfw96,nfw97}, so it is not clear whether a finite $L_{\\rm box}$ will have an effect at all. Nevertheless, it is important to investigate this question and determine how severe a limitation it might be.\\\\  The outline of this paper is as follows; in the next section, we briefly describe the simulations we have used in this study before presenting our results (\\S~\\ref{sec:results}). We confirm the findings of GB94 and BR05 (\\S~\\ref{ssec:bulk}) before examining how various characteristic properties of the halo population -- concentrations, shapes and spins -- are affected by the suppression of long wavelength perturbations on scales larger than $L_{\\rm cut} = f_{\\rm cut}\\,L_{\\rm box}$. Finally, we discuss our results in \\S~\\ref{sec:discussion} and offering some concluding remarks in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The aim of this study has been to establish whether finite box size has a measurable effect on the internal properties of simulated Cold Dark Matter haloes. The results of our analysis suggest that the effects, if present, are small and not statistically significant in most cases. Of the principal quantities we have examined -- concentration, shape and the spin parameter -- we find that spin shows the most prominent effect; the median spin parameter is $\\sim 50\\%$ smaller in our truncated runs, independent of the mass of the system. We argue that this is an imprint of the linear growth phase of the halo's angular momentum by tidal torquing, and that the absence of any measurable trend in concentration or strong trend in shape reflect the importance of virialisation and complex mass accretion histories for these quantities respectively.  These results are useful because they clarify what properties of simulated dark matter haloes are affected by finite box size, and the severity of these effects. Indeed, they are of some importance because they reveal that studies of the internal properties of statistical samples of haloes that do not require clustering information that are based on high resolution simulations of small volumes are not compromised. The abundance of high mass systems will be underestimated, but the properties of low- to intermediate-mass systems are reliable. This is of interest to studies of, for example, the formation of the first generation of population III stars at high redshift \\citep[e.g.][]{abel02}; in these cases, clustering information is not as important (although it would be for reionisation studies, say).  In conclusion, our study has demonstrated that the effects of finite box size are negligible for the internal properties of dark matter haloes."
    },
    "0512/astro-ph0512138_arXiv.txt": {
        "abstract": "We estimate the star-formation rates and the stellar masses of the Extremely Red Objects (EROs) detected in a  $\\approx180 \\rm \\, arcmin^2$ $Ks$-band survey ($Ks\\approx20$\\,mag). This sample is complemented by sensitive 1.4\\,GHz radio observations (12$\\mu$Jy $1\\sigma$ rms) and multiwaveband photometric data ($UBVRIJ$) as part of the Phoenix Deep Survey. For bright $K<19.5$\\,mag  EROs in this sample ($I-K>4$\\,mag; total of 177) we use photometric methods to discriminate dust-enshrouded active systems from  early-type galaxies and to constrain their redshifts. Radio stacking is then employed to estimate mean radio flux densities of $\\approx 8.6$ ($3\\sigma$) and $6.4\\,\\mu$Jy ($2.4\\sigma$) for the dusty and early-type subsamples respectively. Assuming that dust enshrouded active EROs are powered by star-formation the above radio flux density at the median redshift of $z=1$ translates to a radio luminosity of $L_{1.4} = 4.5 \\times  10^{22} \\rm W/Hz$  and a star-formation rate of  $\\rm SFR = 25 \\, M_{\\odot} \\, yr^{-1}$. Combining this result with photometric redshift estimates we find a lower limit to the star-formation rate density of $0.02\\pm0.01 \\rm \\, M_{\\odot}  \\, yr^{-1} \\, Mpc^{-3}$ for the $K<19.5$\\,mag dusty EROs in the range $z=0.85-1.35$. Comparison with the star-formation rate density estimated for previous ERO samples (with similar selection criteria) using optical emission lines, suffering dust attenuation, suggests a mean dust reddening of at least $E(B-V)\\approx0.5$ for this population.  We further use the $Ks$-band luminosity as proxy to stellar mass and argue that the dust enshrouded starburst EROs in our sample are massive systems, $\\rm M \\ga 5 \\times 10^{10} \\, M_{\\odot}$. We also find that EROs represent a sizable fraction (about 50 per cent) of the number density of galaxies more massive than $\\rm M = 5 \\times 10^{10} \\, M_{\\odot}$ at $z\\approx1$, with almost equal contributions from dusty and early type systems. Similarly, we find that EROs contribute about half of the mass density of the Universe at $z\\approx1$ (with almost equal contributions from dusty and early types), after taking into account incompleteness because of the magnitude limit $K=19.5$\\,mag. ",
        "introduction": "The class of Extremely Red Objects (EROs; $R-K>5$, $I-K>4$\\,mag), first identified more than 15 years ago (Elston  et al. 1988), is believed to comprise a heterogeneous population of $z \\ga 1$ systems split between passive galaxies and dust enshrouded AGNs/starbursts (Cimatti et al. 2003). The identification of either type of galaxies (early or dusty) at high-$z$ has important  cosmological implications, thus providing significant impetus in  ERO studies.  For example, finding early-type massive systems at high-$z$ can provide information on both the galaxy formation redshift (Spinrad et al. 1997; Cimatti et al. 2002) and the global mass assembly (e.g. Fontana et al. 2003, 2004; Drory 2004; Glazebrook 2004), thus constraining galaxy formation scenarios: Monolithic collapse early in the Universe ($z_f>2-3$) followed by passive evolution (e.g. Eggen et al. 1962; Larson 1975) versus hierarchical merging and relatively recent formation epochs (Baugh et al. 1996; Kauffmann 1996). Similarly a population of $z \\ga 1$ dusty active galaxies, AGNs or starbursts, that are missing from UV/optical surveys may play an important role  in the global star-formation history (e.g. Haarsma et al. 2000; Smail et al. 2002; Hopkins 2004), the evolution of AGNs and the interpretation of the diffuse X-ray Background with respect to the (still) elusive population of high-$z$ obscured QSOs (e.g. Hasinger et al. 2003).  Recent developments in instrumentation have yielded large ERO samples allowing  systematic study of their statistical properties in a cosmological context. Cimatti et al. (2002) used optical spectroscopy from the {\\sf K20} survey to explore the star-formation rates (SFR) of $K<19.2$\\,mag EROs and to assess their contribution to the  global star-formation history. Under conservative assumptions about reddening, these authors find that  dust-enshrouded EROs represent a small but non-negligible fraction (about 10 per cent) of the SFR density at $z\\approx1$. Dust obscuration issues, however, make this result uncertain. Smail et al. (2002) expanded on the Cimatti et al. (2002) study using deep radio imaging to explore the SFR of $K<20.5$\\,mag EROs independent of dust induced biases. Using this deeper $K$-band sample they estimate star-formation densities higher than those of Cimatti et al. (2002) and suggest that obscured galaxies make a sizable contribution to the total SFR density at $z\\approx1$. More recently Caputi et al. (2005) used the GOODS-South data to investigate the evolution of $K$-band selected galaxies. They argue that EROs among their sample constitute a sizable fraction (about 50-70 per cent) of galaxies with stellar mass $\\rm M>5\\times 10^{10}\\,M_{\\odot}$ at $z=1-2$, suggesting that they represent a major component of the stellar mass build-up at these redshifts.  The observational developments above are also complemented with efforts to model EROs using either semi-analytical (e.g. Somerville et al. 2004a) or hydrodynamical (e.g. Nagamine et al. 2005) methods. Despite significant progress, accounting for both the red colours and the number density of EROs remains a challenge for these numerical simulations. Part of the difficulty lies in our poor understanding of some of the properties of EROs. Open questions include what are the dust properties of these systems, what is the number density of dusty and early type EROs, what is the relative contribution of these sub-populations to the mass density.  In this paper we add to the discussion on the cosmological significance of EROs by combing an $\\approx 180 \\, \\rm arcmin^{2}$ deep ($Ks\\approx20$\\,mag) $Ks$-band survey with ultra-deep 1.4\\,GHz radio data ($\\approx \\rm 60 \\, \\mu Jy$) carried out as part of the Phoenix Deep Survey (Hopkins et al. 2003). This sample has already been used to explore the clustering properties and the environment of EROs (Georgakakis et al. 2005). The advantage of our survey is depth combined with wide area coverage reducing cosmic variance issues. Additionally, the ultra-deep radio data allow SFR estimates for EROs free from the dust obscuration effects that are expected to be important in this class of sources. Throughout the paper we adopt $\\rm H_{o}=70\\,km\\,s^{-1}\\,Mpc^{-1}$, $\\rm \\Omega_{M}=0.3$ and  $\\rm \\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "In this paper we use an $\\approx180 \\rm \\, arcmin^2$ $Ks$-band survey overlapping with ultra-deep radio observations and  multiwaveband photometry ($UBVRIJ$) to estimate the SFRs and the stellar masses of EROs with $K<19.5$\\,mag. Template SEDs are fit to the broad-band optical/NIR photometric data to discriminate between dusty and early-type EROs and to  determine their photometric redshifts. This classification is found to be in fair agreement with methods using optical/NIR colours (e.g. $I-K$ vs $J-K$).  About 8 per cent of the $K \\la 19.5$\\,mag EROs have radio counterparts to the flux density limit of about $60\\rm \\, \\mu Jy$. For the remaining sources we use radio stacking analysis to constrain their mean radio properties. We estimate a stacked signal of about $9$ ($3\\sigma$) and $6 \\, \\rm \\mu Jy$ ($2.5\\sigma$) for dusty and early-type EROs respectively. Assuming that the radio emission in the dusty sub-population is due to star-formation activity we estimate a mean star-formation rate of $\\rm SFR  =  25\\,M_{\\odot} \\, yr^{-1}$ at $z=1$ free from dust obscuration effects. The radio detected EROs at $z\\approx1$ have, on average, much higher SFRs in the range $\\rm 100-1000 \\, M_{\\odot} \\, yr^{-1}$. Combining these results with the photometric redshift estimates we find that the SFR density of the Universe in the range $0.85-1.35$ due to $K < 19.5$\\,mag EROs is $\\rm 0.02\\pm0.01 \\, M_{\\odot} \\, yr^{-1} \\, Mpc^{-3}$. This should be considered a lower limit since EROs fainter than $K = 19.5$\\,mag are not taken into account. Comparison with deeper samples suggests that correcting for these biases will likely increase the estimate above by at least a factor of 2. We also argue that the systems responsible for the observed star-formation density are dusty starbursts more massive than  $\\rm M = 5 \\times  10^{10} \\, M_{\\odot}$. Less massive dusty EROs lie below the magnitude limit $K=19.5$\\,mag at $z\\approx1$ and are the systems responsible for the missing (at least factor of 2) SFR density at $z\\approx1$. Comparison of the dust-independent SFR density estimated here with that of similarly selected ERO samples using optical emission lines, suffering dust attenuation, suggests a mean dust reddening of at least $E(B-V)\\approx0.5$ for this population.  We further use mass-to-light ratios of the best-fit template SED to convert the $Ks$-band luminosity of EROs to stellar mass. We find that EROs contribute about 50 per cent to the total number density of galaxies with stellar mass $\\rm M >  5 \\times  10^{10} \\,  M_{\\odot}$ at $z\\approx1$. This fraction is almost equally split between dusty and early type systems. We further estimate that the $K<19.5$ EROs represent about 50 per cent of the global mass density in the redshift range $z=0.85-1.35$, after taking into account incompleteness due to the magnitude limit $K=19.5$\\,mag. This indicates  that these systems are also a non-negligible component of the Universe mass build-up at these redshifts. The ERO mass density above is also almost equally split between the dusty and early type subpopulations."
    },
    "0512/astro-ph0512412_arXiv.txt": {
        "abstract": "Photoelectric photometry on the extended Str\\\"omgren system ($uvbyCa$) is presented for 7 giants and 21 main sequence stars in the old open cluster, NGC 752. Analysis of the $hk$ data for the turnoff stars yields a new determination of the cluster mean metallicity. From 10 single-star members, [Fe/H] = $-0.06\\pm0.03$, where the error quoted is the standard error of the mean and the Hyades abundance is set at [Fe/H] = +0.12. This result is unchanged if all 20 stars within the limits of the $hk$ metallicity calibration are included. The derived [Fe/H] is in excellent agreement with past estimates using properly-zeroed $m_1$ data, transformed moderate-dispersion spectroscopy, and recent high dispersion spectroscopy. ",
        "introduction": "Our increasingly sophisticated understanding of the Galaxy's chemical history has been advanced by observational techniques of varying sensitivity and reach.  High-dispersion spectroscopic studies provide the most detailed picture of elemental abundances but are often restricted to sample sizes of dozens (e.g., \\citet{spit,luck}) or, occasionally, hundreds of stars \\citep{edv93,val5}. Broad-band techniques can simultaneously reach thousands of stars but can be subject to lower parametric sensitivity and more severe coupling of the variables of interest. Somewhere between the exquisite spectral resolution of high-dispersion studies and the more robust broad-band photometric survey tools, the contributions of intermediate-band and narrow-band photometry can provide surprisingly precise estimates of chemical composition and foreground reddening for individual stars, leading to improved determinations of distance and age for individual stars and stellar aggregates in the Milky Way.  Since 1991, we have augmented the highly successful standard Str\\\"omgren $uvby$H$\\beta$ bandpasses with a relatively narrow filter, $Ca$, centered on the near-ultraviolet ionized lines of Calcium \\citep{att91}.  Constructed as a color difference analogous to the Str\\\"omgren $m_1$ index with the $v$ filter replaced by $Ca$, $hk$ retains sensitivity to metallicity for very low metal abundance stars, the original motivation for adding the filter (e.g., \\citet{atmp}), to metallicities as high as solar- and Hyades-abundance F dwarfs \\citep{at02}.  The extended Str\\\"omgren system was established by \\citet{att91} with photoelectric observations of $\\sim$150 stars, and later augmented by nearly 2000 stars with photoelectric observations \\citep{att95}, predominantly in the southern hemisphere. Despite the relatively large size of this photoelectric sample, the availability of $hk$ standards is limited to fairly bright field stars scattered around the sky -- not an optimally efficient situation for modern CCD observations.  In large part, the present work was motivated by the desirability of developing $hk$ photometry in a northern open cluster for future use as a secondary standard field for CCD observations.  Although we did not secure as many observations per star as would be desired for secondary standards, the photoelectric photometry discussed in this paper has already been used to supplement field star standards in the reduction of CCD data obtained at the WIYN 0.9 meter telescope in 2003 and 2004; the resulting analysis of $Cavby$H$\\beta$ photometry in NGC 2420 is described in a separate publication \\citep{att06}.  NGC 752 presents the mixed blessing of being relatively near: it is dominated by moderately bright stars easily within reach of moderate size telescopes while being relatively low in concentration class. While it extends over nearly a degree in diameter (larger than many CCD fields), its bright uncrowded field lends itself to relatively simple, aperture-based measurement strategies.  A review of the cluster's accepted fundamental parameters, as well as a new derivation of the cluster metallicity from the $hk$ index and an updated discussion of the cluster age, (Sec. 4) will follow a brief discussion of the data acquisition and standardization (Sec. 2), and a comparison with previous photometric studies (Sec. 3). Section 5 summarizes our results. ",
        "conclusions": "Photoelectric photometry on the extended Str\\\"omgren system has been obtained of stars covering a wide range in temperature and luminosity class in the old open cluster NGC 752. While the primary motivation for the observations is the establishment of potential standard stars for CCD applications in the northern sky, the new data also supply an opportunity for evaluating the cluster metallicity, a subject of interest for the last 20 years.  Using metallicity calibrations for the $hk$ index derived and refined over the last decade, the mean [Fe/H] from 10 stars, selected as single stars without evidence for spectroscopic anomalies, becomes --0.06$\\pm$0.03 (s.e.m.), on a scale where the Hyades has [Fe/H] = +0.12. If all 20 stars within the calibration range of the index are used, the mean and the dispersion remain the same. This is slightly higher than the value ([Fe/H] = $-0.15\\pm$0.05) found in the comprehensive study of NGC 752 by \\citet{dan}, and the same within the errors as the abundance derived in \\citet{taat} from the weighted average of DDO photometry and transformed moderate-dispersion spectroscopy ([Fe/H] = $-0.09 \\pm 0.02$). These values are also in excellent agreement with the previous $uvby$ analyses of over two dozen stars at the cluster turnoff by both \\citet{cb70} and \\citet{ni88}, whose photometry generates [Fe/H]$ = -0.08 \\pm 0.02$ and [Fe/H]$= -0.05 \\pm 0.03$, respectively. The only large, high-dispersion spectroscopic survey of the cluster is that of \\citet{sest}, who finds [Fe/H]$ = -0.01 \\pm 0.04$, on a Hyades scale of +0.12. This abundance is clearly consistent within the uncertainties with the estimates obtained by earlier $uvby$ studies and recent cluster compilations when transformed to a common scale. Thus, the best estimate for the metallicity of NGC 752 that encompasses the photometric and spectroscopic results is [Fe/H]$ = -0.05 \\pm 0.05$, on a scale where the Hyades has [Fe/H] = +0.12."
    },
    "0512/astro-ph0512142_arXiv.txt": {
        "abstract": "We determine the carbon isotopic ratios in the atmospheres of some evolved stars in both globular clusters and the disk of our Galaxy. Analysis of $^{12}$CO and $^{13}$CO bands at 2.3 micron was carried out using fits to observed spectra of red giants and Sakurai's object (V4334 Sgr). The dependence of theoretical spectra on the various input parameters was studied in detail. The computation of model atmospheres and a detailed abundance analysis was performed in a self-consistent fashion. A special procedure for determining the best fits to observed spectra was used. We show, that globular cluster giants with [Fe/H] $<$ -1.3 have a low $^{12}$C/$^{13}$C = 4 $\\pm$ 1 abundance ration. In the spectra of Sakurai's object (V4334 Sgr) taken between 1997-98, the 2.3 micron spectral region is veiled by hot dust emission. By fitting UKIRT spectra we determined $^{12}$C/$^{13}$C = 4 $\\pm$ 1 for the July, 1998 spectrum. CO bands in the spectra of ultracool dwarfs are modelled as well. ",
        "introduction": "The carbon isotopes $^{12}C$ and  $^{13}C$ are produced inside low and intermediate mass stars mainly. The carbon isotope ratio $^{12}C$/$^{13}C$ is often used as a tracer of the nucleosynthetic and chemical mixing processes inside stars. Fast evolving chemically peculiar stars provide a unique tool for the verification of our knowledge about the process of stellar evolution. In contrast, some low mass dwarfs do not fuse $^1$H, these are known as brown dwarfs. These objects preserve their initial $^{12}C$ and $^{13}C$, and can be used for testing the Galaxy evolution theories. ",
        "conclusions": ""
    },
    "0512/astro-ph0512374_arXiv.txt": {
        "abstract": "High-$z$ galaxy redshift surveys open up exciting possibilities for precision determinations of neutrino masses and inflationary models. The high-$z$ surveys are more useful for cosmology than low-$z$ ones owing to much weaker non-linearities in matter clustering, redshift-space distortion and galaxy bias, which allows us to use the galaxy power spectrum down to the smaller spatial scales that are inaccessible by low-$z$ surveys. We can then utilize the two-dimensional information of the linear power spectrum in angular and redshift space to measure the scale-dependent suppression of matter clustering due to neutrino free-streaming as well as the shape of the primordial power spectrum. To illustrate capabilities of high-$z$ surveys for constraining neutrino masses and the primordial power spectrum, we compare three future redshift surveys covering 300~square degrees at $0.5<z<2$, $2<z<4$, and $3.5<z<6.5$. We find that, combined with the cosmic microwave background data expected from the Planck satellite, these surveys allow precision determination of the total neutrino mass with the projected errors of $\\sigma(m_{\\nu, {\\rm tot}})=0.059$, 0.043, and 0.025~eV, respectively, thus yielding a positive {\\it detection} of the neutrino mass rather than an upper limit, as $\\sigma(m_{\\nu, {\\rm tot}})$ is smaller than the lower limits to the neutrino masses implied from the neutrino oscillation experiments, by up to a factor of 4 for the highest redshift survey. The accuracies of constraining the tilt and running index of the primordial power spectrum, $\\sigma(n_s)=(3.8,~3.7,~3.0)\\times 10^{-3}$ and  $\\sigma(\\alpha_s)=(5.9,~5.7,~2.4)\\times 10^{-3}$ at $k_0=0.05~{\\rm Mpc}^{-1}$, respectively, are smaller than the current uncertainties by more than an order of magnitude, which will allow us to discriminate between candidate inflationary models. In particular, the error on  $\\alpha_s$ from the future highest redshift survey is not very far away from the prediction of a class of simple inflationary models driven by a massive scalar field with self-coupling, $\\alpha_s=-(0.8-1.2)\\times 10^{-3}$. ",
        "introduction": "We are living in the golden age of cosmology. Various data sets from precision measurements of temperature and polarization anisotropy in the cosmic microwave background (CMB) radiation as well as those of matter density fluctuations in the large-scale structure of the universe mapped by galaxy redshift surveys, Lyman-$\\alpha$ forests and weak gravitational lensing observations are in a spectacular agreement with the concordance $\\Lambda$CDM model \\cite{WMAP,Tegmark,Seljak,jarvis/jain:2005}. These results assure that theory of cosmological linear perturbations is basically correct, and can accurately describe the evolution of photons, neutrinos, baryons, and collisionless dark matter particles \\cite{Peebles,SZ,BE}, for given initial perturbations generated during inflation \\cite{liddle/lyth:2000,dodelson}. The predictions from linear perturbation theory can be compared with the precision cosmological measurements, in order to derive stringent constraints on the various basic cosmological parameters. Future observations with better sensitivity and higher precision will continue to further improve our understanding of the universe.  Fluctuations in different cosmic fluids (dark matter, photons, baryons, and neutrinos) imprint characteristic features in their power spectra, owing to their interaction properties, thermal history, equation of state, and speed of sound. A remarkable example is the acoustic oscillation in the photon-baryon fluid that was generated before the decoupling epoch of photons, $z\\simeq 1088$, which has been observed in the power spectrum of CMB temperature anisotropy \\cite{WMAPcl}, temperature--polarization cross correlation \\cite{WMAPte},  and distribution of galaxies \\cite{bao1,bao2}.  Yet, the latest observations have shown convincingly that we still do not understand much of the universe. The standard model of cosmology tells us that the universe has been dominated by four components. In chronological order the four components are: early dark energy (also known as ``inflaton'' fields), radiation, dark matter, and late-time dark energy. The striking fact is that we do not understand the precise nature of three (dark matter, and early and late-time dark energy) out of the four components; thus, understanding the nature of these three dark components has been and will continue to be one of the most important topics in cosmology in next decades. Of which, one might be hopeful that the next generation particle accelerators such as the Large Hadron Collider (coming on-line in 2007) would find some hints for the nature of dark matter particles. On the other hand, the nature of late-time dark energy, which was discovered by measurements of luminosity distance out to distant Type Ia supernovae \\cite{sn1,sn2}, is a complete mystery, and many people have been trying to find a way to constrain properties of dark energy (see, e.g., \\cite{DEreview} for a review).  How about the early dark energy, inflaton fields, which caused the expansion of the universe to accelerate in the very early universe? We know little about the nature of inflaton, just like we know little about the nature of late-time dark energy. The required property of inflaton fields is basically the same as that of the late-time dark energy component: both must have a large negative pressure which is less than $-1/3$ of their energy density. To proceed further, however, one needs more information from observations. Different inflation models make specific predictions for the shape of the power spectrum \\cite{liddle/lyth:2000} (see also Appendix~B) as well as for other statistical properties \\cite{NGreview} of primordial perturbations. Therefore, one of the most promising ways to constrain the physics of inflation, hence the nature of early dark energy in the universe, is to determine the shape of the primordial power spectrum accurately from observations. For example, the CMB data from the Wilkinson Microwave Anisotropy Probe \\cite{WMAP}, combined with the large-scale structure data from the Two-Degree Field Galaxy Redshift Survey \\cite{2dF}, have already ruled out one of the popular inflationary models driven by a self-interacting massless scalar field \\cite{Peiris}. Understanding the physics of inflation better will likely provide an important implication for late-time dark energy.  ``Radiation'' in the universe at around the matter-radiation equality mainly consists of photons and neutrinos; however, neutrinos actually stop being radiation when their mean energy per particle roughly equals the temperature of the universe. The physics of neutrinos has been revolutionized over the last decade by solar, atmospheric, reactor, and accelerator neutrino experiments having provided strong evidence for finite neutrino masses via mixing between different neutrino flavors, the so-called neutrino oscillations \\cite{exp1,exp2,exp3,exp4,exp5}. These experiments are, however, only sensitive to mass square differences between neutrino mass eigenstates, implying $\\Delta m_{21}^2\\simeq 7\\times 10^{-5}$~eV$^2$ and $\\Delta m_{32}^2\\simeq 3\\times 10^{-3}$~eV$^2$; thus, the most fundamental quantity of neutrinos, the absolute mass, has not been determined yet. Cosmological neutrinos that are the relic of the cosmic thermal history have distinct influences on the structure formation. Their large energy density, comparable to the energy density of photons before the matter-radiation equality, determines the expansion history of the universe. Even after the matter-radiation equality, neutrinos having become non-relativistic affect the structure formation by suppressing the growth of matter density fluctuations at small spatial scales owing to their large velocity dispersion \\cite{BES,BS,Pog1,Pog2,Ma95,Ma} (see Sec.~II and Appendix~A for more details). Therefore, the galaxy redshift surveys, combined with the CMB data, provide a powerful, albeit indirect, means to constraining the neutrino properties \\cite{HET,Hannestad1,Lahav,LahavSuto,Hannestad2}. This approach also complements the theoretical and direct experimental efforts for understanding the neutrino physics.  In fact, the cosmological constraints have placed the most stringent upper bound on the total neutrino mass, $m_{\\nu,{\\rm tot}}\\simlt 0.6\\,{\\rm eV}$ ($2\\sigma$) \\cite{WMAPp}, stronger than the direct experiment limit $\\simlt 2\\,{\\rm eV}$ \\cite{Bonn}. In addition, the result obtained from the Liquid Scintillator Neutrino Detector (LSND) experiment, which implies $\\bar{\\nu}_\\mu$ to $\\bar{\\nu}_e$ oscillations with $\\Delta m^2 \\simgt 0.2$\\,eV$^2$ \\cite{LSND} in an apparent contradiction with the other neutrino oscillation experiments mentioned above, potentially suggests the need for new physics: the cosmological observations will provide independent tests of this hypothesis.  In this paper we shall study the capability of future galaxy surveys at high redshifts, combined with the CMB data, for constraining (1) the neutrino properties, more specifically the total neutrino mass, $m_{\\nu, {\\rm tot}}$, and the number of non-relativistic neutrino species, $N_\\nu^{\\rm nr}$, and (2) the shape of the primordial power spectrum that is parameterized in terms of the spectral tilt,  $n_s$, and the running index, $\\alpha_s$, motivated by inflationary predictions (see Appendix~B). For the former, we shall pay particular attention to our ability to simultaneously constrain $m_{\\nu, {\\rm tot}}$ and $N_\\nu^{\\rm nr}$, as they will provide important clues to resolving the absolute mass scale as well as the neutrino mass hierarchy. The accuracy of determining  the neutrino parameters and the power spectrum shape parameters will be derived using the Fisher information matrix formalism, including marginalization over the other cosmological parameters as well as the galaxy bias.  Our analysis differs from the previous work on the neutrino parameters in that we fully take into account the two-dimensional nature of the galaxy power spectrum in the line-of-sight  and transverse directions, while the previous work used only spherically averaged, one-dimensional power spectra. The geometrical distortion due to cosmology and the redshift space distortion due to the peculiar velocity field will cause anisotropic features in the galaxy power spectrum. These features help to lift degeneracies between cosmological parameters, substantially reducing the uncertainties in the parameter determinations. This is especially true when variations in parameters of interest cause modifications in the power spectrum shape, which is indeed the case for the neutrino parameters, tilt and running index. The usefulness of the two-dimensional power spectrum, especially for high-redshift galaxy surveys, has been carefully investigated in the context of the prospected constraints on late-time dark energy properties \\cite{SE,Taka03,HuHaiman,Blake,GB05,taka05,Yamamoto}.  We shall show the parameter forecasts for future wide-field galaxy surveys that are already being planned or seriously under consideration: the Fiber Multiple Object Spectrograph (FMOS) on Subaru telescope \\cite{fmos}, its significantly expanded version, WFMOS \\cite{wfmos}, the Hobby--Ebery Telescope Dark Energy eXperiment (HETDEX) \\cite{hetdex}, and the Cosmic Inflation Probe (CIP) mission \\cite{cip}. To model these surveys, we consider three hypothetical galaxy surveys which probe the universe over different ranges of redshift, (1) $0.5\\le z\\le 2$, (2) $2\\le z\\le 4$ and (3) $3.5\\le z\\le 6.5$. We fix the sky coverage of each survey at $\\Omega_{\\rm s}=300$~deg$^2$ in order to make a fair comparison between different survey designs. As we shall show below, high-redshift surveys are extremely powerful for precision cosmology because they allow us to probe the linear power spectrum down to smaller length scales than surveys at low redshifts, protecting the cosmological information against systematics due to non-linear perturbations.  We shall also study how the parameter uncertainties are affected by changes in the number density of sampled galaxies and the survey volume. The results would  give us a good guidance to defining the optimal survey design to achieve the desired accuracies in parameter determinations.  The structure of this paper is as follows. In Sec.~\\ref{nu}, we review the physical pictures as to how the non-relativistic (massive) neutrinos lead to scale-dependent modifications in the growth of mass clustering relative to the pure CDM model. Sec.~\\ref{pps} defines the parameterization of the primordial power spectrum motivated by inflationary predictions. In Sec.~\\ref{formalism} we describe a methodology to model the galaxy power spectrum observable from a redshift survey that includes the two-dimensional nature in the line-of-sight and transverse directions. We then present the Fisher information matrix formalism that is used to estimate the projected uncertainties in the cosmological parameter determination from statistical errors on the galaxy power spectrum measurement for a given survey. After survey parameters are defined in Sec.~\\ref{survey}, we show the parameter forecasts in Sec.~\\ref{results}. Finally, we present conclusions and some discussions in Sec.~\\ref{conc}.  We review the basic properties of cosmological neutrinos in Appendix~A, the basic predictions from inflationary models for the shape of the primordial power spectrum in Appendix~B, and the relation between the primordial power spectrum and the observed power spectrum of matter density fluctuations in Appendix~C.  In the following, we assume an adiabatic, cold dark matter (CDM) dominated cosmological model with flat geometry, which is supported by the WMAP results \\cite{WMAP,WMAPp}, and employ the the notation used in \\cite{HE98,HE99}: the present-day density of CDM, baryons, and non-relativistic neutrinos, in units of the critical density, are denoted as $\\Omega_{\\rm c}$, $\\Omega_{\\rm b}$, and $\\Omega_{\\nu}$, respectively. The total matter density is then $\\Omega_{\\rm m}=\\Omega_{\\rm c}+\\Omega_{\\rm b}+\\Omega_{\\nu}$, and $f_\\nu$ is the ratio of the massive neutrino density contribution to $\\Omega_{\\rm m}$: $f_\\nu=\\Omega_\\nu/\\Omega_{\\rm m}$. ",
        "conclusions": "\\label{conc}  The non-relativistic neutrinos and the tilt and running index of the primordial power spectrum cause scale-dependent modifications in the linear power spectrum probed by galaxy redshift surveys. We have shown that one can determine these parameters  precisely by fully exploiting the two-dimensional information of galaxy clustering in angular and redshift directions from high-redshift galaxy surveys, when combined with the CMB data from the Planck experiment. The main results are summarized in Tables~\\ref{tab:fnu005} and ~\\ref{tab:fnu001}, and may be graphically viewed in Figures~\\ref{fig:fnu-Nnu} and \\ref{fig:alpha}. Our conclusions are two fold.  The first conclusion is for the neutrino parameters. We have found that the future galaxy surveys with $\\Omega_{\\rm s}=300$ deg$^2$ can provide very tight constraints on the total neutrino mass. The neutrino oscillation experiments have given the solid lower bound on the total neutrino mass, $m_{\\nu, {\\rm tot}}\\simgt 0.06~$eV, which is actually larger than the projected error on the total neutrino mass expected from the high-redshift galaxy surveys we have considered, up to by a factor of 2.5 for the space-based survey targeting galaxies at $3.5<z<6.5$. If two neutrino species have nearly equal masses (the inverted hierarchy), then the lower bound from the neutrino oscillation experiments, $m_{\\nu, {\\rm tot}}\\simgt 0.1~$eV, is up to 4 times larger than the projected errors of cosmological experiments. Overall, the high-redshift galaxy surveys combined with Planck allow a positive {\\em detection} of the total neutrino mass rather than the upper limit, improving the constraints on the total neutrino mass by a factor of 20--40 compared with the current cosmological constraints. The error on $m_{\\nu, {\\rm tot}}$ that we have found is smaller than that shown in the previous work (e.g., see \\cite{HET}) by a factor of 5 -- 10, despite the fact that the survey volume we assumed is larger than theirs only by a factor 1.5 -- 3. The main reason for the significant improvement is because our analysis exploits the full two-dimensional information in the galaxy power spectrum in redshift space. In particular, the redshift space distortion due to peculiar velocity field significantly helps improve the parameter determinations by breaking degeneracies between the cosmological parameters and the galaxy bias (also see Table \\ref{tab:noCMB}).  In addition, we have carefully investigated how one can use the future surveys to constrain the number of non-relativistic neutrino species, which should play an important role in resolving the neutrino mass hierarchy problem as well as the neutrino absolute mass scale, independently of the total neutrino mass. While we have found that the accuracy needed to discriminate between two models with the same $m_{\\nu, {\\rm tot}}$ but different $N_\\nu^{\\rm nr}$, $\\sigma(N^{\\rm nr}_\\nu)\\simlt 1$, is going to be difficult to achieve for the nominal survey designs we considered, one may achieve the desired precision by enlarging the survey volume. It should be stressed here that it is extremely important to exploit independent constraints on $m_{\\nu, {\\rm tot}}$ and $N_\\nu^{\\rm nr} $ from future cosmological observations, as any results unexpected from the point of view of the standard three-flavor neutrino model would imply anomalies in our understanding of the neutrino physics and hints for the new physics. Needless to say, controlling the systematics in such observations is also a crucial issue.  The second conclusion is for the shape of the primordial power spectrum. We have graphically summarized the expected performance of the future galaxy surveys for constraining the tilt, $n_s$, and the running index, $\\alpha_s$, of the primordial power spectrum in Figure \\ref{fig:alpha}. Compared with the constraints from the CMB data alone, the galaxy surveys we have considered can improve the determinations of $n_s$ and $\\alpha_s$ by a factor of 2 and 3, % yielding $\\sigma(n_s)\\sim 0.003$ and $\\sigma(\\alpha_s)\\simeq 0.002$, respectively. The high-redshift galaxy surveys allow us to probe galaxy clustering in the linear regime down to smaller length scales than at low redshifts. It is also important to note that the galaxy survey and CMB are sensitive to the primordial power spectrum shape at different $k$ ranges, and therefore these two probes are nicely complementary to each other in terms of constraining the tilt and the running index at different $k$. We have explicitly shown that the space-based galaxy survey (SG), such as CIP, has a similar level of precision for constraining $n_s$ and $\\alpha_s$ at the pivot scale $k_{\\rm 0 pivot}\\sim 0.5$ Mpc$^{-1}$, to the Planck experiment at $k_0=0.05~$Mpc$^{-1}$ (see Table \\ref{tab:k0pivot}). While it is sometimes argued that simple inflationary models should predict negligible amounts of running index which are out of reach of any experiments, our results show that the projected error on $\\alpha_s$ from the SG survey with five times more survey volume  is actually as small as $\\alpha_s$ predicted by the simple model based on a massive, self-interacting scalar field, $\\alpha_s=-(0.8-1.2)\\times 10^{-3}$. Our results for $n_s$  and $\\alpha_s$ are still on the conservative side,  as we have ignored the baryonic oscillations in the analysis. We have found that one can reduce the uncertainties in $n_s$ and $\\alpha_s$ by a factor of 2--3 by including the baryonic oscillations --- however, in that case one might also have to worry about the effect of curvature of the universe and dark energy properties. We report the results of the full analysis elsewhere.  Throughout this paper, we have assumed the standard {\\em three}-flavor neutrinos, some of which have become non-relativistic by the present epoch, as favored by neutrino oscillation experiments and the current theory of particle physics. On the other hand, cosmological observations such as CMB and the large-scale structure can also put independent constraints on the number of {\\it relativistic}, weakly interacting particles just like neutrinos, as a change in the relativistic degrees of freedom directly affects the expansion rate of the universe during the radiation era. For example, an increase in the relativistic degrees of freedom delays the matter-radiation equality, to which the CMB and large-scale structure observables are sensitive. The number of relativistic particles (minus photons) is conventionally expressed in terms of the effective number of neutrino species (i.e., the temperature of additional relativistic species is assumed to be $(4/11)^{1/3}$ of the temperature of photons, just like ordinary neutrinos). The current cosmological bound is $N_{\\nu}=4.2^{+1.2}_{-1.7}$ (95\\% C.L.) \\cite{Hannestad05}, while the standard model predicts $N_{\\nu}=3.04$ \\cite{QED} (also see \\cite{Dodelson} for the constraint on the abundance and mass of the sterile neutrinos based on the recent cosmological data sets). Hence, it is interesting to explore how one can simultaneously constrain the number of relativistic degrees of freedom before the photon decoupling epoch  as well as the number of non-relativistic particle species at low redshifts, $N^{\\rm nr}_\\nu$, using the future cosmological data sets. Properties of such weakly interacting particles are still difficult to measure experimentally.  It has been shown in the literature that the weak gravitational lensing \\cite{TJ,Kev,Song} or the number count of galaxy clusters \\cite{Fukugita,Haiman} can also be a powerful probe of the cosmological parameters. Different methods are sensitive to the structure formation at different ranges of redshifts and wavenumbers, and have different dependence on the cosmological parameters. More importantly, they are subject to very different systematics. Hence, by combining several  methods including the galaxy surveys considered in this paper, one can check for  systematics inherent in one particular method. The combination of different methods may also reduce statistical errors on the cosmological parameters. While it is worth exploring this issue carefully, we can easily imagine that, for the neutrino parameters, the weak lensing or the cluster number count does not improve $N_\\nu^{\\rm nr}$ very much. To constrain $N_\\nu^{\\rm nr}$ better one has to find a way to probe very large spatial scales, larger than the neutrino free-streaming scales; however, these methods probe only very small spatial scales where fluctuations have already become non-linear. (They may provide improvements in the determination of $m_{\\nu,{\\rm tot}}$.)  There is a promising way to check for systematics and improve the parameter determinations using the galaxy survey alone. While we have been assuming that perturbations are strictly linear and hence perfectly Gaussian, in reality small non-linearity always exists. The corrections to the power spectrum due to such small non-linearity may be calculated analytically using the higher-order cosmological perturbation theory (see \\cite{Bernardeau} for a review), which works extremely well at $z>2$ (Komatsu and Jeong, in preparation). Using the same higher-order cosmological perturbation theory, one may also compute the higher-order statistics, such as the bispectrum, which is a very powerful tool to check for systematics due to non-linearity in matter clustering, redshift space distortion, and galaxy bias \\cite{mat1,mat2}. Also, the bispectrum and power spectrum have different cosmological dependences \\cite{TJ,dolney}; thus, it is naturally expected that combining the two would improve the determinations of cosmological parameters. The results of our investigation along these lines will be reported elsewhere.  No matter how powerful the bispectrum could be in terms of checking for systematics, better models for non-linearity in redshift-space distortion and galaxy bias are definitely required for our projected errors to be actually realized. A good news is that we do not need a fully non-linear description of either component: we always restrict ourselves to the ``weakly non-linear regime'' where perturbation theory should still be valid. Having an accurate model for the redshift space distortion in the weakly non-linear regime is important for the precise determination of the total neutrino mass, as the redshift distortion plays a major role in lifting the degeneracy between the galaxy bias and the matter power spectrum amplitude (see Table~\\ref{tab:noCMB} what happens when the redshift space distortion is ignored). Recently, a sophisticated model of the distortion effect including the weak non-linear correction was developed in \\cite{Roman} based on the analytic method as well as the simulations. Likewise, it will be quite possible to develop a sufficiently accurate, well-calibrated model of the distortion effect at least on large length scales, based on adequate simulations. While we have employed a scale-independent linear bias throughout this paper, this model must break down even at weakly non-linear regime. Analytical \\cite{mat3} as well as numerical \\cite{white05,SE05,white05b} work has shown that deviations from a scale-independent bias do exist even on large spatial scales. This effect would also become particularly important for the precision measurements of the baryonic oscillations for constraining properties of dark energy. Therefore, careful and systematic investigations based on both analytical and numerical tools are needed to understand the realistic effect of scale-dependent bias on estimates of the cosmological parameters. As we have mentioned already, information from the higher-order statistics in the galaxy clustering would be a powerful diagnosis tool to check for systematics due to non-linear bias. As perturbation theory predicts that the galaxy power spectrum and bispectrum should depend on the galaxy bias differently, one can directly determine the galaxy bias and cosmological parameters simultaneously, by combining the two statistical quantities \\cite{mat1,mat2,dolney,Feldman,verde}. This method should also allow us to see a potential scale-dependent biasing effect from the epoch of reionization \\cite{babich/loeb:2005}.  Finally, let us comment on survey parameters. In order to make our discussion general, we have considered three hypothetical surveys which are different in their redshift coverage and the number density of targeted galaxies (see Sec.~\\ref{survey} for the survey definition). We have also explored how the parameter errors would change when the survey parameters are varied from the fiducial values (see Table \\ref{tab:vn} and Figure \\ref{fig:fnu-Nnu_VN}). Increasing a survey volume has a greater impact on the parameter errors compared with increasing the numbers of targeted galaxies for a given survey volume. We are hoping that our results provide useful information to help to define an optimal survey design to attain the desired accuracy on the parameter determinations, given the  limited observational resources and budget.  {\\it Acknowledgments:} We thank K.~Ichiki, T.~Murayama, T.~Yoshida, and M.~White for helpful discussions, and also thank an anonymous referee for useful comments. This work was supported in part by a Grand-in-Aid for Scientific Research (17740129) of the Ministry of Education, Culture, Sports, Science and Technology in Japan as well as by the COE program at Tohoku University.  M.T. acknowledges the warm hospitality  of University of Texas at Austin where this work was partly done. E.K. acknowledges support from the Alfred P. Sloan Foundation and NASA grant NNGO4GQ39G.  We acknowledge the use of the publicly-available CMBFAST code \\cite{cmbfast}.  \\appendix"
    },
    "0512/astro-ph0512004_arXiv.txt": {
        "abstract": "We investigate the extent to which galaxies' star-formation histories and morphologies are determined by their clustocentric distance and their nearest cluster's richness.  We define clustocentric distance as the transverse projected distance between a galaxy and its nearest cluster center within $\\pm 1000~\\mathrm{km\\,s^{-1}}$.  We consider a cluster to be a bound group of at least $\\N$ galaxies brighter than $M_r=-19$~mag; we employ different richness criteria of $\\N=5$, 10, and 20.  We look at three tracers of star-formation history ($^{0.1}i$ band absolute magnitude $M_{^{0.1}i}$, $^{0.1}[g-r]$ color, and \\Halpha\\ emission line equivalent width) and two indicators of galaxy morphology (surface brightness and radial concentration) for 52,569 galaxies in the redshift range of $0.015<z<0.068$.  We find that our morphology indicators (surface brightness and concentration) relate to the clustocentric distance only indirectly through their relationships with stellar population and star formation rate.  Galaxies that are near the cluster center tend to be more luminous, redder and have lower \\Halpha\\ EW (\\ie, lower star-formation rates) than those that lie near or outside the virial radius of the cluster.  The detailed relationships between these galaxy properties and clustocentric distance depend on cluster richness.  For richer clusters, we find that \\textsl{(i)}~the transition in color and \\Halpha\\ EW from cluster center to field values is more abrupt and occurs closer to the cluster virial radius, and \\textsl{(ii)}~the color and \\Halpha\\ EW distributions are overall narrower than in less rich clusters.  We also find that the radial gradient seen in the luminosity distribution is strongest around the smaller clusters and decreases as the cluster richness (and mass) increases.  We find there is a ``characteristic distance'' at around one virial radius (the infall region) where the change with radius of galaxy property distributions is most dramatic, but we find no evidence for infall-triggered star bursts.  These results suggest that galaxies ``know'' the distance to, and the size of, their nearest cluster and they express this information in their star-formation histories. ",
        "introduction": "The statistical properties of galaxies are closely related to the densities of their surrounding environments.  Since regions of the Universe with different densities evolve at different rates, we expect these environment dependencies to contain crucial information about galaxy formation and evolution.  Much of previous environment-related work focused on the relationship between morphology and environment (\\citealt{hubble36a, oemler74a, dressler80a, postman84a}; or the more recent work of \\citealt{hermit96a, guzzo97a, giuricin01a, trujillo02a}).  All these works find that bulge-dominated galaxies are more strongly clustered than disk-dominated galaxies.  Spectroscopic and photometric properties of galaxies are strongly correlated with morphology so it is not surprising to find that the spectroscopic and photometric properties of galaxies are also functions of clustering \\citep{kennicutt83a, hashimoto98a, balogh01a, martinez02a, lewis02a, norberg02a, hogg03b, hogg04a, kauffmann04a, zehavi05b}.  While trying to understand \\emph{why} these properties relate with environment, it is important to ask which properties are correlated with environment independently of the others.  Previous work has found that, out of color, luminosity, surface brightness, and S\\'ersic index (concentration), the color and luminosity of a galaxy appear to be the only properties that are directly related to the local overdensity \\citep{blanton05b}.  Surface brightness and concentration appear to be related to the environment only through their relationships with color and luminosity.  The clustocentric distance---the distance to the nearest rich cluster center---is a fundamental and precisely measurable environment indicator; indeed it can be measured at much higher signal-to-noise than estimates of overdensity on fixed scales.  Previous observational results that use the clustocentric distance as an environment indicator have found that the star-formation rates of galaxies are correlated with the clustocentric distance in the following ways: \\textsl{(i)} the star-formation rates of galaxies near the centers of clusters are low and they rise with increasing clustocentric distance. \\textsl{(ii)} the correlation is steeper for galaxies with higher star-formation rates, and \\textsl{(iii)} there is a `characteristic break' at a distance of 3-4 cluster virial radii at which the change in the correlation is greatest \\citep{lewis02a, gomez03a}.  It has been suggested that gas-stripping plays a significant role in this rapid truncation of star-formation, as well as the morphological transformation that is observed in the cluster population \\citep{vogt04a}, although most statistical studies find that this process cannot be extremely rapid \\citep{kodama01a, balogh00a}. Models generally find that the star-formation rates of cluster galaxies depend primarily on the time elapsed since their accretion onto the massive system, but that the cessation of star formation may take place gradually over a few billion years \\citep{balogh00a}.  In this short paper, we build upon previous studies in several ways. We use the \\citet{berlind06a} cluster catalog that was created with a different algorithm than previous studies.  In particular, the algorithm did not use galaxy colors in the group and cluster identification, so our galaxy colors are unbiased even in the cluster cores.  Moreover, the cluster catalog was designed to recover clusters of galaxies that occupy the same underlying dark matter halo.  Consequently, the abundance of these clusters is unbiased relative to the abundance of halos and can thus be used to obtain robust estimates of the cluster masses.  The method we use to estimate cluster masses differs significantly from the methods used in previous studies.  This difference leads to a different conclusion regarding the activity observed in the theoretically active infall regions around clusters.  Finally, we examine the relationship between the morphology--environment relation and the color--environment relation.  We use the Sloan Digital Sky Survey (SDSS) data to measure the clustocentric distances of $\\sim52,000$ galaxies relative to $\\sim950$ clusters within the redshift range $0.015<z<0.068$.  We use high quality photometric and spectroscopic data to examine the dependence of $^{0.1}[g-r]$ color, absolute magnitude $M_{^{0.1}i}$, \\Halpha\\ equivalent width (EW), concentration (S\\'ersic index), and surface brightness $\\mu_{^{0.1}i}$ on both clustocentric distance and cluster richness.  The large solid angular coverage of the SDSS permits the study of galaxies from cluster centers through the ``infall regions'' \\citep{poggianti99a, balogh00a, kodama01a} and out to huge clustocentric distances.  In what follows, a cosmological world model with $(\\Omega_\\mathrm{M},\\Omega_\\mathrm{\\Lambda})=(0.3,0.7)$ is adopted, and the Hubble constant is parameterized $H_0\\equiv100\\,h~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$, for the purposes of calculating distances and volumes with $h=1$ except where otherwise noted \\citep[\\eg,][]{hogg99cosm}. ",
        "conclusions": "Using a complete sample of 52,569 galaxies in the redshift range of $0.015<z<0.068$ we examine how $^{0.1}i$ band absolute magnitude $M_{^{0.1}i}$, $^{0.1}[g-r]$ color, \\Halpha\\ EW, surface brightness, and concentration are related to both the transverse projected clustocentric distance, \\Dcl, and the richness, \\N, of its nearest cluster.  Similar to previous results \\citep{blanton05b,christlein05a}, we find that our morphology tracers (concentration and surface brightness) appear to be related to environment (\\Dcl\\, in this case) only indirectly through their relationships with the star-formation history tracers. This suggests that the well known morphology--environment relation is a residual of the star-formation-history--environment relation.  We note that, although our simple morphology indicators do not show direct dependence on clustocentric distance, this result does not disagree with that the original morphology--density studies \\citep{hubble36a, oemler74a, dressler80a, postman84a} because these studies did not attempt to separate morphological and star-formation dependences.  Our results also do not disagree with previous work \\citep{vogt04a} which suggests that gas stripping plays a significant role in the morphological transformation and rapid truncation of star formation by showing asymmetries in HI and \\Halpha\\ flux on the leading edge of infalling spiral galaxies.  There is no disagreement because our very blunt morphology indicators are insensitive to these asymmetries.  We have shown that galaxies know about the distance to, \\emph{and richness of}, their nearest cluster and they express that knowledge most clearly through their star-formation histories.  Some morphological evidence suggests that they also know the direction as well \\citep{vogt04a}.  Where our results overlap those of previous investigators \\citep{lewis02a, gomez03a, blanton05b}, we mostly find good agreement. In particular, we find that galaxy properties depend on clustocentric distance much the same way as they do on other environmental indicators \\citep{blanton05b}.  Our result improves on this previous one because the clustocentric distance is measured at much higher signal-to-noise and probes environments on scales much smaller than local overdensity measurements.  Where our results disagree with previous work, it can mainly be attributed to the different methods of computing cluster virial radii.  While it may seem that this is a minor discrepancy, these disagreements become important at the interpretation level.  In particular, the claim that clusters affect galaxies well beyond the cluster virial radii \\citep{lewis02a, gomez03a} depends strongly on the calculation of virial radius.  Our virial radii are estimated from cluster abundances (and an assumed ``concordance'' cosmological model), whereas in previous work they are estimated from cluster velocity dispersions.  In order to check the difference between these two methods, we computed a second set of virial radii for our clusters using the \\cite{berlind06a} velocity dispersions and the \\cite{girardi98a} approximation that relates virial radius to velocity dispersion.  We find that the virial radii based on velocity dispersions are systematically lower than those based on abundances by a factor of two to three.  This is sufficient to explain the difference in interpretation between \\cite{gomez03a} and this work.  We choose to measure our virial radii using cluster abundances because the \\cite{berlind06a} cluster catalog was tuned to produce unbiased abundances and we, therefore, believe them to be less subject to systematic errors than velocity dispersion estimates.  However, this point deserves to be studied in more detail.  After analyzing these results, we are in a position to address the theoretically active ``infall region'' \\citep{poggianti99a, balogh00a, kodama01a} around virialized systems.  All three trends (color, \\Halpha, and absolute magnitude) as a function of \\Dcl\\ have a ``characteristic break'' at about the same virial-radius-normalized distance from the cluster center but we find no clear evidence for an ``increase'' in star formation activity at the infall region.  Here we have investigated how galaxy colors, luminosities, and star-formation rates (\\Halpha\\ EW) relate to clustocentric distance and how these relations depend on cluster richness.  We find that for larger clusters the transitions in the typical values of color and star-formation rate from cluster centers to the field are more abrupt and occur closer to the viral radius than those for smaller clusters. We also find the increase in typical luminosity when looking from the field to cluster centers is greater around smaller clusters.  A question that naturally arises is: What physical physical processes are involved in creating these richness-dependent variations?  It could be that \\textsl{(i)}~the possible transformation mechanisms are stronger in larger clusters (\\eg, more frequent tidal interactions or a hotter intercluster medium), \\textsl{(ii)}~galaxies in larger clusters have been in this environment longer and therefore have had more time to evolve to their long-term properties, or (more likely) a combination of the two."
    },
    "0512/astro-ph0512232_arXiv.txt": {
        "abstract": "It has been suggested that Gamma Ray Bursts (GRB) may enable the expansion rate of our Universe to be measured out to very high redshifts ($z \\gsim 5$) just as type Ia supernovae have done at $z \\sim$1--1.5. We explore this possibility here, and find that GRB have the potential to detect dark energy at high statistical significance, but they are unlikely to be competitive with future supernovae missions, such as SNAP, in measuring the properties of the dark energy. The exception to this conclusion is if there is appreciable dark energy at early times, in which case the information from GRB's will provide an excellent complement to the $z\\sim 1$ information from supernovae. ",
        "introduction": "By observing distant type Ia supernovae (SNe Ia), the expansion history of our Universe has been studied out to redshift of order unity. These observations have showed that the expansion rate of our Universe is accelerating \\cite{accelerating}, which is generally attributed to the presence of dark energy. These observations allow us to measure the expansion rate only over recent times, however. If our Universe's expansion rate could be measured over a range of higher redshifts, that could provide a useful discriminator of various models of dark energy -- especially those models with a sizable quantity of dark energy at early times.  With this in mind, we here consider using Gamma Ray Bursts (GRB) as cosmological probes. The primary advantage of GRB over SNe Ia is that they are considerably brighter and thus can be observed at much higher redshifts. The {\\it average} GRB observed by the currently operating Swift satellite has a redshift of $\\sim2.8$ \\cite{swiftmean}, whereas the highest redshift SNe Ia among the 157 in the \"gold\" sample of Ref.~\\cite{gold} is less than 1.8. Swift is expected to detect several GRB each year at $z > 5$.  GRB have substantial disadvantages in comparison to Ia SNe, as well. SNe Ia are powerful probes of our Universe's expansion because their intrinsic luminosity can be inferred independently of any redshift measurement -- ie. they are standardizable candles. The degree to which GRB can be used as standard candles is not yet fully known. Although the quantity of (geometrically corrected) energy released in the jets of long duration GRB is clustered around $1.3\\times 10^{51}$ ergs \\cite{ejet}, the dispersion of this quantity is too large to be of much use in studying cosmology. This dispersion can be considerably reduced, however, by making use of the various known correlations between the luminosity of a burst and its other observable parameters. Such correlations include those associated with the variability \\cite{variability} and spectral lag \\cite{lag} of a GRB. In addition to these, two other well-known GRB luminosity indicators have appeared in the literature \\cite{other}. We will not discuss these here, as they are either too inaccurate or poorly established to be adopted reliably at this time. It appears at least plausible, however, that observations by Swift and other experiments will improve our ability to accurately determine the intrinsic luminosity of GRB. The degree to which this can be accomplished will determine the usefulness of GRB to cosmology.  Here we study how much can be learned about cosmological expansion from GRB's. In particular, we ask whether future Swift observations will pin down properties of the dark energy, and if so, how these projected constraints compare with the expected results from Ia SNe. ",
        "conclusions": "In this article, we have studied the possibility of constraining the properties of dark energy using future observations of Gamma Ray Bursts (GRB). GRB are brighter than Supernovae (SNe), and thus can be observed at much higher redshifts. The Swift satellite, for example, can measure several GRB at redshift above 5 each year. This is in contrast to type Ia SNe observations, which are limited to redshifts smaller than 1.7, even with an experiment such as SNAP.  The weakness of using GRB as cosmological probes lies in the degree of scatter in their intrinsic luminosities. The degree to which GRB are standardizable candles is not yet well known. Currently known relationships between GRB luminosity and other independent observables (such as variability and spectral lag) have led to magnitude dispersions on the scale of $\\sigma \\sim 0.6$. With this level of dispersion, GRB should be able to detect dark energy at a significance above 5$\\sigma$, thus providing an independent verification of SNe observations.  It is quite plausible, however, that future observations will enable GRB luminosites to be determined with substantially higher accuracy. Even if the magnitude dispersion of GRB can be substantially reduced, however, these objects will probably never compete with SNe as probes of dark energy. Part of the reason for this is that lensing by large scale structure becomes more important for sources at high redshifts. This results in a floor on the total dispersion of high redshift GRB and ultimately on the statistical error on the dark energy equation of state.  The one caveat to this conclusion is if there is appreciable early ($z>1.5$) dark energy. In that case, GRB will provide a useful and complementary probe of the expansion history of our Universe.    {\\it Acknowledgments:} We would like to thank Chris Quigg, Mike Stamatikos and Gajus Miknaitis for helpful discussions. This work has been supported by the US Department of Energy and by NASA grant NAG5-10842."
    }
}